{
    "0410/astro-ph0410421_arXiv.txt": {
        "abstract": "{We show that the shape of the observed distribution of Milky Way (MW) satellites is inconsistent with being drawn from a cosmological sub-structure population with a confidence of 99.5 per cent. Most of the MW satellites therefore cannot be related to dark-matter dominated satellites. ",
        "introduction": "\\label{sec:intro}  Calculations of structure formation within the framework of cold dark matter (CDM) cosmology show that Milky Way (MW) type systems have the same scaled theoretical distribution of sub-haloes as rich galaxy clusters, and within 500~kpc they should contain about 500 sub-haloes with masses $M\\simgreat 10^8\\,M_\\odot$ (Moore et al.\\ \\cite{moore99}; Klypin et al.\\ \\cite{klypin99}; Governato et al. 2004). However, only 13~dwarves have been found within a distance of 500~kpc around the MW. The observed dwarves may only sample a sub-set of the actually present CDM sub-structures (Stoehr et al. 2002; Hayashi et al. 2003; Bullock et al. 2000; Susa \\& Umemura 2004; Kravtsov, Gnedin \\& Klypin 2004). Such biasing could be the result of complex early baryonic physics that cannot, at present, be treated theoretically in sufficient detail, but Kazantzidis et al. (2004) point out that this cannot be the entire solution.  An additional path to testing predictions of CDM cosmology is to compare the shape of the observed satellite distribution to the theoretical shapes (Zaritsky \\& Gonzalez 1999; Hartwick 2000; Sales \\& Lambas 2004).  The sub-structures fall inwards from filaments that are thicker than the virialised regions of the hosts. However, within its virialised region, the number distribution of sub-structure in a theoretical host halo follows that of its dark-matter (DM) distribution.  CDM models predict the host DM haloes to be oblate with flattening increasing with increasing mass and radius (Combes 2002; Merrifield 2002). The ratio of minor to major axis of the DM density distribution has the value $q_{\\rm d}=0.7\\pm0.17$ for MW-sized haloes within the virial radius.  The intermediate to major axis ratio is $q_{\\rm d}'\\simgreat 0.7$ (Bullock 2002).  When dissipational baryonic physics is taken into account the haloes become more axis-symmetric (larger $q_{\\rm d}'$) and more flattened, $q_{\\rm d}=0.5\\pm0.15$ within the virial radius.  The minor axis is co-linear to the angular momentum of the baryonic disk (Dubinski 1994). Prolate haloes do not emerge.  The empirical evidence is that the MW dark halo is somewhat flattened (oblate) with $q_{\\rm d}\\simgreat 0.8$ within $R\\simless 60$~kpc (Olling \\& Merrifield 2000, 2001; Ibata et al.  2001; Majewski et al.  2003; Mart{\\'{\\i}}nez-Delgado et al. 2004).  Beyond this distance the shape is likely to be more oblate (Bullock 2002), but invoking continuity the axis ratio $q_{\\rm d}$ cannot change drastically.  The theoretical sub-structure distribution of MW-type hosts must therefore be essentially isotropic (Ghigna et al. 1998; Zentner \\& Bullock \\cite{zentner03}; Diemand, Moore \\& Stadel 2004; Kravtsov et al. 2004; Aubert et al.\\ \\cite{aubert04}).  If the MW dwarves do indeed constitute the shining fraction of DM sub-structures, then their number-density distribution should be consistent with an isotropic (i.e.  spherical) or oblate power-law radial parent distribution.  This is assumed to be the case by most researchers, given the relatively small number of satellites. With this contribution we show that, despite its smallness, the MW satellite sample is inconsistent with a cosmological sub-structure population. We do this by concentrating on the most elementary facts, namely purely on the positions of the satellites. ",
        "conclusions": "\\label{sec:concs}  Cosmological models can be tested among other ways by comparing the theoretical sub-structure distribution with observed satellite distributions. The theoretical distribution contains about 500 sub-haloes within approximately 500~kpc of a MW-type galaxy and follows an approximately power-law radial distribution with $p\\simless 2$ and is essentially isotropic. The well-known MW distribution contains only a dozen dwarves, is indeed consistent with the theoretical radial distribution but is highly anisotropic. The anisotropy is such that the MW dwarves form a disk-like structure with a root-mean-square height of $10-30$~kpc which lies nearly perpendicularly to the plane of the MW. The pole of this great disk lies close to the orbital poles of the LMC, the SMC, Draco and Ursa Minor. The distance of closest approach of the plane to the Galactic centre, $D_{\\rm P}\\la 2$~kpc, is much smaller than the radial extent of the Galactic disk ($\\approx 20$~kpc) or even the root-mean square height, $\\Delta$, of the disk of satellites ($D_{\\rm P}\\ll \\Delta$). This is a strong indication that the sample of dwarves within about 250~kpc is relaxed in the Galactic potential.  Their orbits must be confined within the great disk because the likelihood of obtaining such a disk-like dwarf distribution given a true underlying isotropic distribution (that ought to match the sphericity of the MW DM halo) is less than 0.5~per cent. This result persists even after removing the kinematically related SMC and UMi from the analysis.  A distribution of polar orbits with arbitrary eccentricities and orientation of orbital planes is also excluded with the same confidence because it leads to an isotropic distribution of dwarves.  An oblate MW dark matter halo would yield an even larger discrepancy with the disk of satellites.  An alternative approach is taken by Hartwick (2000) who argues that the 10~satellites within 400~kpc (the LMC and SMC are combined to one satellite) map the MW DM halo shape and form a highly inclined and highly prolate system with minor/major axis ratio $q_{\\rm d}\\approx 0.03-0.05$.  However, the extreme triaxiality derived in this way is completely inconsistent with the observational and theoretical shapes of CDM host-haloes and sub-structure distributions (\\S~\\ref{sec:intro}).  The approach taken here differs by noting the very significant mismatch between (i)~the disk-like satellite distribution, (ii)~the {\\it independent}\\, empirical constraints on the shape of the MW dark matter halo, and (iii)~the theoretical shapes of CDM host haloes (\\S~\\ref{sec:intro}). In the view presented here, the mismatch between the number {\\it and}\\, spatial distribution of MW dwarves compared to the theoretical distribution challenges the claim that the MW dwarves are cosmological sub-structures that ought to populate the MW halo.  A more natural and more conservative (by not resorting to exotic physics) explanation for the MW dwarf distribution in a great disk with a ratio of height to radius of~0.1--0.2 would appear to be in terms of a causal connection between most of them.  This could be the case if most of the dwarves stem from one initial gas-rich parent satellite on an eccentric near-polar orbit that interacted with the young MW, perhaps a number of times, forming tidal arms semi-periodically as its orbit shrank. The early gas-rich tidal arms may have condensed in regions to tidal dwarf galaxies, as is observed in present-day interacting gas-rich galaxies (e.g. Knierman et al. 2003; Weilbacher, Duc \\& Fritze-v.~Alvensleben 2003).  The LMC may be the most massive remnant of this larger satellite, while the lesser dwarves may be its old children (Lynden-Bell \\cite{lyndenbell76}). The Magellanic Stream may be just such a newly formed but meagre tidal feature (Kunkel \\cite{kunkel79}), and the alignment of the disk of satellites with the surrounding matter distribution (Hartwick 2000) may simply result from the gas-rich parent satellite coming-in from that direction.  The different chemical enrichment and star-formation histories of the various dwarves (e.g.\\ Ikuta \\& Arimoto \\cite{ikuta02}; Grebel et al.\\ \\cite{grebel03}) may in this case be a result of their different initial masses that will have been significantly larger than their present-day baryonic masses (Kroupa \\cite{kroupa97}) and the complex interplay between stellar evolution, tides, gaseous stripping and gas accretion during the orbits within the MW halo, none of which are presently understood in much detail. The simulations of Kroupa have shown that ancient tidal-dwarf galaxies may appear similar to some of the observed dSph satellites.  The sub-structure under-abundance problem extends to fossil galaxy groups where early photo-evaporation could not have removed baryons from the sub-structures (D'Onghia \\& Lake 2004), and a sub-structure-overabundance is evident for rich clusters (Diemand et al. 2004). CDM cosmology thus faces a sub-structure challenge on all mass scales."
    },
    "0410/astro-ph0410617_arXiv.txt": {
        "abstract": "We present here a new calculation of the gamma-ray spectrum from $pp \\rightarrow \\pi^0$ in the Galactic ridge environment. The calculation includes the diffractive $p$-$p$ interaction and incorporates the Feynman scaling violation for the first time. Galactic diffuse gamma-rays come, predominantly, from $\\pi^0 \\rightarrow \\gamma \\gamma$ in the sub-GeV to multi-GeV range. Hunter et al. found, however, an excess in the GeV range (``GeV Excess'') in the EGRET Galactic diffuse spectrum above the prediction based on experimental $pp \\rightarrow \\pi^0$ cross-sections and the Feynman scaling hypothesis. We show, in this work, that the diffractive process makes the gamma-ray spectrum harder than the incident proton spectrum by $\\sim 0.05$ in power-law index, and, that the scaling violation produces 30$-$80\\% more $\\pi^0$ than the scaling model for incident proton energies above 100~GeV.  Combination of the two can explain about a half of the ``GeV Excess'' with the local cosmic proton (power-law index $\\sim 2.7$).  The excess can be fully explained if the proton spectral index in the Galactic ridge is a little harder ($\\sim 0.2$ in power-law index) than the local spectrum.  Given also in the paper is that the diffractive process enhances $e^+$ over $e^-$ and the scaling violation gives $50-100$\\% higher $\\bar{p}$ yield than without the violation, both in the multi-GeV range. ",
        "introduction": "Gamma-rays from neutral pions produced by cosmic-ray proton interactions with ISM have been predicted to dominate the diffuse Galactic emission in the sub-GeV to GeV band since 1960's. Early pioneers including \\citet{Ginzburg67} and \\citet{Hayakawa69} have estimated the gamma-ray flux from $\\pi^0$ together with other important mechanisms around that time.  First quantitative observation-based studies of the diffuse gamma-ray spectrum covering the sub-GeV and GeV band were made, eg. by \\citet{Strong78}, \\citet{SB81}, \\citet{Dermer86}, and \\citet{Stecker89}. They compared the data from COS-B \\citep{COS-B} with their models based on experimental $pp \\rightarrow \\pi^0$ data from accelerators, their extension to higher energies on the Feynman scaling hypothesis \\citep{Feynman69}, estimations on cosmic ray proton and electron fluxes, and the ISM distribution obtained by radio surveys. Within the uncertainties in the data and modeling, the studies cited above concluded that gamma-rays from $\\pi^0$ are a dominant component in the Galactic ridge spectrum above 100 MeV. The bremsstrahlung emission by $e^+/e^-$ off ISM atoms and the inverse-Compton scattering of infra-red and optical photons by $e^+/e^-$ are also expected to contribute significantly in the sub-GeV to GeV energy range \\citep{Hayakawa69,MurthyWolfendale93,Schoenfelder01}  The limited statistics and energy coverage of the COS-B gamma-ray data permitted only a crude consistency check of the $pp \\rightarrow \\pi^0 \\rightarrow \\gamma$ hypothesis within a factor $\\sim 2$ and left large ambiguity on the mix of emission mechanisms \\citep{SB81}.  The spatial distribution of the energy-integrated gamma-ray intensity, on the other hand, gave a higher statistical accuracy on which \\citet{Hasselwander82}, \\citet{Strong82}, \\citet{Bloemen84}, \\citet{Bloemen85}, and others set the path to the Galactic gamma-ray astronomy.  We refer to \\citet{MurthyWolfendale93}, \\citet{Schoenfelder01}, and \\citet{Schlickeiser02} for general references on the topics of this work.  When the much improved data obtained with EGRET \\citep{EGRET} were studied by \\citet{Bertsch93} and \\citet{Hunter97}, an excess of about $\\times (1.5-2)$ became apparent in the data in the GeV band relative to the Galactic gamma-ray emission models cited above. This excess is visible along the entire Galactic plane ($-10<b<10$ deg. and $-90<\\ell<+90$ deg.) but most pronounced in the central region of the plane ($-40<\\ell<+40$ deg.).  Here $b$ and $\\ell$ are the Galactic latitude and longitude, respectively. In literature, the excess is referred to as the ``GeV Excess''.  \\citet{Mori97} studied the diffuse gamma-ray emission by using Monte Carlo $p$-$p$ interaction simulators devloped for accelerator experiments and confirmed that the EGRET spectrum can not be reproduced with the conventional cosmic proton spectrum.  Strong and Moskalenko developed a computer program, GALPROP, where the cosmic ray propagation and interaction are numerically calculated in a model Galaxy to obtain the spatial distribution and spectrum of secondary particles including gamma-rays and radio isotopes \\citep{Galprop1,Galprop2}. The diffuse Galactic gamma-ray spectra have been calculated separately for the $\\pi^0$ decay, bremsstrahlung, and inverse-Compton with various parameter settings of GALPROP by \\citet{SMR00}.  They noted that: a) the local interstellar spectra of electrons and protons (power-law index $\\ge 2.5$ above $10-20$ GeV) do not reproduce the gamma-ray spectrum along the Galactic plane observed by EGRET; b) a very hard electron spectrum (power-law index $\\sim 1.8$) and a modified nucleon spectrum will be needed to minimize the difference between the prediction and the data in the GeV band; and c) the GeV excess in the central Galactic ridge can not be reproduced within the constraint on the proton spectral index imposed by recent cosmic proton measurements and the limit on the electron spectral index ($\\sim 1.9$) from radio and local cosmic ray observations.  They noted that the excess persists at higher Galactic latitude ($|\\ell|>5$ deg).  \\citet{Buesching01} have also noted that the EGRET spectrum is incompatible with the locally measured cosmic proton spectrum.  The GeV Excess has led to new optimizations of the Galactic gamma-ray emission models and speculations on possible new gamma-ray sources in the Galaxy. One choice is to assume a harder proton spectrum in the Galactic ridge region as has been noted by \\citet{Mori97}. Another choice is to assume a much higher electron flux with a broken power-law spectrum in the Galactic ridge \\citep{SMR04}. \\citet{Buesching01} has proposed to introduce a mix of spectral indices for protons which leads to a convex gamma-ray spectrum similar to that observed by EGRET. Possible contributions of unidentified pulsars \\citep{Pohl97} and dark-matter particle annihilation \\citep{deBoer03, Cesarini04} have also been studied.  We came to note that all calculations of the Galactic gamma-ray emission cited above have not included an important component of inelastic $p$-$p$ interaction, the diffractive interaction, nor incorporated the Feynman scaling violation in the non-diffractive inelastic interaction. Another important finding was that these calculations assume an obsolete $p$-$p$ non-diffractive inelastic cross-section model taking a constant value of $\\sim 24$~mb for $T_p \\gg 10$~GeV. Updating these shortfalls and inaccuracy can change the gamma-ray spectrum from the proton ISM interaction at high energies in the following ways: the diffractive process will add gamma-rays in the highest end of the spectrum; the scaling violation and the up-to-date non-diffractive inelastic cross-section will increase the gamma-ray yield in the GeV to multi-GeV range.   We built a model (model A) representing the latest knowledge on the $p$-$p$ inelastic interaction, and calculated the gamma-ray spectrum due to $\\pi^0$ produced in the cosmic-ray proton interaction with ISM. The $p$-$p$ interaction is simulated separately for the non-diffractive inelastic and for the diffractive processes. The non-diffractive process is calculated by two computer programs: Pythia 6.2 by \\citet{Pythia62} for the proton kinetic energy ($T_p$) range $512$~TeV $\\ge T_{p} \\ge 52.6$~GeV and the model by \\citet{SB81} with the parametrization of \\citet{Blattnig00} for $52.6$~GeV $\\ge T_{p} > 0.488$~GeV. The diffractive process is simulated by a program written for this study \\footnote{Simulation Program DiffDissocSimulNew.py for model A is available upon request from author.}: it is based on the formulae given in \\citet{Goulianos83}, \\citet{Goulianos95}, and \\citet{Goulianos99}. The non-diffractive and diffractive gamma-ray spectra are added according to the cross-section model for model A shown in Columns 2 and 3 of Table 1 and Fig.1a\\footnote{The first 2 data sets in Figure 1 are from \\citet{Hagiwara02} and the last from the Spires database (see http://www.slac.stanford.edu/spires).}.  \\begin{small}  \\begin{table} \\caption{$p$-$p$ Model Cross-Sections and Galactic Proton Spectral Models} \\begin{tabular}{rrrrrrr} \\hline\\hline $T_p$ & \\multicolumn{2}{c}{$\\sigma$(model A)} & $\\sigma$(model B) & \\multicolumn{3}{c}{Factors for Proton Spectra} \\\\ \\hline (GeV) & $\\sigma$(NonDiff) & $\\sigma$(Diff) & $\\sigma$(NonDiff) & Ind=2.0 & LIS & Trial4GR \\\\ \\hline 4.88E-01 & 5 & 0 & 5 & 2.05E+03 & 5.54E+04 & 3.05E+04\\\\ 6.90E-01 & 20 & 0 & 20 & 1.45E+03 & 4.58E+04 & 2.01E+04\\\\ 9.80E-01 & 23.4 & 0 & 23.4 & 1.02E+03 & 3.72E+04 & 1.32E+04\\\\ 1.38E+00 & 25 & 0 & 24 & 7.25E+02 & 2.91E+04 & 8.75E+03\\\\ 1.95E+00 & 27.57 & 0 & 24.6 & 5.13E+02 & 2.13E+04 & 5.78E+03\\\\ 2.76E+00 & 28.57 & 0 & 24.4 & 3.62E+02 & 1.49E+04 & 3.81E+03\\\\ 3.91E+00 & 29.27 & 0 & 24.1 & 2.56E+02 & 9.54E+03 & 2.51E+03\\\\ 5.52E+00 & 29.76 & 0 & 23.8 & 1.81E+02 & 5.85E+03 & 1.66E+03\\\\ 7.81E+00 & 23.96 & 6.13 & 23.4 & 1.28E+02 & 3.51E+03 & 1.09E+03\\\\ 1.11E+01 & 23.65 & 6.68 & 23 & 9.05E+01 & 2.04E+03 & 7.21E+02\\\\ 1.56E+01 & 23.29 & 7.22 & 22.6 & 6.40E+01 & 1.18E+03 & 4.75E+02\\\\ 2.21E+01 & 22.91 & 7.76 & 22 & 4.53E+01 & 6.52E+02 & 3.04E+02\\\\ 3.13E+01 & 22.53 & 8.29 & 21.6 & 3.20E+01 & 3.62E+02 & 1.81E+02\\\\ 4.42E+01 & 22.27 & 8.72 & 21.6 & 2.26E+01 & 2.01E+02 & 1.08E+02\\\\ 6.25E+01 & 22.23 & 8.96 & 21.6 & 1.60E+01 & 1.11E+02 & 6.40E+01\\\\ 8.84E+01 & 22.2 & 9.19 & 21.6 & 1.13E+01 & 6.18E+01 & 3.81E+01\\\\ 1.25E+02 & 22.14 & 9.43 & 21.6 & 8.00E+00 & 3.43E+01 & 2.26E+01\\\\ 1.77E+02 & 22.14 & 9.67 & 21.6 & 5.66E+00 & 1.90E+01 & 1.35E+01\\\\ 2.50E+02 & 22.22 & 9.91 & 21.6 & 4.00E+00 & 1.06E+01 & 8.00E+00\\\\ 3.54E+02 & 22.36 & 10.15 & 21.6 & 2.83E+00 & 5.86E+00 & 4.76E+00\\\\ 5.00E+02 & 22.58 & 10.39 & 21.6 & 2.00E+00 & 3.25E+00 & 2.83E+00\\\\ 7.07E+02 & 22.88 & 10.64 & 21.6 & 1.41E+00 & 1.80E+00 & 1.68E+00\\\\ \\hline 1.00E+03 & 23.24 & 10.88 & 21.6 & 1.00E+00 & 1.00E+00 & 1.00E+00\\\\ \\hline 1.41E+03 & 23.67 & 11.12 & 21.6 & 7.09E-01 & 5.58E-01 & 5.97E-01\\\\ 2.00E+03 & 24.18 & 11.36 & 21.6 & 5.00E-01 & 3.08E-01 & 3.54E-01\\\\ 2.80E+03 & 24.75 & 11.6 & 21.6 & 3.57E-01 & 1.74E-01 & 2.13E-01\\\\ 4.00E+03 & 25.4 & 11.85 & 21.6 & 2.50E-01 & 9.47E-02 & 1.25E-01\\\\ 5.66E+03 & 26.1 & 12.09 & 21.6 & 1.77E-01 & 5.25E-02 & 7.43E-02\\\\ 8.00E+03 & 26.88 & 12.33 & 21.6 & 1.25E-01 & 2.92E-02 & 4.42E-02\\\\ 1.13E+04 & 27.72 & 12.57 & 21.6 & 8.85E-02 & 1.62E-02 & 2.63E-02\\\\ 1.60E+04 & 28.63 & 12.82 & 21.6 & 6.25E-02 & 8.97E-03 & 1.56E-02\\\\ 2.26E+04 & 29.6 & 13.06 & 21.6 & 4.42E-02 & 4.98E-03 & 9.29E-03\\\\ 3.20E+04 & 30.64 & 13.3 & 21.6 & 3.13E-02 & 2.76E-03 & 5.52E-03\\\\ 4.53E+04 & 31.74 & 13.54 & 21.6 & 2.21E-02 & 1.53E-03 & 3.29E-03\\\\ 6.40E+04 & 32.9 & 13.79 & 21.6 & 1.56E-02 & 8.50E-04 & 1.95E-03\\\\ 9.05E+04 & 34.12 & 14.03 & 21.6 & 1.11E-02 & 4.72E-04 & 1.16E-03\\\\ 1.28E+05 & 35.41 & 14.27 & 21.6 & 7.81E-03 & 2.62E-04 & 6.91E-04\\\\ 1.81E+05 & 36.76 & 14.51 & 21.6 & 5.53E-03 & 1.45E-04 & 4.11E-04\\\\ 2.56E+05 & 38.18 & 14.76 & 21.6 & 3.91E-03 & 8.05E-05 & 2.44E-04\\\\ 3.62E+05 & 39.65 & 15 & 21.6 & 2.76E-03 & 4.47E-05 & 1.45E-04\\\\ 5.12E+05 & 41.19 & 15.24 & 21.6 & 1.95E-03 & 2.48E-05 & 8.63E-05\\\\ \\hline\\hline \\end{tabular} \\end{table} \\end{small}   \\begin{figure} \\epsscale{.80} \\plotone{f1.eps} \\caption{The $p$-$p$ cross-section models: (a) for model A and (b) for model B. Curves are for the total ({\\it{upper solid}}), non-diffractive inelastic ({\\it{dot-dashed}}), elastic ({\\it{dashed}}), all diffractive ({\\it{lower solid}}), and single diffractive ({\\it{dotted}}) processes. Note that model B is made only of non-diffractive inelastic process. Data are for the total ({\\it{circles}}), elastic ({\\it{triangles}}), and single diffraction ({\\it{crosses}}). \\label{fig1}} \\end{figure}  The gamma-ray spectrum from $pp \\rightarrow \\pi^0$ will be presented for 3 cosmic proton spectra: power-law spectrum of index 2.0 representing the acceleration site (see, eg., \\citet{Ellison04}); the local interstellar spectrum (referred to as LIS) obtained from the recent primary cosmic proton measurements (see eg. \\citet{MSOP02}); and a trial broken power-law spectrum for the Galactic ridge region (index=2.5 and 2.2 for above and below $T_p=20$ GeV, respectively).  We also built a reference model, model B, consisting only of the non-diffractive process as all previous models have assumed \\citep{Strong78,SB81,Dermer86,Stecker89}. The non-diffractive cross-section assumed in model B approaches a constant value (see Column 4 of Table 1 and Fig.1b) similar to that assumed in the above references. The computer programs for the model A non-diffractive process are used to calculate the model B non-diffractive process except for the Pythia parameters as will be described later in Sec.5.  In the following sections, we will describe how the total inelastic cross-section is broken down in our model A; relate the Feynman scaling hypothesis and numerical simulation codes in a historic perspective; describe how the scaling violation is incorporated in Pythia; introduce our simulation code of the diffractive interaction; and present the gamma-ray spectra obtained with models A and B. We then compare the predictions of model A, \\citet{SB81}, \\citet{SMR04} and model B with the EGRET gamma-ray spectrum (EGRET Archive) \\footnote{See EGRET archived data at ftp://cossc.gsfc.nasa.gov/compton/data/egret/high\\_level/combined\\_data/\\\\ galactic/counts.g1234\\_30.g001. We have subtracted all point sources listed in the EGRET 3rd Catalog from the data above to make the diffuse gamma-ray spectrum in the Galactic ridge. This point-source subtracted data will be made available in the publication in preparation.}  Finally discussion on the results, our conclusions, and possible implications on $\\bar{p}$, neutrino, $e^+$, and $e^-$ spectra produced by $p$-$p$ interactions will be given. Technical details on models A and B will be given in Appendix. ",
        "conclusions": "We conclude on the analyses presented here that an accurate modeling of the $p$-$p$ interaction (model A) with the diffractive process and the Feynman scaling violation makes the gamma-ray spectrum harder and produces 30$-$80\\% more gamma-rays (Figs.6, 7, and 8) than previous predictions \\citep{Strong78,SB81,Dermer86,Stecker89,Mori97} for incident protons with $T_p>100$~GeV. Combination of the two can explain $\\sim 50$~\\% of the ``GeV Excess'' in the EGRET Galactic ridge spectrum within the conventional cosmic proton and electron spectra as shown in Fig.8.  The above statement is only relative to other $pp \\rightarrow \\pi^0$ production models: the absolute prediction of the Galactic ridge gamma-ray spectrum is contingent on the absolute normalization, or the absolute cosmic ray fluxes, the absolute ISM density, and the absolute radiation field density.  As far as the gamma-ray spectral shape is concerned, the remaining discrepancy (50\\%) requires some modification to the conventional cosmic ray spectra: one possibility is to assume the proton spectrum in the Galactic ridge to be a little harder than that of the solar neighborhood, eg. $\\sim 2.5$ in power-law index as Trail4GR in Fig.8.  We have compared model A critically with data from accelerator experiments (Figs.1a, 3, 4, and 5) and confirmed that important aspects of experimental data are reproduced much better by model A than model B which crudely reflects the Feynman scaling hypothesis.  We believe that all future cosmic $p$-$p$ interaction models must include the diffractive process and incorporate the scaling violation.  Model A with LIS predicts the $\\bar{p}$ flux to be higher by a factor of $\\sim 1.5-2$ relative to model B. We will compare the above prediction with the recent measurements on the $\\bar{p}$ flux \\citep{Orito00, Asaoka02, Basini99, Boezio01} and the prediction by \\citet{MSOP02} in a separate publication (T. Kamae et al. 2005, in preparation).  In this work, we have neglected contributions of $\\alpha$ particles and helium atoms/ions to the gamma-ray spectrum. We note that the abundance of helium atoms/ions is about 7\\% of hydgrogen atoms/ions and so is the $\\alpha$ to $p$ ratio. The spectral index of $\\alpha$ particles is known to be lower by about $0.1$ (see eg., \\citet{Mori97,Schlickeiser02}). Since there are no measurement on the $\\pi^0$ inclusive cross-section for $pHe$, $\\alpha p$, and $\\alpha He$ interactions at high energies, we refer to an estimate on possible deviation from that of $p$-$p$. We take the results obtained on $pd$ and $p$-$p$ interactions by \\citet{DiffDissoc6}.  From the reference, we learn that: the $(\\sigma_T(pd)/\\sigma_T(pp))^2$ remains constant over the incident momentum range of the experiment; and the coherent factor in the diffractive process is less than 10\\%. Since the coherent factor decreases rapidly as the momentum transfer increases, we expect it to be much smaller than 10~\\% for the non-diffractive process. By inference we conclude that the gamma-ray spectra produced by $pHe$, $\\alpha p$, and $\\alpha He$ interactions can safely be represented by those by $p$-$p$ interactions as has been done in this work.  \\begin{figure} \\epsscale{.80} \\plotone{f10.eps} \\caption{Spectra of electrons ({\\it{dashed curve}}) and positrons ({\\it{solid curve}}) produced in 6400 diffractive interactions by protons with $T_p=512$~TeV.} \\end{figure}  The findings of this work make the following interesting predictions on the spectra of other cosmic ray particles (T. Kamae et al. 2005, in preparation): \\begin{itemize} \\item The diffractive interaction and scaling violation (model A) nearly double the TeV gamma-ray yield from interaction of power-law protons with index=2.0, compared to that predicted with model B or our approximation to the scaling model (see Fig.6a). We also expect all decay products of charged pions ($e^+$, $e^-$, $\\nu_e$, $\\bar{\\nu_e}$, $\\nu_\\mu$, and $\\bar{\\nu_\\mu}$) to increase by a similar proportion near the highest end of their spectra. \\item Combination of the inherently low multiplicity of the diffractive process and the charge conservation predicts $\\sim 1.5$ times as many $\\pi^+$ as $\\pi^-$ for $T_p=512$~TeV protons as shown in Fig.10. We hence expect $e^+/e^-$ and $\\nu_e/\\bar{\\nu_e}$ to increase near the highest end of their spectra where the diffractive process contributes most. We note that possible increase of $e^+$ relative to $e^-$ has been reported in the cosmic $e^+$ spectrum above 5~GeV. \\citep{Coutu99}. \\end{itemize}"
    },
    "0410/astro-ph0410567_arXiv.txt": {
        "abstract": "This paper explores three different strategies for the inversion of spectral lines (and their Stokes profiles) using artificial neural networks. It is shown that a straightforward approach in which the network is trained with synthetic spectra from a simplified model leads to considerable errors in the inversion of real observations. This problem can be overcome in at least two different ways that are studied here in detail. The first method makes use of an additional pre-processing auto-associative neural network to project the observed profile into the theoretical model subspace. The second method considers a suitable regularization of the neural network used for the inversion. These new techniques are shown to be robust and reliable when applied to the inversion of both synthetic and observed data, with errors typically below $\\sim$100~G. ",
        "introduction": "\\label{sec:intro}  The analysis of the spectral properties of the light intensity and its polarization state is the basis of modern solar physics. For over two decades, least-squares profile fitting has been the method of choice (see \\citeNP{SN01a}; \\citeNP{dTI03} and references therein). Many different inversion codes, based on a variety of physical models, have been developed and used extensively for the determination of magnetic and thermodynamic conditions in the atmosphere.  While the use of least-squares fitting has important benefits, there is an increasing demand for alternative procedures that are faster and more robust. By robust I mean capable of operating reliably without human intervention on a routine basis. This demand is driven by the development of a new generation of spectro-polarimeters, which will deliver enormous data flows (SOLIS, \\citeNP{K98}; Solar-B, \\citeNP{LES01}; Sunrise, \\citeNP{SBK+03}; DLSP, \\citeNP{SEL+03}).  The attention of solar physicist has turned in recent years towards a new breed of diagnostic techniques based on pattern recognition and machine learning. A considerable number of papers has been devoted to the investigation of these techniques (\\citeNP{RLAT+00}; \\citeNP{SNLAL01}; \\citeNP{CS01}; \\citeNP{SLA02}; \\citeNP{SN03}; \\citeNP{dTILA03}; \\citeNP{SN04d}). Most of those works deal with the Principal Component Analysis, which is a series expansion of the spectra. Another line of research explores the use artificial neural networks (ANNs), which show considerable promise for the inversion of spectral observations (\\citeNP{CS01}; \\citeNP{SN03}).  The work presented here demonstrates the applicability of ANNs to actual observed data in typical working conditions. The radiative transfer computations that are needed to synthesize training profiles have been carried out using the Milne-Eddington approximation. This is helpful to simplify the problem and to allow for the synthesis of thousands of profiles in a reasonable computing time. All the calculations done for this work, and the CPU times quoted below, were obtained using a Pentium~IV processor running at 1.2~GHz. The ANN inversions require very little computational resolurces both in terms of processor speed and memory storage. The training algorithm, on the other hand, can be very demanding. ",
        "conclusions": "\\label{sec:conc}  This paper shows that ANNs are a viable alternative to least-squares fitting for the routine analysis of large amounts of data. While their accuracy is somewhat lower, the ability to process much larger datasets will probably present advantages for some application. The CPU times required for the ANN inversions presented here are $\\sim$10~seconds, compared to $\\sim$5~hours for the original Milne-Eddington inversion.  Furthermore, one does not have to be concerned with the algorithm finding secondary minima or not converging. Strictly speaking, the ANN inversion is not an inverse problem. It is rather a case of interpolating a multidimensional function that maps a set of gridpoints from the space of spectra into the space of models.  It is important to emphasize that pattern recognition techniques are not meant to replace traditional least-squares fitting algorithms, but to complement them. For detailed studies of a smaller data subset or individual profiles, or if one needs to consider more realistic model atmospheres (with line-of-sight gradients, Non-LTE effects, etc), it is still necessary to use a least-squares inversion."
    },
    "0410/astro-ph0410084_arXiv.txt": {
        "abstract": "The purpose of this letter is to check the quality of different methods for estimating stellar masses of galaxies. We compare the results of (a) fitting stellar population synthesis models to broad band colors from SDSS and 2MASS, (b) the analysis of spectroscopic features of SDSS galaxies \\citep{Kauffmannetal03a}, and, (c) a simple dynamical mass estimate based on SDSS velocity dispersions and effective radii.  Knowing that all three methods can have significant biases, a comparison can help to establish their (relative) reliability. In this way, one can also probe the quality of the observationally cheap broadband color mass estimators for galaxies at higher redshift.  Generally, masses based on broad-band colors and spectroscopic features agree reasonably well, with a rms scatter of only $\\sim 0.2$~dex over almost 4 decades in mass. However, as may be expected, systematic differences do exist and have an amplitude of $\\sim 0.15$~dex, corrleting with \\Ha\\ emission strength.  Interestingly, masses from broad-band color fitting are in better agreement with dynamical masses than masses based on the analysis of spectroscopic features. In addition, the differences between the latter and the dynamical masses correlate with \\Ha\\ equivalent width, while this much less the case for the broad-band masses.  We conclude that broad band color mass estimators, provided they are based on a large enough wavelength coverage and use an appropriate range of ages, metallicities and dust extinctions, can yield fairly reliable stellar masses for galaxies. This is a very encouraging result as such mass estimates are very likely the only ones available at significant redshifts for some time to come. ",
        "introduction": "\\label{sec:introduction}  The stellar mass of galaxies at the present epoch and the build-up of stellar mass over cosmic time has become the focus of intense research in the past few years.  In the local universe, results on the stellar mass function of galaxies were published using the new generation of wide-angle surveys in the optical (Sloan Digital Sky Survey; SDSS, \\citealp{SDSS}; 2dF, e.g.~\\citealp{2dF99}) and near-infrared (Two Micron All Sky Survey; 2MASS, \\citealp{TwoMASS}). \\citet{2dF01} combined data from 2MASS and 2dF to derive the local stellar mass function, \\citet{BMKW03} used the SDSS and 2MASS to the same end.  At $z > 0$, a number of authors studied the stellar mass density as a function of redshift \\citep{BE00,MUNICS3,Cohen02,DPFB03,Fontanaetal03,Rudnicketal03} reaching $z \\sim 3$, while others, using wider field surveys, investigated the evolution of the mass function of galaxies \\citep{MUNICS6,K20-04} to $z \\sim 1.5$.  Generally, the high-redshift work relies on fits of multi-color photometry to a grid of composite stellar population (CSP) models to determine a stellar mass-to-light ratio, since large and complete spectroscopic samples of galaxies are not yet available.  A similar approach was chosen by \\citet{2dF01} and \\citet{BMKW03}, too, at $z \\sim 0$.  Taking advantage of the availability of photometry and spectroscopy for galaxies in the SDSS, \\citet[][K03 hereafter]{Kauffmannetal03a} utilized spectroscopic diagnostics (4000\\AA~Break, \\D, and the \\Hd\\ Balmer absorption line index ) to estimate the mean stellar age and the fraction of stars formed in recent bursts in each galaxy.  By comparison of the colors predicted by their best-fit model to the object's broad-band photometry they determine the amount of extinction by dust and hence the stellar mass-to-light ratio.  The purpose of this letter is to compare the stellar masses determined by this spectroscopic technique to masses obtained from multi-passband photometry and to compare both methods to a simple dynamical estimate of mass. Knowing that none of these methods yields a fiducial (stellar) mass, a comparison helps to establish the (relative) reliability of each method and makes us aware of potential differences between these estimators. Moreover, it can show us whether one can use observationally cheaper estimators as surrogates for more expensive (or unobtainable) ones, which is particularly important when dealing with high-redshift datasets.  Specifically, we want to know how the two estimators compare to each other, if using K-band \\ML\\ yields better masses than using the g-band \\ML\\ which is accessible at high $z$, and how these compare to a simple dynamical mass estimator, $M \\sim \\sigma^2 R_e / G$.  \\begin{figure*}[t] \\centering \\epsscale{0.8} \\plotone{f1a.eps}\\hspace*{0.5cm} \\plotone{f1b.eps} \\caption{\\label{fig:loglik}% Illustration of the model-fitting technique used to estimate stellar masses. The lower panels show the comparison of the best fitting model to the photometric data. The red and blue lines represent the main and burst component, respectively. The green line represents the combined SED. The left hand side shows a young object with a burst component, the right side an older object. The upper panels show projections of the likelihood function onto four planes, age vs.\\ dust in the main component, burst fraction vs.\\ burst extinction, age of the main component vs.\\ burst fraction, and star formation timescale, $\\tau$, vs.\\ burst fraction. The resulting likelihood distributions of \\ML\\ are shown in the upper right panel on each side ($M/L_g$ blue; $M/L_i$ green; $M/L_K$ red). The \\ML\\ of the best fitting model is indicated by vertical lines.} \\end{figure*}  This letter is laid out as follows. In Sect.~\\ref{sec:galaxy-sample} we describe the sample of galaxies we use in this work.  In Sect.~\\ref{sec:deriving-masses} we give a brief overview of how we derive stellar masses by fitting CSP models to multi-band photometry. In Sect.~\\ref{sec:stellmass} we compare these masses to the values in K03 and in Sect.~\\ref{sec:dynmass} to a simple dynamical estimate of mass based on the SDSS velocity dispersions. We also discuss the implications of these comparisons.  We assume $\\Omega_{\\mathrm{M}} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70~\\mathrm{km\\ s^{-1}\\ Mpc^{-1}}$ throughout this work. ",
        "conclusions": ""
    },
    "0410/hep-ph0410114_arXiv.txt": {
        "abstract": " ",
        "introduction": "For some time it has been apparent that the inferred values of the cosmological baryon and dark matter densities are strikingly similar. For example, the latest WMAP-determined range for the dark matter density, $0.129 > \\Omega_{\\rm{CDM}} h^2 > 0.095$, is within a factor of a few of the combined WMAP and big-bang nucleosynthesis determined value of the baryon density, $0.025 > \\Omega_{\\rm{b}} h^2 > 0.012$ \\cite{WMAP,BBN}. In the vast majority of models of the early universe, the cosmological baryon and dark matter densities are independently determined. The surviving baryon density is set by a baryon asymmetry generated during baryogenesis, and thus depends upon unknown, and {\\it a priori} unknown and possibly small CP-violating phases, as well as unknown baryon-number violating dynamics. In contrast, the dark matter density is set by the `freeze-out' of the interactions that keep the dark matter in equilibrium, and is independent of the dynamics of baryogenesis. In particular, although weakly-interacting massive particles at the TeV-scale naturally have a relic density of $\\cal{O}$(1) times the critical density, this is not at all the case for baryons.  The Boltzmann equations show the size of the baryon relic density in the absence of an asymmetry is ${\\cal O}(10^{-11})$. Thus the comparative closeness of the baryon and dark matter densities poses a puzzle. One possible approach to this problem is to more closely integrate the dynamics of baryogenesis with that of the origin of dark matter. In particular, it is natural to consider models where the dark matter and baryon sectors share a quantum number, either continuous or discrete, which provides a relation between their surviving number densities and thus energy densities.\\footnote{An alternate approach is to invoke a form of the Anthropic Principle, see, for example, Ref. \\cite{wilczek}. Other ideas are explored in Refs. \\cite{Kus,Enq,Oak2,Oak,farrar}.}  Specifically, in this letter we consider models of dark matter possessing a particle-antiparticle asymmetry where this asymmetry strongly effects the dark matter density, and through the electroweak (EW) anomaly, determines the baryon asymmetry, thus naturally linking $\\Omega_{\\rm b}$ and $\\Omega_{\\rm CDM}$. (A noteworthy early attempt along these lines which shares some features with our model is contained in Ref.~\\cite{kaplan}.) The shared quantum number clearly requires that the dark matter particle not be it's own antiparticle, so the putative dark matter candidate cannot be the standard LSP neutralino of the broken Minimal Supersymmetric Standard Model (MSSM).  In fact, we will show in sections~3 and 4 that sneutrinos can play the role of such dark matter in a previously studied variant of the MSSM in which the light neutrino masses result from higher-dimensional supersymmetry-breaking terms~\\cite{ahhmsw,borz,sw, mrw,hmrw}.  This model preserves all the successes of the MSSM, such as stability of the weak scale and unification of gauge couplings, while being, at least in part, testable at the LHC.   Before we proceed to the calculation of the relic dark matter and baryon densities and the details of our model, it is useful to consider some aspects of the idea of a shared quantum number determining the ratio of baryon to dark matter.  First, consider the simplified case in which, to a very good approximation, the dark matter and baryon sectors cannot exchange their conserved quantum number after its first production.  As an example, consider the situation in which the asymptotic baryon number $B=1$ states of the Standard Model (SM) have global charge, $q$, while the lightest asymptotic states (with mass $m_{\\rm dm}$) in a hidden sector carry charge, $Q$.  Then conservation of global charge implies that $q\\left(n_{\\rm b} - n_{\\rm \\bar{b}}\\right) =-Q\\left( n_{\\rm dm} - n_{\\rm \\bar{dm}} \\right)$, where the $n$'s are the number densities of the indicated (anti)particles. Further assume that interactions in both sectors are strong enough such that almost all antiparticles are eliminated by annihilation with their particles (assuming $Q/q<0$ for simplicity). This implies \\beq n_{\\rm b} =  c n_{\\rm dm} ~~{\\rm with}~~ c = |Q/q| , \\label{naive} \\eeq which in turn leads to \\beq \\frac{\\Omega_{\\rm b}}{\\Omega_{\\rm dm}} = \\frac{m_{\\rm b}}{m_{\\rm dm}} \\frac{n_{\\rm b}}{n_{\\rm dm}} = c \\frac{m_{\\rm b}}{m_{\\rm dm}} . \\label{naive2} \\eeq Thus the energy densities are indeed related, but the ratio is only naturally ${\\cal O}(1)$ if the ratio of the baryon to dark matter masses is not too small.  It is important to note that this disfavours models of the above type where the dark matter candidate is a `hidden sector' particle and favours a particular class of models where the dark matter candidate has weak scale mass, its mass arising either from electroweak symmetry breaking, or from the dynamics which drives electroweak breaking, such as supersymmetry softly broken at the weak-scale.\\footnote{ In extensions of the SM such as the MSSM where many gauge-non-neutral fermionic and scalar states also gain mass at this scale, the beta-function coefficient for $SU(3)_{\\rm QCD}$ sharply increases as energies are reduced below the weak scale, and the confinement scale and thus baryon mass is naturally not very much smaller than $m_{\\rm weak}$.}  As a consequence, the naive statement of Eq.(\\ref{naive}) also requires modification as the approximation of negligible interactions between the two sectors, and thus negligible exchange of the shared conserved quantum number, does not hold in realistic models.  In this situation, the ratio of the conserved quantum number stored in the two sectors, $c$, is not simply given by $Q/q$, but is determined by the `chemical' equilibrium conditions between the two sectors just before the freeze-out of the relevant interactions.  We will further discuss this point in section 2.  Moreover, two other alterations to the naive relation Eq.(\\ref{naive2}) can be present.  First, the shared quantum number may only be conserved modulo $n$, in which case `self' annihilations may occur, and second, it is not always the case that dark matter interactions are efficient enough such that $n_{\\rm \\bar{dm}} \\ll n_{\\rm dm}$. Nevertheless, as we argue in the next section, a careful consideration of such models shows that a version of the naive statement Eq.(\\ref{naive2}) can indeed hold. ",
        "conclusions": "In the standard freeze-out calculation for a weakly interacting dark matter relic, there is little reason to expect a density of dark matter which is similar to the density of baryons, short of anthropic arguments. In this letter, however, we have presented an explanation for the similarity between these two quantities. Our solution introduces an asymmetry between dark matter particles and anti-particles which is related to the baryon-antibaryon asymmetry. This leads to a natural dark matter relic density of the same order of magnitude as the baryon density.  As an example, we considered a mixed sneutrino dark matter candidate which transfers its particle-antiparticle asymmetry to the baryons through the electroweak anomaly. We carry out the relic density calculation for such a candidate and find substantial regions of parameter space in which the observations of WMAP can be satisfied. With no asymmetry, only a narrow strip of parameter space can provide the observed relic density.  \\vskip 0.05in \\begin{center} {\\bf Acknowledgments} \\end{center} \\vskip0.05in We wish to thank Yuval Grossman for a useful correspondence. SW is supported by PPARC Studentship Award PPA/S/S/2002/03530."
    },
    "0410/astro-ph0410037_arXiv.txt": {
        "abstract": "Wide-field imaging from space should not forget the dispersive dimension. We consider the capability of space-based imaging with a slitless grism: because of the low near-infrared background in space and the high sky-density of high redshift emission line galaxies this makes for a very powerful redshift machine with no moving parts. A small 1m space telescope with a 0.5 degree field of view could measure redshifts for $10^7$ galaxies at $0.5<z<2$ per year, this is a MIDEX class concept which we have dubbed `The Baryon Oscillation Probe' as the primary science case would be constraining dark energy evolution via measurement of the baryonic oscillations in the galaxy power spectrum. These ideas are generalizable to other missions such as SNAP and DESTINY. ",
        "introduction": "\\def\\Ha{H$\\alpha$} \\def\\micron{$\\mu$m}  Slitless dispersive imaging has a long history extending back to the early Schmidt plate objective prism surveys. The set up is an imaging system, a dispersing element (ideally located at a pupil of the imaging system to avoid aberrations, but this is not essential in slow beams) and a filter to delimit the bandpass.  In essence the object forms its own slit due to its angular size, to first order the background over the field of view is uniform and is equal to the undispersed background through the same filter. From the ground this means the background remains bright when the object's light is dispersed and so one takes an enormous signal-to-noise (S/N) penalty.  Because of this such surveys have generally been restricted to bright emission line galaxies in the local Universe (see the KISS survey\\cite{Weg03} for a recent example). For unresolved emission lines the S/N is independent of the spectral resolution of the disperser.  A low background environment greatly increases the sensitivity of slitless surveys. In space the background is dominated by solar scattered light from zodiacal dust and at 1.6 microns is $\\sim 1000\\times$ darker than on the ground \\cite{JWST} (there is some dependence on ecliptic latitude).  On the Hubble Space Telescope both the ACS optical camera and the NICMOS near-infrared camera are equipped with grisms for slitless spectroscopy. This mode has been used on ACS for spectroscopy of Type 1a supernovae to $z=2$ \\cite{Riess04} and on NICMOS for detection of an abundant population of $z>1$ \\Ha\\ emission line galaxies \\cite{Mc99}. The strong evolution of the cosmological star-formation rate out to $z=1$ means that their are many more bright emission line galaxies than would otherwise be the case\\cite{Hop00}. The space background peaks at 0.6\\micron\\ and so as one follows \\Ha\\ to high-redshift both the strong evolution in the source population and the diminishing background act to counteract the cosmological dimming for the S/N. Going from $z=0.2$ to $z=2$ these effects work to give us a S/N boost of a factor of 30 and make the high-redshift regime accessible to small space telescopes. ",
        "conclusions": "We have outlined a concept (BOP) for a near-infrared space mission using slitless spectroscopy to do wide area redshift surveys an order of magnitude larger than existing ground based surveys such as SDSS and 2dFGRS in both number and volume and probing high-redshift ($0.5<z<2$). It represents a factor of ten improvement on what could be done on the ground. Instrumentally it represents a very simple hardware system which by taking advantage of the low background in space has no requirement for complex fiber or slit configuring machinery.   Although lower backgrounds could be achieved with wide field spectroscopic fiber or slit designs that have recently been proposed\\cite{Mcc04} these technologies are immature by spaceflight standards and as such must be considered risky at the present time.  BOP would serve diverse science goals: some other examples include studying the evolution of galaxy clustering, measuring the evolving luminosity function of star-forming galaxies as a function of metallicity and environment, providing redshift depth information for deep wide imaging surveys and searching for Ly$\\alpha$ emission from young galaxies at $z>7$. We have presented in detail BOP's ability to constrain dark energy using baryon oscillation measurements. It would deliver precision on the equation of state comparable to the SNAP mission but via a completely independent technique.  The BOP concept is generalizable to any wide-field space mission as long as a grism could be inserted in to the imaging system.  We argue that in any such space mission the science-to-cost ratio of adding such a grism is enormous, and should be considered.  \\begin{ack} We would like to thank Chris Blake for assistance in producing Figures~\\ref{fig:pk} and \\ref{fig:wz} and Bob Woodruff from Lockheed Martin for useful design discussions. KG acknowledges generous funding from the David and Lucille Packard foundation. \\end{ack}"
    },
    "0410/astro-ph0410201_arXiv.txt": {
        "abstract": "FUSE observations have been instrumental in advancing our understanding of the physical properties and behavior of active binary systems, including cataclysmic variables (CVs) and X-ray binaries (XRBs). FUSE data have allowed observers to study: accretion disks, magnetically-channeled accretion flow, and white dwarf accretors, and how these respond to accretion fluctuations and disk outbursts; the role of binary evolution in determining system properties; the vertical and azimuthal structure in disks and disk winds; and what the variations in active binary properties reveal about accretion physics in compact systems.  Results of FUSE observations of active binaries are reviewed here. ",
        "introduction": "Active binaries are interacting binary systems ($P_{orb} \\simeq 1$ -- 20~hr) in which a donor star transfers mass to a compact object.  In cataclysmic variables (CVs), the compact object is a white dwarf (WD), while in X-ray binaries (XRBs), the compact object is a neutron star or a stellar-mass black hole. In CVs and low-mass XRBs, mass is transferred via Roche lobe overflow of the donor star and in most cases is accreted onto the compact object through an accretion disk. Active binaries are excellent test beds for the study of binary evolution, SN Type Ia progenitors, and accretion and relativistic physics in nearby, unembedded objects.  In the FUV, active binaries show emission from the WD, the channeled accretion flow (in magnetic systems), the inner accretion disk ($\\leq20 R_{WD}$), and accretion disk winds. The dominant FUV sources and the system morphologies and behaviors vary dramatically in active binaries as a function of compact object type, orbital period, evolutionary history, mass accretion rate, and viewing inclination.  To date, more than 90 active binaries have been observed by FUSE: 24 dwarf novae, 25 non-magnetic novalikes, 12 intermediate polars, 17 polars, 4 super-soft XRBs, 5 high-mass XRBs, 2 low-mass XRBs, and 3 unclassified systems.  Below, we summarize some of the scientific results obtained from these observations and discuss future work. ",
        "conclusions": "After five years, FUSE has observed most of the bright CVs and all of the bright XRBs available to the telescope.  In doing so, it has revealed much about the properties and behavior of active binaries: the response of disks and WDs to accretion events; the role of binary evolution in setting system abundances; the importance of vertical structure in disks and disk outflows; and the strengths and limitations in our picture of the launching mechanism for disk winds.  Future work will concentrate on utilizing the large FUSE database to probe the general properties of active binaries and on expanding our understanding of accretion physics through intensive studies and modeling of single systems.  A survey of the FUV properties of CVs as a function of mass accretion rate, orbital period, evolutionary history, and binary viewing inclination is ongoing \\citep{froning2003}. Also ongoing are large programs devoted to observations of single systems that will undertake detailed analysis and modeling of time variability to determine the structure and properties of the WDs, disk and magnetic accretion regions, and outflows in active binaries.  FUSE has opened a new window into the study of active binaries and will continue to provide cutting-edge scientific observations of these systems for some time to come."
    },
    "0410/astro-ph0410171_arXiv.txt": {
        "abstract": "We present simultaneous $UBV$ polarimetric and photometric observations of the pre-main-sequence binary AK~Sco, obtained over 12 nights, slightly less than the orbital period of 13.6 days. The polarization is a sum of interstellar and intrinsic polarization, with a significant intrinsic polarization of 1\\% at 5250\\AA, indicating the presence of circumstellar matter distributed in an asymmetric geometry. The polarization and its position angle are clearly variable on time scales of hours and nights, in all 3 wavelengths, with a behavior related to the orbital motion. The variations have the highest amplitudes seen so far for pre-main-sequence binaries ($\\approx$1\\%, $\\approx30$\\arcdeg) and are sinusoidal with periods similar to the orbital period and half of it. The polarization variations are generally correlated with the photometric ones: when the star gets fainter, it also gets redder and its polarization increases. The color-magnitude diagram $B-V$, $V$ exhibits a ratio of total to selective absorption $R=4.3$ higher than in normal interstellar clouds ($R=3.1$).  The interpretation of the simultaneous photometric and polarimetric observations is that a cloud of circumstellar matter passes in front of the star, decreasing the amount of direct, unpolarized light, and hence increasing the contribution of scattered (blue) light. We show that the large amplitude of the polarization variations can not be reproduced with a single scattering model and axially symmetric circumbinary or circumstellar disks. ",
        "introduction": "AK~Sco (HD~152404 = IRAS~16514-3648 = HBC~271 = HIP~82747) is a double-lined spectroscopic binary with a period of 13.6093~d, an eccentricity of 0.469 (Mathieu 1994), and a projected separation of 0.143 AU (Jensen et al. 1996a). At a distance of $145^{+39}_{-25}$~pc (Bertout, Robichon, \\& Arenou 1999), it might be an outlying member of the Upper Scorpio subgroup of the Sco-Cen association (Andersen et al. 1989).  AK~Sco is estimated to be $6\\times10^6$ yr old (Andersen et al. 1989) and has strong Li absorption lines (Herbig \\& Rao 1972). However, its classification is unclear. The NIR excess and large irregular light variations are more typical of a classical T~Tauri star (CTTS), but the weak emission lines make it a weak-line T~Tauri star (WTTS) (Andersen et al. 1989). The spectral type is F5 V (Herbig \\& Rao 1972, Hamann \\& Persson 1992) which puts it in a group intermediate between T~Tauris stars and F-type stars according to Th\\'e, de Winter, \\& P\\'erez (1994).  AK~Sco has been known for decades to be highly variable at optical and IR wavelengths. Periods of roughly constant optical brightness at an average level are interrupted at irregular intervals by rapid variations (Bibo \\& Th\\'e 1991), with no correlation (at least permanent) with orbital period (see Jensen \\& Mathieu 1997). The large photometric variations are attributed to variable obscuration by circumstellar (CS) dust condensations (Andersen et al. 1989, Hutchinson et al. 1994).  The submm continuum emission must arise outside the binary orbit, in a circumbinary (CB) disk (Jensen \\& Mathieu 1997) with a predicted dynamically-cleared gap between 0.032 and 0.48~AU. This gap can reproduce the observations if it is filled with optically thin material to reproduce a strong 10 \\micron\\ silicate feature. However, since there is no evidence for a NIR deficit of emission, this gap is not absolutely necessary (Jensen \\& Mathieu 1997).  Numerical simulations of the evolution of circumbinary disks by G\\\"unther \\& Kley (2002) show that in the dynamically-cleared gap in the CB disk, spiral arms bring CS matter from the CB disk to the CS disks around each star, in agreement with Jensen \\& Mathieu's (1997) interpretation of the SED with a gap filled with optically thin material. Accretion is predicted to be synchronized with apastron, but AK~Sco does not exhibit clear periodic excesses of luminosity at those phases.  A model matching spectroscopic and photometric observations includes two identical stars ($M = 1.5M_{\\odot}$, $R = 2.2R_{\\odot}$, $L = 0.9L_{\\odot}$) with an orbital inclination of $\\sim$63\\arcdeg, embedded in a dense cloud with fine structure near the stars and grains larger than those of standard interstellar clouds. No uncertainty is given for the orbital inclination, but eclipses, which have not been observed so far, would be expected if $i \\ga 70$\\arcdeg. The cloud has a dimension similar to the size of the orbit and was determined to be the cause of the irregular variations. A cooler dust component also exists at a larger distance at $\\approx 10$~AU (Andersen et al. 1989).  AK~Sco's SED was also fit by Gregorio-Hetem \\& Hetem (2002) to give stellar radii of 1.69$R_{\\odot}$, a disk radius of 10~AU and an envelope radius of 1600~AU. The optical depth was $\\tau=0.38$ and the inclination 61\\arcdeg. The authors also classify AK~Sco as a young main sequence star.  Alencar et al. (2003) have also used photometric and spectroscopic observations to find the physical parameters for the two near-identical stars: $M = 1.35\\pm0.07M_{\\odot}$, $R = 1.59 \\pm 0.35 R_{\\odot}$, with an orbital inclination 65\\arcdeg $< i < $ 70\\arcdeg. The dust obscuration was also computed and reveals the existence of substructure at a scale of a stellar diameter. One side of the orbit would also be more obscured than the other.  Polarimetry of binaries, spectroscopic or visual, can be very useful to learn about the geometry of their systems (see the recent review by Manset 2004, in preparation).  In a recent series of papers (Manset \\& Bastien 2000, 2001; hereafter Papers I and II), we have studied numerically the polarization variations produced by electrons and dusty envelopes surrounding binary stars, and we have also developed tools to compare observations of pre-main sequence (PMS) binaries with analytical and numerical models and applied them to about two dozens PMS binaries (Manset \\& Bastien 2001b, 2002, 2003; hereafter Papers III, IV, and V).  In this paper, we present simultaneous polarimetric and photometric observations of AK Sco (Sec. 2), use the tools developed earlier to analyze these data (Sec. 3), and compare with the photometric data (Sec. 4). Finally, we compare our results with those of other PMS binaries (Sec. 5) and with numerical simulations (Sec. 6). ",
        "conclusions": "We have presented simultaneous photometric and polarimetric observations of a PMS spectroscopic binary, AK~Sco. The mean observed polarization ($N$=27) is 0.70\\% at 132\\arcdeg\\ at 3550\\AA, 0.84\\% at 132\\arcdeg\\ at 4300\\AA, and 1.0\\% at 133\\arcdeg\\ at 5250\\AA. If an estimate of the interstellar polarization is removed, the intrinsic polarization for AK~Sco is still significant, 0.70\\% at 108\\arcdeg\\ at 3550\\AA, 0.83\\% at 109\\arcdeg\\ at 4300\\AA, and 0.94\\% at 112\\arcdeg\\ at 5250\\AA. This indicates the presence of circumstellar matter distributed in an asymmetric geometry (like a flattened envelope or disk).  AK~Sco's polarization cannot be of interstellar origin only because (1) the position angles for the interstellar and observed polarizations are different, (2) AK~Sco's polarization $P$ is variable, (3) its position angle also varies, as a function of time and wavelength.  The polarization is clearly variable at all wavelengths, on time scales of hours, days, months, and years, which is a typical behavior for PMS stars. The data were obtained on 12 consecutive nights, almost covering the 13.6 day period. Regular variations are seen and have amplitudes of $\\Delta P \\approx 1.0$\\%, $\\Delta \\theta \\approx 30^\\circ$, the highest seen so far for PMS binaries. Those variations are more apparent in the blue (3550$-$4300\\AA) than in the yellow part of the spectrum (5250\\AA) and are compatible with the presence of circumstellar matter located close to the stars.  The variations are not simply double-periodic (as produced by a simple model of 2 equal mass stars in a circular orbit, at the center of an axisymmetric circumbinary envelope made of electrons), but include single-periodic variations. The presence of single-periodic variations could be due to non equal mass stars, the presence of dust grains, an asymmetric configuration of the circumstellar or circumbinary material, or the eccentricity of the orbit ($e=0.469$).  The polarimetric data show some scatter but less than for many of the other PMS binaries that were observed over longer periods of time (Papers~IV and V). We believe this indicates that observations of binaries gathered over a small interval of time compared to the time scale characteristic of modifications in the CS environment will suffer less from epoch-to-epoch variations in the average polarization, and therefore show less scatter and clearer periodic variations.  The polarization varies as a function of wavelength, and is almost always lower in the blue and higher in the red.  The position angle also varies as a function of the wavelength, with different dependencies on the wavelength. Both AK~Sco's polarization level and position angle vary as a function of time and wavelength.  Large amplitude variations are also seen in photometry, with typical amplitudes of 0.5~mag in $V$, 0.12~mag in $B-V$ and 0.20~mag in $U-B$. These photometric variations are also correlated to the polarimetric ones: usually, when the star is bright, its polarization is small, and when the polarization is large, the star is fainter and redder. These photometric and polarimetric variations are compatible with occultation by a CS dust cloud in orbit in a dusty circumbinary envelope. The color-magnitude diagram $B-V$, $V$ exhibits a ratio of total to selective absorption of $R=4.3$, higher than in normal interstellar clouds where $R=3.1$, and characteristic of circumstellar (as opposed to interstellar) matter. There is no evidence for eclipses of one component of the binary by the other, in both the photometric and the polarimetric data.  Since numerical simulations produce high-amplitude variations for circumstellar disks rather than for circumbinary disks, this suggests the presence of CS (and maybe also CB) disks, although the short period does not allow large CS disks. Alternatively, this indicates that there is matter close to the stars, for example in a CB disk whose central cavity is not completely evacuated, or in a stream from the CB disk to the star.  A detailed analysis of the polarimetric data shows that the BME formalism, which could be used to find the orbital inclination, is not able to recover that information. Even though an inclination higher than 63\\arcdeg would help explain the anticorrelation between the star's brightness and its polarization if a cloud of dust passes in the line of sight, the BME inclination between 80 and 90\\arcdeg cannot be regarded as a meaningful result. Therefore, we cannot confirm the orbital inclinations found by Andersen et al. (1989) or Alencar et al. (2003), 63\\arcdeg, and 65\\arcdeg\\ $< i <$ 70\\arcdeg\\ respectively.  The BME results concerning the moments of the distribution are interesting. AK~Sco has values of $\\tau_0 G$ and $\\tau_0 H$ a factor 10 higher than for other binaries.  There are many free parameters, and even with the known inclination and eccentricity, our numerical simulations could not reproduce the observations. The maximum polarization amplitude we can obtain with a single scattering model assuming an axially symmetric density distribution of scatterers around the center of mass (CB envelope) or around a single star (CS disk), is 0.20\\%. Two avenues to explore further are: take into account multiple scattering, and include non-axisymmetric density distributions such as a stream of material falling onto a component of the binary (or both)."
    },
    "0410/astro-ph0410492_arXiv.txt": {
        "abstract": "We have obtained time series optical spectra of the cataclysmic variable PQ And in quiescence.  The spectra show a white dwarf continuum with narrow Balmer emission superimposed over strong Balmer absorption.  The emission lines have blue and red components whose strength changes with time.  An analysis of the H$\\alpha$ emission line implies a short orbital period below the period gap.  Given its lack of accretion disk features, its large and infrequent outbursts, and an orbital period below the period gap, PQ And is probably a low accretion rate object similar to WZ Sge.  In addition, white dwarf model fits imply that PQ And is an excellent ZZ Cet candidate. ",
        "introduction": "PQ And was discovered by McAdam on 21 March 1988 \\citep{iauc4570} at a visual magnitude of 10.  Examination of the Palomar Sky Survey plates showed a precursor object with a blue magnitude between 18-19 and a red magnitude of about 20 \\citep{iauc4577,iauc4579}.  Within 19 days the light curve had declined by 2 magnitudes.  The large outburst amplitude and rapid decline led to the initial classification of a classical nova but spectra taken 3.5 months after the outburst by \\citet{iauc4629} lacked the nebular features typically observed in classical novae.  Instead, the spectra had strong Balmer emission surrounded by broad absorption similar to the quiesent spectra of the dwarf nova WZ Sge.  The strong absorption indicated that the energy distribution was primarily from the white dwarf (WD) and implied an extremely low accretion rate. Recently, it has been shown that systems with this type of spectra are good candidates for containing non-radially pulsating DA WDs or ZZ Cet stars \\citep{WW04a,WW04b}.  WZ Sge is a dwarf nova whose outbursts occur on timescales of tens of years and with amplitudes of 7-8 magnitudes.  Cataclysmic variables (CVs) of this type are also known as TOADS or Tremendous Outburst Amplitude Dwarf novae \\citep{HSC95}.  These characteristics are thought to be due to the very close binary separation and the extremely low accretion rate in these systems.  TOAD orbital periods are all below the period gap and are typically of order 90 minutes.  The mass transfer rates implied by the models are only 10$^{-11}$M$_{\\odot}$ yr$^{-1}$ \\citep{HSC95}.  Given PQ And's similarities to a dwarf nova, \\citet{richter90} searched archival plates for evidence of previous outbursts.  Two other outbursts with maximum magnitudes of $\\sim$ 11 were found on 23 August 1938 and 7 March 1967.  With decades long superoutburst timescales, an outburst amplitude of $\\sim$ 9 mag, and a WZ Sge like quiescent spectrum, PQ And clearly has TOAD characteristics.  In this paper we show that the orbital period {\\em appears to be} below the period gap and that the system likely contains a WD in the instability strip. ",
        "conclusions": "PQ And has characteristics similar to WZ Sge type variable stars.  Its quiescent spectra show little indication of an accretion disk implying a low mass transfer rate.  PQ And had three known superoutbursts of amplitude $\\sim$ 9 magnitudes in the 50 years prior to 1988.  Fits to spectra give P$_{orb}$ = 1.7 hours while the period derived from photometry is slightly longer at 2.1 hours. More observations are required to confirm the true period but the available evidence strongly supports an orbital period below the orbital gap. The best fit WD model gives T$_{eff}$ = 12,000 K, log(g) = 7.7, and a distance of 330 pc.  The WD parameters place it within the ZZ Cet instability strip.  With a quiescent magnitude of $V$ = 19, non-radial pulsations could be detected in PQ And with a mid-sized telescope."
    },
    "0410/hep-ph0410358_arXiv.txt": {
        "abstract": "We present a study of the multiplicities, of the lateral distributions and of the ratio of the electromagnetic to the hadronic components in the air showers, generated by the collision in the atmosphere of an incoming high energy cosmic ray and mediated by the formation of a mini black hole, predicted in TeV scale gravity models with large extra dimensions. The analysis is performed via a large scale simulation of the resulting cascades over the entire range ($10^{15}-10^{19}$eV) of ultra high initial energies, for several values of the number of large extra dimensions, for a variety of altitudes of the initial interaction and with the energy losses in the bulk taken into account. The results are compared with a representative of the standard events, namely the shower due to the collision of a primary proton with a nucleon in the atmosphere. Both the multiplicities and the lateral distribution of the showers show important differences between the two cases and, consequently, may be useful for the observational characterization of the events. The electromagnetic/hadronic ratio is strongly fluctuating and, thus, less decisive for the altitudes considered. ",
        "introduction": "In the almost structureless fast falling with energy inclusive cosmic ray spectrum, two kinematic regions have drawn considerable attention for a long time \\cite{explanations}. These regions are the only ones in which the spectral index of the cosmic ray flux shows a sharper variation as a function of energy, probably signaling some ``new physics'', according to many. These two regions, termed the {\\em knee} and the {\\em ankle} \\cite{uhecr} have been puzzling theorists and experimentalists alike and no clear and widely accepted explanation of this unusual behaviour in the propagation of the primaries - prior to their impact with the earth atmosphere - exists yet. A large experimental effort \\cite{auger,euso} in the next several years will hopefully clarify several of the issues related to this behaviour.  While the {\\em ankle} is mentioned in the debate regarding the possible existence of the so called Greisen, Zatsepin and Kuzmin (GZK) cutoff \\cite{gzk}, due to the interaction of the primaries with the cosmic background radiation, the proposed resolutions of this puzzle are several, ranging from a resonant Z-burst mechanism \\cite{zburst} to string relics and other exotic particle decays \\cite{berez, ben, sb, cfp}. The existence of data beyond the cutoff has also been critically discussed \\cite{BW}.  Given the large energy involved in the first stage of the formation of the air showers, the study of the properties of the cascade should be sensitive to any new physics between the electroweak scale and the original collision scale. Especially in the highest energy region of the spectrum, the energy available in the interaction of the primaries with the atmospheric nuclei is far above any conceivable energy scale attainable at future ground-based accelerators. Therefore, the possibility of detecting supersymmetry, for instance, in cosmic ray showers has also been contemplated \\cite{fragfun}. Thus, it is not surprising, that most of the attempts to explain these features of the cosmic ray spectrum typically assume some form of new physics at those energies.  With the advent of theories with a low fundamental scale of gravity \\cite{ED} and large compact or non-compact extra dimensions, the possibility of copiously producing mini black holes (based on Thorne's hoop conjecture \\cite{Thorne}) in collisions involving hadronic factorization scales above 1 TeV has received considerable attention \\cite{DL} \\cite{Sarcevic} and these ideas, naturally, have found their way also in the literature of high energy cosmic rays \\cite{Anchordoqui,pp4} and astrophysics \\cite{DR}. For instance, it was recently suggested that the long known {\\em Centauro} events might be understood as evaporating mini black holes, produced by the collision of a very energetic primary (maybe a neutrino) with a nucleon (quark) in the atmosphere \\cite{Theodore}. Other proposals \\cite{Ewa} also either involve new forms of matter (for example strangelets) or speculate about major changes in the strong interaction dynamics \\cite{White}.  While estimates for the frequencies of these types of processes both in cosmic rays \\cite{Cavaglia,Anchordoqui,Theodore} and at colliders \\cite{Sarcevic} have been presented, detailed studies of the multiplicities of the particles collected at the detectors, generated by the extensive atmospheric air showers following the first impact of the primary rays, are far from covering all the main features of the cascade \\cite{refhtml}. These studies will be useful in order to eventually disentangle new physics starting from an analysis of the geometry of the shower, of the multiplicity distributions of its main sub-components \\cite{CCF} and of its directionality from deep space. For instance, the study of the location of the maxima of the showers at positions which can be detected by fluorescence mirrors \\cite{review_Anchordoqui}, generated as they go across the atmosphere, and their variations as a function of the parameters of the underlying physical theory, may help in this effort \\cite{Cavaglia}; other observables which also contain potential new information are the multiplicities of the various particle sub-components and the opening of the showers as they are detected on the ground \\cite{CCF}. We will focus on this last type of observables.  To summarize: in the context of the TeV scale gravity with large extra dimensions it is reasonable to assume that mini black holes, black holes with mass of a few TeV, can form at the first impact of ultra high energy primary cosmic rays with nucleons in the atmosphere. The black hole will evaporate into all types of particles of the Standard Model and gravity. The initial partons will hadronize and all resulting particles as they propagate in the atmosphere will develop into a shower(s), which eventually will reach the detectors. The nature and basic characteristics of these showers is the question that is the main subject of the present work. What is the signature on the detector of the showers arising from the decay of such mini black holes and how it compares with a normal (not black hole mediated) cosmic ray event, due, for instance, to a primary proton with the same energy colliding with an atmospheric nucleon (the \"benchmark\" event used here). The comparison will be based on appropriate observables of the type mentioned above.  Our incomplete control of the quantum gravity/string theory effects, of the physics of low energy non-perturbative QCD and of the nature of the quark-gluon plasma phase in QCD, makes a fully general analysis of the above phenomena impossible at this stage. To proceed, we made the following simplifying assumptions and approximations. (1) The brane tension was assumed much smaller than the fundamental gravity scale, so it does not modify the flat background metric. It is not clear at this point how severe this assumption is, since it is related to the ``cosmological constant problem'' and to the concrete realization of the Brane-World scenario. (2) The black hole was assumed to evaporate instantly, leading to initial ``partons'', whose number and distributions are obtained semiclassically. No virtual holes were discussed and no back reaction was taken into account. (3) The initial decay products were assumed to fly away and hadronize, with no intermediate formation of a quark-gluon plasma or of a disoriented chiral condensate (DCC). (4) We used standard simulation programs for the investigation of the extensive air showers produced in the cases of interest. To this purpose, we have decided to use the Monte Carlo program CORSIKA \\cite{CORSIKA} with the hadronic interaction implemented in SIBYLL \\cite{SIBYLL} in order to perform this comparison, selecting a benchmark process which can be realistically simulated by this Monte Carlo, though other hadronization models are also available \\cite{QGSJET}. Finally, (5) a comment is in order about our selection of benchmark process and choice of interesting events. In contrast to the case of a hadronic primary, the mini black hole production cross section due to the collision of a $\\geq 10^3$ TeV neutrino with a parton is of the order of the weak interaction neutrino-parton cross section \\cite{Theodore}. It would, thus, be interesting to compare the atmospheric showers of a normal neutrino-induced cosmic ray event to one with a black hole intermediate state. Unfortunately, at present neutrinos are not available as primaries in CORSIKA, a fact which sets a limitation on our benchmark study. However, it has to be mentioned that neutrino scattering off protons is not treated coherently at very high energy, since effects of parton saturation have not yet been implemented in the existing codes \\cite{CCF}. As shown in \\cite{KK} these effects tend to lower the cross section in the neutrino case.  For a proton-proton impact, the distribution of momenta among the partons and the presence of a lower factorization scale should render this effect less pronounced. For these reasons we have selected as benchmark process a proton-to-air collision at the same depth ($X_0$) and with the same energy as the corresponding ``signal event''. In order to reduce the large statistical fluctuations in the formation of the extensive air showers after the collisions, we have chosen at a first stage, in the bulk of our work, to simulate collisions taking place in the lower part of the atmosphere, up to 1 km above the detector, in order to see whether any deviation from a standard scattering scenario can be identified. Another motivation for the analysis of such deeply penetrating events is their relevance in the study of the possibility to interpret the Centauro events as evaporating mini black holes \\cite{Theodore}. A second group of simulations have been performed at a higher altitude, for comparison.  The present paper consists of seven sections, of which this Introduction is the first. In Section 2 we briefly describe the D-brane world scenario, in order to make clear the fundamental theoretical assumptions in our study. A brief review of the properties of black holes and black hole evaporation is offered here, together with all basic semiclassical formulas used in the analysis, with the dependence on the large extra dimensions shown explicitly. In Section 3 a detailed phenomenological description of the modeling of the decay of the black hole is presented, which is complementary to the previous literature and provides an independent characterization of the structure of the decay. Incidentally, a Monte Carlo code for black hole decay has also been presented recently \\cite{Webber}. We recall that this description -as done in all the previous works on the subject- is limited to the {\\em Schwarzschild phase} of the lifetime of the mini black hole. The modeling of the radiation emission from the black hole - as obtained in the semiclassical picture - (see \\cite{kanti2} for an overview) is performed here independently, using semi-analytical methods, and has been included in the computer code that we have written and used, and which is interfaced with CORSIKA. Recent computations of the greybody factors for bulk/brane emissions \\cite{kanti2}, which match well with the analytical approach of \\cite{kanti1} valid in the low energy limit of particle emission by the black hole, have also been taken into account. Section 4 contains our modeling of the hadronization process. The hadronization of the partons emitted by the black hole is treated analytically in the black hole rest frame, by solving the evolution equations for the parton fragmentation functions, making use of a special algorithm \\cite{CC} and of a specific set of initial conditions for these functions \\cite{kkp}. After a brief discussion in Section 5 of the transformation of the kinematics of the black hole decay event from the black hole frame to the laboratory frame, we proceed in Section 6 with a Monte Carlo simulation of the extensive air showers of the particles produced by taking these particles as primaries. The simulations are quite intensive and have been performed on a small computer cluster. As we have already mentioned, in this work we focus on the multiplicities, on the lateral distributions of the events and on the ratio of electromagnetic to hadronic energies and multiplicities and scan the entire ultra high energy part of the cosmic ray spectrum. Our results are summarized in a series of plots and are commented upon in the final discussion Section 7. ",
        "conclusions": "A rather detailed analysis was presented of some of the main observables which characterize the air showers formed, when a high energy collision in the atmosphere leads to the formation of a mini black hole. We have decided to focus our attention on the particle multiplicities of these events, on the geometrical opening of the showers produced, and on the ratio of their electromagnetic to hadronic components, as functions of the entire ultra high energy spectrum of the incoming primary source. We have shown that in a double logarithmic scale the energy vs multiplicity as well as the energy vs shower-size plots are linear, characterized by slopes which depend on the number of extra dimensions. We have compared these predictions with standard (benchmark) simulations and corrected for the energy which escaped in the bulk, or emitted by the black holes at stages prior to the Schwarzschild phase. Black hole events are characterized by faster growing multiplicities for impacts taking place close to the detector; impacts at higher altitudes share a similar trend, but less pronounced. The multiplicities from the black hole are larger in the lower part of the energy range, while they become bigger for higher energies. We should also mention that, given the choice made for our benchmark simulations, here we have been considering the worst scenario: in a simulation with an impacting neutrino it should be possible to discern between the two underlying events, whether they are standard or black hole mediated. The lateral distributions appear to be the most striking signature of a black hole event. Due to the higher $p_T$s involved, they are much larger than in the benchmark standard simulations.  Our analysis can be easily generalized to more complex geometrical situations, where a slanted entry of the original primary is considered and, in particular, to horizontal air showers, which are relevant for the detection of neutrino induced showers, on which we hope to return in a future work.  To address issues related to the Centauros, one has to concentrate on black holes produced closer to the detectors, since the main interaction in these events is believed to have taken place somewhere between 0 and 500 m above the detectors \\cite{CCT2}. The multiplicities in the showers will of course decrease and are expected, based on the above graphs, to get very close to the values observed, while the ratio of $N_{{\\rm em}}/N_{{\\rm hadron}}$ will also approach the observed values, since the kaons produced by the black hole's \"democratic\" evaporation will not have enough time to decay and will be counted as hadrons. Another crucial question in connection with the mini black hole interpretation of the Centauro events is whether the partonic decay products of the black hole passed through a high temperature quark-gluon plasma phase and whether a Disoriented Chiral Condensate (DCC) was formed, before it finally turned into hadrons. In the present analysis it was implicitly assumed that no such phase was developed. However, if a DCC forms, the prediction for the ratio of the electromagnetic to hadronic component will be drastically different.  Needless to say, the subject of black hole production and evaporation touches upon several different subfields of theoretical and experimental high energy physics and astrophysics. The fields of Cosmic ray physics and of the air shower formation in the atmosphere; of quantum gravity/string theory, of non-perturbative low energy QCD and of the quark-gluon plasma phase, just to mention a few, and this may not even be a complete list. Given our incomplete understanding of all of these, it is clear that a lot more has to be done, before one can safely compare the theory to the observational data. The analysis of diffractive interactions at those energies, in particular, will require considerable attention. Nevertheless, in our view, the issues involved are very important and deserve every effort.  \\vspace{2cm}  \\centerline{\\bf Acknowledgments}  \\vspace{0.5cm}  We thank Raf Guedens, Dominic Clancy, Andrei Mironov, Alexey Morozov and Greg Landsberg for helpful discussions. The work of T.N.T. is partially supported by the grant HRPN-CT-2000-00122 from the EU, as well by the Hellenic Ministry of Education grants ``O.P.Education - Pythagoras'' and ``O.P.Education - Heraklitos''. The work of C.C. and A.C. is supported by INFN of Italy (BA21). The simulations have been performed at the INFN-Lecce computer cluster. C.C. thanks the Theory Group at the University of Crete and in particular E. Kiritsis for hospitality, and the Theory Division at the Max Planck Institute in Munich for hospitality while completing this work.  \\vspace{1.5cm} \\centerline{\\bf POSTSCRIPT} After completing these studies we have been informed by D. Heck that a new version of CORSIKA has been released, which is able to deal with neutrino primaries. As we have discussed in our work, the use of a more realistic benchmark based on neutrino primaries does not invalidate the results of our simulations, but should actually render the differences between mini black hole mediated and standard events even more pronounced.        \\clearpage"
    },
    "0410/astro-ph0410353_arXiv.txt": {
        "abstract": "We present radio observations of the black hole X-ray transient XTE J1720$-$318, which was discovered in 2003 January as it entered an outburst. We analyse the radio data in the context of the X-ray outburst and the broad-band spectrum. An unresolved radio source was detected during the rising phase, reaching a peak of nearly 5 mJy approximately coincident with the peak of the X-ray lightcurve. Study of the spectral indices suggests that at least two ejection events took place, the radio-emitting material expanding and becoming optically thin as it faded. The broad-band spectra suggested that the accretion disc dominated the emission, as expected for a source in the high/soft state. The radio emission decayed to below the sensitivity of the telescopes for $\\sim 6$ weeks but switched on again during the transition of the X-ray source to the low/hard state. At least one ``glitch'' was superimposed on the otherwise exponential decay of the X-ray lightcurve, which was reminiscent of the multiple jet ejections of XTE J1859+226. We also present a $K_s$-band image of XTE J1720$-$318 and its surrounding field taken with the VLT. ",
        "introduction": "\\begin{figure} \\begin{center} \\psfig{file=ir-image.ps,angle=0,width=7cm} \\vspace*{0cm} \\caption{VLT/ISSAC image of the field of XTE J1720$-$318 in the $K_s$-band. The position of the infrared source is marked with a cross.} \\label{ir-image} \\end{center} \\end{figure}  X-ray transient sources are X-ray binary systems, typically with a low mass companion star and a black hole for the compact object (although sometimes a neutron star). They are well-known for dramatic outbursts caused by some form of instability within the accretion flow. There are now $\\sim 40$ observed sources, some of which are recurrent although the majority have still only been observed in outburst once.  The canonical ``soft'' X-ray transient outburst is one which displays a soft blackbody spectrum in low energy X-rays and a ``Fast Rise Exponential Decay'' lightcurve morphology (e.g. Chen, Shrader \\& Livio 1997). With modern X-ray telescopes it has been possible to obtain the spectral and temporal coverage to show that the behaviour is more complicated than this. X-ray transients appear to enter the outburst from an initially hard spectral state (the low/hard state; see e.g. van der Klis 1995). On a timescale of $\\sim$ days--weeks the X-ray source then softens in most cases (e.g. Brocksopp et al. 2002). Some X-ray transients, however, do not soften but remain in the low/hard state throughout the outburst (e.g. Brocksopp et al. 2004, 2001).  All black hole X-ray transient sources which have been observed at radio frequencies have been detected at some point during their outburst\\footnote{The one exception to this is XTE J1755$-$324 but only a single upper limit was obtained and the observation took place during the decay, $\\sim 1$ month after the onset of the outburst (Ogley et al. 1997)} (Brocksopp et al. 2002), although the relationship between the X-ray and radio lightcurves is not a simple correlation. In most of these a jet origin can be inferred, either from synthesis imaging or from study of the spectral index of the radio emission. Low/hard X-ray states are associated with a flat-spectrum compact jet whereas the high/soft (now also known as thermal-dominant) state is not (Fender 2003, 2001; McClintock \\& Remillard 2003). Short-lived discrete ejections are also observed and appear to take place during the transition from the hard to the very high state (or steep power-law state). In particular XTE J1859+226 was observed to make a series of 5--6 radio ejections, each of which was associated with a temporary epoch of X-ray hardening superimposed on a general softening during the decay (Brocksopp et al. 2002).  \\subsection{XTE J1720$-$318}  XTE J1720$-$318 was detected for the first time in 2003 January by the All Sky Monitor on-board the Rossi Timing X-ray Explorer ({\\sl RXTE}/ASM; Remillard et al. 2003). The average 2--12 keV flux was $\\sim130$ mCrab initially, later rising to $\\sim430$ mCrab. The source was hard at first but was later observed by the {\\sl RXTE}/PCA to have a soft spectrum and a low level of high-frequency variability, consistent with the high/soft state (Markwardt \\& Swank 2003; see also e.g. van der Klis 1995). A later {\\sl RXTE}/PCA observation revealed an iron line at 6.2 keV with a FWHM of 2.5 keV and equivalent width of 95 eV (Markwardt 2003).  The radio counterpart was discovered by Rupen et al. (2003) and was found to display variability on a timescale of $\\sim$days; the flux density peaked at 4.9 mJy at 4.9 GHz. The infrared counterpart was discovered by Kato et al. (2003) with $J$, $H$ and $K_s$ band magnitudes in the range 15--17. O'Brien et al. (2003) confirmed the $K_s$ band detection with VLT observations and also refined the radio position of the source to RA=17:19:58.994, DEC=$-$31:45:01.25. Further infared monitoring by Nagata et al. (2003) revealed a decaying lightcurve, superimposed by a secondary maximum $\\sim40$ days after the outburst; the infrared light was interpreted as emission from an irradiated accretion disc.  Finally, XTE J1720$-$318 was observed towards the end of its decay from outburst by {\\sl INTEGRAL} (during pointed observations of H1743$-$322) and was found to be in the low/hard state (Cadolle Bel et al. 2004). ",
        "conclusions": "We have presented radio observations of the 2003 outburst of the X-ray transient XTE J1720$-$318. The radio source was unresolved and reached a peak of $\\sim 5$ mJy. Study of the spectral index showed that, while the source was optically thin throughout, there was some significant variability which we interpret as partially self-absorbed emission at the onset of two discrete ejection events. Following a period of non-detection, the radio source switched on again contemporaneously with the transition to the low/hard state. The broadband spectrum showed that the infrared emission was significantly brighter than the radio synchrotron spectrum; we suggest that the infrared emission was dominated by the accretion disc and/or irradiation of the companion star as expected for a ``soft'' X-ray transient event. The variability during the decay was reminiscent of that of XTE J1859+226; in the case of XTE J1859+226 this variability was interpreted as a sequence of jet ejections. This emphasizes the need for high S/N X-ray hardness and radio monitoring during the ``glitches'' which are often observed superimposed on the decay of the X-ray source."
    },
    "0410/astro-ph0410679_arXiv.txt": {
        "abstract": "The  outflows from comets in orbit around G-type main-sequence stars can be detected when they produce  transient OH absorption lines in the spectrum near 3100 {\\AA} of the host star.  There is only about a  3 ${\\times}$ 10$^{-8}$ probability of detecting an analog to comet  Hale-Bopp orbiting an analog to the Sun.  However,  for young  solar-type stars with very large numbers of  comets, possibly delivering water to terrestrial planets, there is as much as a 1\\% chance that any sufficiently sensitive, randomly-timed observation may detect such transient absorption. ",
        "introduction": "A fundamental goal in modern astronomy is to learn more about the origin and evolution of the solar system and other planetary environments  where life may exist.    Circumstellar dust around main sequence stars was discovered with IRAS (Aumann et al. 1984) and has been studied extensively to learn indirectly about the origin and evolution of larger parent bodies  which might resemble asteroids or comets (Lagrange, Backman \\& Artymowicz 2000; Zuckerman 2001).    To date, however, very little is directly known about these parent bodies.  The goal of this paper is to describe an observational strategy to identify and characterize individual extrasolar comets.  Water-rich asteroids and comets are probably responsible for delivering to the Earth most of its water (Morbidelli et al. 2000).  By studying comets and related objects around young stars, it may be possible to learn more about this process which is fundamental for the development of life as we know it.  There have been a number of previous discussions about possible observational signatures of comets around other stars.  Alcock, Fristrom \\& Siegelman (1986) suggested that the presence of metals in the atmospheres of white dwarfs could be explained by the collision of comets with the star, but this hypothesis is not yet supported by  additional evidence (see Zuckerman et al. 2003). Melnick et al. (2001)  detected gas-phase H$_{2}$O in the outflow from IRC+10216, a carbon-rich mass-losing red giant, which they interpreted as being produced by the sublimation of comets (Stern, Shull \\& Brandt 1990, Ford \\& Neufeld 2001). However, the observed water might  result from chemical reactions in the dense outflow (Willacy 2004).  Jura (2004) has argued that the lack of excess 25 ${\\mu}$m radiation in IRAS data for first ascent red giants means that these stars typically have less than 0.1 M$_{\\oplus}$ of   comet-like Kuiper Belt Objects in orbits at ${\\sim}$ 45 AU.  There is compelling evidence for Falling Evaporating Bodies (FEB's) around ${\\beta}$ Pic, a 12 Myr-old star (Barrado y Navascues et al. 1999, Zuckerman et al. 2001) with a substantial amount of circumstellar dust (Artymowicz 1997, Vidal-Madjar, Lecavelier des Etangs \\& Ferlet 1998). Transient absorption lines which can be attributed to infalling planetesimal also have been detected around other young A-type stars (see, for example, Grady et al. 1996, Roberge et al. 2002). Although Lecavelier des Etangs, Vidal-Madjar \\& Ferlet (1996) and Li \\& Greenberg (1998) have proposed that much of the dust around ${\\beta}$ Pic arises from comets, the observed FEB's   contain a substantial amount of refractory material, and their degree of  resemblance to ice-rich comets in the solar system is  uncertain (Karmann, Beust \\& Klinger 2001, 2003, Thebault, Augereau \\& Beust 2003).  There are very substantial uncertainties regarding  the Oort belt comets in the solar system.  Estimates of the total mass in this region range from   0.6 M$_{\\oplus}$  (Stern \\& Weissman 2001) to 40 M$_{\\oplus}$ Weissman (1996), while there are a  factor of 100 fewer observed ``dormant\"  comets than  predicted with the simplest models (Levison et al. 2002).  Given the major unknowns, in this paper we describe  schematic rather than  detailed models.  We consider  the production of  gas-phase absorptions in the observed spectra of main sequence stars caused by transiting comets which are warm enough that water-ice sublimates (see Beust, Karmann \\& Lagrange 2001).     Previously, the effects of the photometric variations by cometary dust have been computed by Lamers,  Lecavelier des Etangs \\& Ferlet (1997) and Lecavelier des Etangs, Vidal-Madjar \\& Ferlet (1999); here, we consider the possibility of detecting comets by their spectral line absorptions (see also Smith, Black \\& Oppenheimer 1981).   We show results for  OH because it is abundantly produced in comets  and because it can be detected from the ground. ",
        "conclusions": "We have considered the transient OH absorption signature of comets similar to those found in the solar system.  Such events are very rare for stars like the Sun. For young solar-type stars with a substantial 24 ${\\mu}$m excess  where there may be comets  delivering water to terrestrial planets,  a single, sensitive, randomly-timed observation may have a 1\\% chance of detecting OH flowing away from a large comet.  This work has been partly supported by NASA."
    },
    "0410/astro-ph0410486_arXiv.txt": {
        "abstract": "\\noindent We present a simple way to calculate non-Gaussianity in inflation using fully non-linear equations on long wavelengths with stochastic sources to take into account the short-wavelength quantum fluctuations. Our formalism includes both scalar metric and matter perturbations, combining them into variables which are invariant under changes of time slicing in the long-wavelength limit. We illustrate this method with a perturbative calculation in the single-field slow-roll case. We also introduce a convenient choice of variables to graphically present the full momentum dependence of the three-point correlator. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410165_arXiv.txt": {
        "abstract": "We analyzed the spatial distribution of $28500$ photometrically selected galaxies with magnitude $23.5<{\\cal R}_{\\rm AB}<25.5$ and  redshift $1.4<z<3.5$ in 21 fields with a total area of $0.81$ square degrees.  The galaxies were divided into three subsamples, with mean redshifts $\\bar z=1.7$, $2.2$, $2.9$, according to the $U_nG{\\cal R}$ selection criteria of Adelberger et al. (2004) and Steidel et al. (2003). Combining the galaxies' measured angular clustering with redshift distributions inferred from $1600$ spectroscopic redshifts, we find comoving correlation lengths at the three redshifts of $r_0 = 4.5\\pm 0.6$, $4.2\\pm 0.5$, and $4.0\\pm 0.6 h^{-1}$ Mpc, respectively, and infer a roughly constant correlation function slope of $\\gamma=1.6\\pm 0.1$. We derive similar numbers from the $1600$ object spectroscopic sample itself with a new statistic, $K$, that is insensitive to many possible systematics.  Galaxies that are bright in ${\\cal R}$ ($\\lambda_{\\rm rest}\\sim 1500$--$2500$\\AA) cluster more strongly than fainter galaxies at $z=2.9$ and $z=2.2$ but not, apparently, at $z=1.7$.  Comparison to a numerical simulation that is consistent with recent WMAP observations suggests that galaxies in our samples are associated with dark matter halos of mass $10^{11.2}$--$10^{11.8} M_\\odot$ ($z=2.9$), $10^{11.8}$--$10^{12.2} M_\\odot$ ($z=2.2$), $10^{11.9}$--$10^{12.3} M_\\odot$ ($z=1.7$), and that a small fraction of the halos contain more than one galaxy that satisfies our selection criteria. Adding recent observations of galaxy clustering at $z\\sim 0$ and $z\\sim 1$ to the simulation results, we conclude that the typical object in our samples will evolve into an elliptical galaxy by redshift $z=0$ and will already have an early-type spectrum by redshift $z=1$. We comment briefly on the implied relationship between galaxies in our survey and those selected with other techniques. ",
        "introduction": "\\label{sec:intro} Early investigators studied the spatial distribution of galaxies because they hoped to learn about the structure of the universe on the largest scales.  Their influential work was superseded, in the end, by its competition. Problems began with the demonstration that galaxies contained only a small fraction of the matter in the universe. Galaxy formation remained too poorly understood to quell doubts about  how faithfully galaxies traced underlying distribution of dark matter. Other observations improved---gravitational lensing, peculiar velocities, intergalactic absorption lines, and so on---and seemed easier to relate to matter fluctuations. Computers became fast enough to predict the evolution of the large-scale matter distribution from the the initial conditions that microwave-background missions were measuring with increasing precision. As it became clear that the simulations and observations agreed remarkably well, most researchers concluded that the large-scale structure of the universe could be understood completely as the product of gravitational instability amplifying small inflationary perturbations. Galaxies, once believed to be the primary constituent of the universe, came to be seen as small test particles swept into ever larger structures by converging dark matter flows.  The spatial distribution of galaxies remains interesting because it can teach us about galaxy formation. Galaxy formation must be closely related to larger process of gravitational structure formation, since the formation of a galaxy begins with gas streaming into a massive potential well and ends with stars drifting in the cosmic flow. Lessons from 20 years of numerical investigations into structure formation should therefore carry over to the analysis of galaxy clustering. One example is the known positive correlation between clustering strength and mass for virialized dark matter halos. Since galaxies reside in dark matter halos, their clustering strength provides an indication of the mass of the halos that contain them.  The resulting mass estimate depends on the assumptions that the microwave background and other observations have given us reliable estimates of cosmological parameters and of the initial matter power-spectrum, that numerical simulations can correctly trace the evolution of the matter distribution, at least for moderate densities, and that no process can significantly separate baryons from dark matter on Mpc scales---assumptions that are at least as plausible as those behind competing techniques for mass estimation. Another example is the evolution in the clustering strength of a population of galaxies once it has formed.  This is driven solely by gravity and is easy to predict from numerical simulations.  Comparing the clustering of (say) galaxies in the local universe and those at high redshift can therefore suggest or rule out possible links between them.  Few other methods are as useful for unifying into evolutionary sequences the galaxy populations we observe at different look-back times. An excellent review of the history of these techniques has been written by Giavalisco (2002; pp 620-624).  This paper has two aims.  The first is to present measurements of the clustering of UV-selected star-forming galaxies in a redshift range $1.4<z<3.5$ that is only partially explored.  Section~\\ref{sec:data} describes the way we obtained our data, \\S~\\ref{sec:methods} describes and justifies the techniques we used to estimate the spatial clustering in our galaxy samples, and \\S~\\ref{sec:results} presents our estimates of the galaxy correlation function at redshifts $z\\sim 1.7$, $z\\sim 2.2$, and $z\\sim 2.9$. The survey analyzed here is several times larger than its predecessors; the surveyed area of $0.81$ square degrees is roughly 700 times larger than the area analyzed by Arnouts et al. (2002) and 4 times larger than the areas analyzed by Giavalisco \\& Dickinson (2001) and Ouchi et al. (2001). The second aim is to discuss what our measurements imply about the galaxies and their descendants. In~\\S~\\ref{sec:correspondence} we show that the galaxies' correlation functions are indistinguishable from those of virialized dark matter halos with mass $M\\sim 10^{12} M_\\odot$.  In~\\S~\\ref{sec:evolution} we show that the galaxies, dragged by gravity for billions of years, caught in the press of structure formation, would by redshift $z=0$ have a correlation function that is indistinguishable from that of the elliptical galaxies that surround us. Our results are summarized in~\\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} The first part of the paper was concerned with measuring the spatial clustering of large samples of star-forming galaxies at redshifts $z\\sim 1.7$, $2.2$, and~$2.9$. We fit a three-dimensional correlation of the form $\\xi(r)=(r/r_0)^{-\\gamma}$ to the galaxies' angular clustering with standard techniques and to the galaxies' redshift clustering with a new estimator.  The new estimator, $K$, is insensitive to many of the possible systematic biases in our spectroscopic surveys. We reached consistent conclusions about the correlation function with the two approaches, but adopted the angular results since their random uncertainties were somewhat smaller. As given in \\S~\\ref{sec:angularresults}, the best-fit correlation-function parameters from the angular clustering are \\begin{equation} r_0/(h^{-1}{\\rm Mpc}),\\gamma=\\cases{ 4.0\\pm 0.6, 1.57\\pm 0.14& LBG \\cr 4.2\\pm 0.5, 1.59\\pm 0.08& BX  \\cr 4.5\\pm 0.6, 1.55\\pm 0.10& BM \\cr} \\label{eq:correlationlengths} \\end{equation} where BM, BX, and LBG are the names we have given the $U_nG{\\cal R}$ color-selection criteria that we used to find galaxies at $z\\sim 1.7$, $2.2$, and~$2.9$ (\\S~\\ref{sec:observeddata}). The quoted $1\\sigma$ errors include random uncertainties, uncertainties in the integral constraint corrections, and uncertainties in the shapes of the selection functions.  Since the value of $r_0$ depends on the apparent-magnitude limit of the samples, at least in the two higher-redshift bins (figure~\\ref{fig:clust_seg}), the reported values of $r_0$ are somewhat arbitrary. We chose to limit each sample to a fixed range of apparent magnitudes, $23.5<{\\cal R}<25.5$, on the grounds that this resulted in a similar comoving density in each redshift bin.  Different magnitude limits would have resulted in different correlation lengths.  Readers should be aware that the numbers we give are appropriate to the samples as we have defined them, not to the general galaxy population at high redshifts.  The second part of the paper was based on the proposition that WMAP (Spergel et al. 2003) and other experiments have given us reliable measurements of the cosmological parameters and of the shape of the dark matter power-spectrum. This implies that we know what sorts of virialized dark-matter halos existed at different epochs in the past and how their spatial distribution evolved over time.  Since galaxies reside within dark-matter halos, they will have the same correlation function as the halos on large scales (equation~\\ref{eq:xigxih}).  The galaxies' clustering should therefore tell us what sort of halo they reside within. We found a good match (figures~\\ref{fig:xi_halo} and~\\ref{fig:wthet_halo}) between the correlation functions of the galaxies and of halos with threshold masses ranging from $10^{11.5}M_\\odot$ (LBG) to $10^{12.1}M_\\odot$ (BM).\\footnote{When we say (for example) that the galaxies have the same correlation function as halos with threshold mass $M=10^{11.5}M_\\odot$, we mean the subset of halos with mass $M>10^{11.5}M_\\odot$.  The median mass of this subset is $M=10^{11.86}M_\\odot$ in the GIF-LCDM simulation at $z=2.97$.  Halo subsets can be defined with schemes more elaborate than our simple mass threshold (e.g., Kauffmann et al. 1999, Bullock, Wechsler, \\& Somerville 2002), but the differences between the possible schemes are too small to affect our analysis when the subsets they produce are constrained to have the same clustering on large scales.} Equation~\\ref{eq:mass_scales} gives rough $1\\sigma$ limits on the halos' total masses. Similar masses for Lyman-break galaxies have been derived with the same approach by Jing \\& Suto (1998), Adelberger et al. (1998), Giavalisco \\& Dickinson (2000), Porciani \\& Giavalisco (2002), and others. Although the estimated masses were derived solely from the galaxy clustering, they seem reasonable on other grounds. They cannot be much higher. The number density of halos would be lower than the number density of LBGs, for example, if the halo mass were greater than $10^{11.8}M_\\odot$. Such large halo masses would be possible only if significantly more than one LBG resided in the typical halo, and that is something that our observations rule out (\\S~\\ref{sec:correspondence}). Nor can they be much lower.  The halos that contain LBGs would not have enough baryons to form the median LBG stellar mass of roughly $10^{10}M_\\odot$ (Shapley et al. 2001\\footnote{ This mass is for a Baldry-Glazebrook (2003; eq. 3) IMF, and is therefore $1.82$ times lower than the value Shapley et al. calculated for an IMF with a Salpeter slope between $0.1$ and $100M_\\odot$.  Their assumed Salpeter IMF is probably unrealistic since the IMF in the solar neighborhood is known to flatten near $\\sim 1M_\\odot$ and eventually turn over at lower masses. See Leitherer (1998) or Renzini (2004) for further discussion.}) unless their total mass were greater than about $10^{11}M_\\odot$.  We should mention in passing that our best-fit halo masses seem to imply that only a small fraction of the baryons in the halos are associated with the observed galaxies.  For example, the best-fit mass threshold of $10^{11.5}M_\\odot$ for LBGs corresponds to a median total mass of $10^{11.86}M_\\odot$ and median baryonic mass of $1.2\\times 10^{11}M_\\odot$ (for $\\Omega_b/\\Omega_M\\simeq 0.17$, Spergel et al. 2003), roughly ten times larger than the observed stellar masses of LBGs. Since the $10^8$ supernovae that explode during the assembly of the typical LBG's stellar mass will eject roughly $10^8 M_\\odot$ of metals (e.g., Woosley \\& Weaver 1995), enough to enrich at most $1.3\\times 10^{10} M_\\odot$ of gas to LBGs' typical metallicities of $0.4Z_\\odot$ (Pettini et al. 2002), their observed interstellar gas cannot contain a large fraction of the remaining baryons.  These baryons need not be associated with other objects in the halo, however.  They may be locked in dim stars that formed in previous episodes of star-formation (e.g., Papovich, Dickinson, \\& Ferguson 2001), or may have been heated by various processes to undetectably high temperatures. The latter is presumably the case for nearby galaxies, whose ratios of mass in stars and gas to total mass are usually also smaller than the WMAP value $\\Omega_b/\\Omega_M\\simeq 0.17$. The Milky Way, for example, has a total mass of $10^{12}M_\\odot$ (Zaritsky 1999; Wilkinson \\& Evans 1999) and a mass in gas and stars of only $\\sim 8\\times 10^{10} M_\\odot$ (K. Freeman 2004, private communication), yet few would assert that its missing baryons belong to another galaxy in its halo.  After establishing plausible total masses for the halos associated with the galaxies, we considered some of the implications. Our arguments were not new (see, e.g., Moustakas \\& Somerville 2002, Martini \\& Weinberg 2001, Adelberger et al. 1998). They seemed worth revisiting only because our knowledge of the cosmogony, of the local universe, and of high-redshift galaxies has improved so much in the last few years.  We began by estimating the completeness of our surveys from a comparison of the galaxies' number densities to the number densities of halos with similar clustering strength (figure~\\ref{fig:nhalo}).  Similar number densities would imply that almost all of the most massive halos contained a galaxy that satisfied our selection criteria; a much lower galaxy number density would imply that most of the galaxies in massive halos are missed by our survey.  Defining $\\eta$ as the ratio of galaxy to halo number density, we found rough $1\\sigma$ limits of $0.2<\\eta_{\\rm LBG}<1$, $0.6<\\eta_{\\rm BX}<3$, and $0.5<\\eta_{\\rm BM}<2.5$.  These limits were derived from the clustering at radii $r\\simgt 1h^{-1}$ comoving Mpc.  The clustering on smaller scales, sensitive to the possible presence of more than one galaxy in a halo, implies that the upper limits on $\\eta_{\\rm BX}$ and $\\eta_{\\rm BM}$ should be revised downwards to $\\sim 1.25$.  The data appear consistent with the claim of Franx et al. (2003) that our selection criteria find roughly half of the most massive galaxies at $z\\sim 2$. A completeness of order $50$\\% seems plausible to us for other reasons as well.  Shapley et al. (2001) estimate a lifetime for the typical LBG of $\\sim 3\\times 10^8$ yr, for example, which implies that the typical LBG will be bright enough for us to detect for only about half of the time that elapsed between the survey selection limits of $z\\sim 3.4$ and $z\\sim 2.6$.   We considered next the way the clustering of the galaxies would evolve (figure~\\ref{fig:r0_vs_z}).  Analysis of the GIF-LCDM simulation suggested that the correlation length of LBG descendants would be similar by $z\\sim 2.2$ to the correlation length of galaxies in the BX sample and by $z\\sim 1.7$ to the correlation length of galaxies in the BM sample.  The spatial clustering is therefore consistent with the idea that we are seeing the same population at all three redshifts, though the selection criteria's $\\sim 50$\\% incompleteness leaves room for the populations to be distinct and the difference in stellar masses between the LBG ($10^{10}M_\\odot$) and BX ($2\\times 10^{10} M_\\odot$; Steidel et al. 2005, in preparation) samples may not be consistent with continuous star-formation at observed rates through the elapsed time. Turning our attention to lower redshifts, we found that at $z\\sim 1$ our descendants' clustering would most closely match the observed clustering of galaxies that are red and bright and have early-type spectra (figure~\\ref{fig:z1gals}). By $z\\sim 0.2$ the estimated clustering of the descendants suggested elliptical galaxies as the most likely counterparts (figure~\\ref{fig:z0gals}).   The correspondence is especially hard to dispute for the descendants of the brightest and most strongly clustered galaxies in the high-redshift samples.  One conclusion seems difficult to escape:  the descendants of the galaxies in our samples must have significantly larger stellar masses than their high-redshift forebears. Only $\\sim 25$\\% of the total stellar mass in the local universe is found in galaxies with stellar masses smaller than $2\\times 10^{10}M_\\odot$ (Kauffmann et al. 2003), similar to the values in our high-redshift samples, and these faint galaxies are too weakly clustered to have descended from the galaxies we find at $1.4<z<3.5$. Only elliptical galaxies have a spatial distribution consistent with our expectations for the descendants, and the characteristic stellar mass of ellipticals is $10^{11}M_\\odot$ (Padmanabhan et al. 2004). The increase in stellar mass from $z\\sim 2$ to $z\\sim 0$ could have been produced by ongoing star formation or by mergers.  Our results at redshift $z\\sim 1$ may favor the latter, but in any case our findings are strongly inconsistent with traditional notions of monolithic collapse.  It would be unfair to close without mentioning one population of high-redshift galaxies that we have ignored completely. These are the bright ($K\\simlt 20$) near-infrared-selected galaxies. Their reported star-formation rates ($\\sim 200M_\\odot$/yr), correlation lengths ($>9 h^{-1}$ Mpc, Daddi et al. 2004) and stellar masses ($2\\times 10^{11}M_\\odot$; van Dokkum et al. 2004) are extraordinary, far larger than the corresponding values for typical galaxies in our samples.\\footnote{These values are not larger, however, than those for galaxies in our samples with similar $K$ magnitudes.  See, e.g., Shapley et al. 2004 and Adelberger et al. 2005} Galaxies with similarly extreme properties are not a negligible component of the high redshift universe. The shape of the $850\\mu$m background implies that up to a third (Cowie, Barger, \\& Kneib 2002) of all stars could have formed in objects with star-formation rates greater than $\\sim 200 M_\\odot$; halos at $z\\sim 3$ with the large masses $M\\simgt 10^{12.7}M_\\odot$ implied by $r_0\\sim 8$--$9h^{-1}$ Mpc contain in total almost $20$\\% as many baryons as the more numerous and smaller halos with $M\\sim 10^{11.5} M_\\odot$ that contain LBGs; objects with stellar masses $M_\\ast>2\\times 10^{11}M_\\odot$ contain nearly $5$\\% of all stars in the local universe (Kauffmann et al. 2003) and $20$\\% of the stars in local elliptical galaxies (Padmanabhan et al. 2004). No treatment will be entirely complete if it neglects galaxies similar to those found in near-IR surveys. The galaxies we studied are neither the most massive, nor the most rapidly star-forming, nor the most clustered galaxies in the high redshift universe, but it is precisely this that makes them plausible progenitors for the early-type galaxies that surround us today.  \\bigskip \\bigskip We are grateful to the Virgo consortium for its public release of the GIF simulation data and to G. Kauffmann for bringing the data to our attention.  A. Coil and T. Budav\\'ari responded helpfully to our questions about their data. M. Giavalisco, the referee, gave us an insightful report. KLA, AES, and NAR were supported by fellowships from the Carnegie Institute of Washington, the Miller Foundation, and the National Science Foundation.  DKE and CCS were supported by grant AST 03-07263 from the National Science Foundation."
    },
    "0410/astro-ph0410215_arXiv.txt": {
        "abstract": "Although galactic dark matter halos are basically Newtonian structures, the study of their interplay with large scale cosmic evolution and with relativistic effects, such as gravitational lenses, quintessence sources or gravitational waves, makes it necessary to obtain adequate relativistic descriptions for these self--gravitating systems. With this purpose in mind, we construct a post--Newtonian  fluid framework for the ``Navarro--Frenk--White'' (NFW) models of galactic halos that follow from N--body numerical simulations. Since these simulations are unable to resolve regions very near the halo center, the extrapolation of the fitting formula leads to a spherically averaged ``universal'' density profile that diverges at the origin. We remove this inconvenient feature by replacing a small central region of the NFW halo with an interior Schwarzschild solution with constant density, continuously matched to the remaining NFW spacetime. A model of a single halo, as an isolated object with finite mass, follows by  smoothly matching the NFW spacetime to a Schwarzschild vacuum exterior along the virial radius, the physical ``cut--off'' customarily imposed, as the mass associated with NFW profiles diverges asymptotically. Numerical simulations assume weakly interacting collisionless particles, hence we suggest that NFW halos approximately satisfy an ``ideal gas'' type of equation of state, where mass--density is the dominant rest--mass contribution to matter--energy, with the internal energy contribution associated with an anisotropic kinetic pressure. We show that, outside the central core, this pressure and the mass density roughly satisfy a polytropic relation. Since stellar polytropes are the equilibrium configurations in Tsallis' non--extensive formalism of Statistical Mechanics, we argue that NFW halos might provide a rough empirical estimate of the free parameter $q$ of Tsallis' formalism. ",
        "introduction": "A large amount of compelling evidence based on direct and indirect observations: rotation velocity profiles, microlensing and tidal effects affecting satellite galaxies and galaxies within galaxy clusters, reveals that most of the matter content of galactic systems is made up of dark matter (DM). Since the physical nature of DM so far remains uncertain, this issue has become one of the most interesting open problems in astrophysics and cosmology \\cite{KoTu,Padma1,Peac,Ellis,Fornengo}. Among a wide variety of proposed explanations we have: thermal sources, meaning a colissionless gas of weakly interacting massive particles (WIMP's), which can be very massive ($m\\sim 100-200$ GeV) and supersymmetric~\\cite{Ellis} (``cold dark matter'' CDM) or self--interacting less massive ($m\\sim$ KeV) particles~\\cite{scdm,wdm} (``warm'' DM).  The DM contribution dispersed in galactic halos is about 90--95 \\% of the matter content of galactic systems, while visible baryonic matter (stars and gas) is clustered in galactic disks. It is then a good approximation to consider the gravitational field of a galaxy as that of its DM halo (for whatever assumptions we might make on its physical nature), while visible matter can be thought of as ``test particles'' in this field~\\cite{HGDM, Lake}.   Assuming the CDM paradigm, we can distinguish two types of halo models: idealized models obtained from a Kinetic Theory approach, whether based on specific theoretical considerations or on convenient ansatzes that fix a distribution function satisfying Vlassov's equation~\\cite{BT}, or models based on ``universal'' mass density profiles obtained empirically from the outcome of N--body numerical simulations~\\cite{nbody_1,nbody_2,nbody_3}. In this paper we will study the equilibrium configurations that emerge from the latter approach, based on the well known numerical simulations of Navarro, Frenk and White (NFW)~\\cite{nbody_1,LoMa}. It is important to mention that these simulations yield virialized equilibrium structures that reasonably fit CDM structures at a cosmological scale ($\\agt 100$ Mpc), though some of their predictions in smaller scales (``cuspy'' density profiles and excess substructure) seem to be at odds with observations~\\cite{cdm_problems_1,cdm_problems_2}, especially those based on galaxies with low surface brightness (LSB), which are supposed to be overwhelmingly dominated by DM and so well suited to examine the predictions of various DM models~\\cite{vera,LSB}.  We consider in this paper that DM halos are spherically symmetric equilibrium configurations, a reasonable approximation since their global rotation is not dynamically significant~\\cite{urc}. Galactic halos in virialized equilibrium are also Newtonian systems characterized by typical velocities, ranging from a few km/sec for dwarf galaxies up to about $1000-3000$ km/sec for rich clusters. So, why bother with a general relativistic treatment? First, there is a purely theoretical interest in incorporating these important self--gravitating systems into General Relativity, the best available gravitational theory. In fact, important experimental tests of General Relativity are currently and customarily carried on (within a weak field post--Newtonian approach) in Solar System bound Newtonian systems. Secondly, a post--Newtonian description of galactic halos can be, not only useful and interesting, but essential for studying their interaction with physical effects that lack a Newtonian equivalent, such as gravitational lenses or gravitational waves. In fact, the post--Newtonian halo models that we present in this paper can be readily used in lensing studies, or can provide the unperturbed zero order configuration in the examination of the perturbative effect of gravitational waves on galactic halos. Also, a post--Newtonian description is necessary in the study of the interplay between galactic structures and large scale ($> 100\\,$ Mpc) cosmological evolution dominated by a repulsive ``dark energy'', modeled by sources like quintessence and/or a cosmological constant, whose Newtonian description might be inadequate. Finaly,  since galactic halos are customarily examined as Newtonian structures, we feel it is important to show the readership of General Relativity journals how to construct spacetimes, in a post--Newtonian approximation, that are suitable for the description and study of these important self--gravitating systems.  Given the NFW mass density profile, we show in section II how all dynamical variables of the Newtonian  NFW halo can be derived. In section III we construct a post--Newtonian fluid relativistic generalization of a NFW halo, under the assumption that the gas of collisionless WIMPs should satisfy an ``ideal gas'' type of equation of state~\\cite{RKT,Padma2,Padma3}. For isotropic velocity distributions, this assumption allows us to determine the internal energy density by means of the hydrostatic equations themselves. For the case with anisotropic velocities we follow the same procedure with regards to the ``radial'' component of the stress tensor, determining the ``tangential'' stress by a suitable empirical ansatz (section VI--B).  The fact that the NFW spherically averaged mass density profile diverges at the halo center follows from interpolating a fitting formula associated with numerical simulations that have a finite resolution limit at the halo center~\\cite{NSres}. Although this behavior of the density profile does not imply that simulations predict an infinite central density, it is none--the--less an undesirable feature, which we ammend in section IV  by replacing a small region around the center of the NFW halo with a spherical section of a spacetime with constant matter--energy density, {\\it i.e.} a section of the ``interior'' Schwarzschild solution, that is continuously matched to the remaining of the NFW post--Newtonian spacetime.  Galactic halos are hierarchical structures: small halos lie within galaxy clusters, which might be part of superclusters, etc, the asymptotic field of a typical NFW halo should somehow merge with a mean field of a larger substructure, or with a mean cosmological field. However, the dynamical input from a suitable cosmological background is well imprinted in the theoretical design of NFW simulations~\\cite{nbody_1,nbody_2,nbody_3}, while the effect of a cosmological constant on the equilibrium of virialized halo structures is know to be negligible (see \\cite{SussHdez} and references quoted therein).  Hence, at galactic scales the empiric NFW profiles are assumed to be valid only up to the ``virial radius'', a physical ``cut--off'' scale associated with a virialization process~\\cite{Padma1,Peac,BT,Padma2,Padma3}, ignoring altogether their transition to background fields associated with larger structures or to a cosmological background. While a post--Newtonian approach also allows one to impose this virial cut--off scale and to ignore its asymptotic behavior, we show in section IV how well behaved asymptotically flat NFW configurations can be constructed. Also, since any localized self--gravitating system (even if belonging to large substructures) can be approximately described as an isolated system, we also examine an alternative cut--off by matching generic NFW halos to a Schwarzschild vacuum exterior at the virial radius.  In section V we provide the equilibrium equations given in terms of suitable dimensionless variables for the post--Newtonian NFW halos, using the matching with the ``interior'' and ``exterior'' Schwarzschild solutions defined in section IV. Analytic solutions of these equations are obtained in section VI, for isotropic velocities (subsection A) and for a well defined case with anisotropic velocities (subsection B). We show in both cases that (outside the central region) the radial pressure and mass density satisfy approximately a polytropic relation characteristic of stellar polytropes~\\cite{BT}. Even if NFW halos exhibit (in general) deviations from an isotropic velocity distribution, while velocities in polytropes are strictly isotropic, we argue in section VII that the resemblance of outer regions of NFW halos to stellar polytropes might be significant, since virialized self--gravitating systems exhibit non--extensive forms of energy and entropy, and stelar polytropes are the equilibrium states in the application to astrophysical systems of the non--extensive Statistical Mechanics formalism developed by Tsallis~\\cite{Tsallis,PL,TS1,TS2} (see \\cite{Chavanis} for a critical approach to this formalism). ",
        "conclusions": "In the previous sections we have constructed adequate post--Newtonian generalizations for the galactic halo models that emerge from the well known NFW numerical simulations. We have shown how the issues of lack of a regular center (because of interpolating an empiric density profile) and an unbounded halo mass can be resolved by suitable matchings with a section of an interior Schwarzschild solution with constant density, and with a vacuum Schwarzschild exterior. Even if galactic halos are essentially Newtonian systems, we feel it is important for relativists to see how they can also be described and studied in General Relativity within the framework of a post--Newtonian weak field regime. Such a description can be very valuable in studying their interaction with physical effects (gravitational lenses and gravitational waves) and dark energy sources, all of which lack an adequate Newtonian description.  Following our proposal that NFW halos satisfy the ideal gas type of equation of state (\\ref{NRIGES}), we have shown empiricaly (see figure 3) that outside their central core region these halos approximately satisfy the polytropic relation (\\ref{isotpoly}) with $n\\approx 10$. This might be quite significant, since virialized self--gravitating systems are characterized by non--extensive forms of energy and entropy~\\cite{Padma2,Padma3}, and as mentioned before, stellar polytropes are the equilibrium state associated with the non--extensive entropy functional in  Tsallis' formalism~\\cite{Tsallis,PL,TS1, TS2} (see~\\cite{Chavanis} for a critical appraisal). However,  the consequences of this rough polytropic relation should be looked carefully, since stellar polytropes are solutions of Vlassov equation with an isotropic velocity distribution~\\cite{BT}, while NFW halos follow from numerical simulations and exhibit (in general) anisotropic velocity distributions (even if these anisotropies are not too large~\\cite{LoMa}). In the application of Tsallis formalism to self--gravitating collisionless systems~\\cite{TS1, TS2},  the free parameter $q=(2n-1)/(2n-3)$ denotes the departure from the extensive Boltzmann--Gibbs entropy associated with the isothermal sphere (which follows as the limiting case $n\\to\\infty$, or equivalently, as $q\\to 1$). Assuming Tsallis theory to be correct, the empiric verification  provided by figure 3 might indicate that in the region outside  the central core NFW numerical simulations yield self--gravitating configurations that approach an equilibrium state characterized by the Tsallis parameter $q\\approx 1.1$.  While the central cusps in the density profile predicted by NFW simulations seem to be at odds with observations~\\cite{cdm_problems_1,cdm_problems_2,vera,LSB}, there is no conflict between these observations and the $1/x^3$ scaling of the NFW density profile outside the core region. Although the issue of the cuspy cores is still controversial, if it turns out that galactic halos do exhibit flat density cores, their density profiles could adjusted to stellar polytropes  and this might be helpful in providing a better empirical verification of Tsallis' formalism. However, this idea must be handled with due case, since stellar polytropes are very idealized configurations.  Although we have only dealt with NFW halos, the methodology that we have followed here can be applied, in principle, to any Newtonian model of galactic halos. For a deeper study of galactic halo models (NFW, as well as other empiric or theoretical models), it is important to consider a wider theoretical framework, not only using a post--Newtonian approach, but including also the usual thermodynamics of self--gravitation systems~\\cite{Padma2,Padma3}, as well as alternative approaches such as Tsallis' formalism~\\cite{Tsallis,PL,TS1,TS2}. This study might provide interesting theoretical clues for understanding the Statistical Mechanics associated with numerical simulations and/or gravitational clustering. An improvement and extension of the present study of NFW halos are being pursued elsewhere~\\cite{enproceso}."
    },
    "0410/astro-ph0410023_arXiv.txt": {
        "abstract": "During the low optical brightness states of AM Herculis systems (polars) when accretion has declined to a very low value, the underlying magnetic white dwarf photosphere can be modelled without the complication of thermal bremstrahlung and cyclotron emission from the luminous accretion column. The far ultraviolet spectra can be modelled with high gravity solar composition photospheres. In this way, I present new temperatures and the first chemical abundance estimates for the white dwarfs in three selected polars from the IUE NEWSIPS archive. For the white dwarf in V834 Cen with T$_{eff} = 16,000$K, Si/H = 0.1 solar, C/H = 0.5 solar, for BY Cam, T$_{eff} = 17,000$K, Si/H = 0.1, C/H = 5 solar and for RX J1313-32, T$_{eff} = 22,000$K, Si/H = 0.1 solar, C/H = 0.1 solar. The temperature distribution of 24 white dwarfs in polars with known temperatures above and below the period gap is compared with the distribution of the white dwarf temperatures in dwarf novae during quiescence. In both cases, the magnetic white dwarfs in polars are significantly cooler than the non-magnetics. For all CV white dwarfs, magnetic and non-magnetic with T$_{eff} < 12,500$K, 91\\% of the objects (10 out of 11) are magnetics in polars. This suggests that long term accretion heating and cooling of white dwarfs in polars differs from the effects of long term accretion in non-magnetic disk accretors. ",
        "introduction": "The physics of magnetic accretion at the highly magnetic white dwarf surface in polars is poorly known. We do not yet understand the lateral spread (sideways diffusion) and downward diffusion of accreted matter, and the long term heating of the magnetic white dwarf (compared with the heating due to disk accretion).  Yet it is these very systems which offer us tantalizing possibilities for increasing our understanding through studies during their low optical brightness states. This is because, for reasons which remain unexplained, accretion declines to an extremely low rate during these low states and hence the emission contribution from thermal bremstrahlung and cyclotron sources from the accretion column is no longer evident in their far ultraviolet spectra. All that is left to observe is the bare magnetic white dwarf with relatively rapidly cooling polar accretion regions. By using white dwarf model atmospheres to analyze far UV spectra obtained with IUE and HST during low states, new insights are possible on all of the above questions.  In this talk, I present the first results on the chemical abundances of the accreted atmospheres of the white dwarfs in polars starting with the prototype AM Herculis itself, V834 Cen, BY Cam and RX J1313-32. For comparison, the metal abundance derived by Sion et al. (1998) for the magnetic white dwarf in the pre-CV (pre-IP?) V471 Tauri will also be compared. Next, I present the latest information on the low state temperature distribution of the magnetic white dwarfs in polars and compare their temperature distribution with the white dwarf temperatures in dwarf novae. I do this in two ways: number versus T$_{eff}$ distribution histograms and by using the distribution function of magnetic and non-magnetic white dwarf temperatures versus orbital period. ",
        "conclusions": ""
    },
    "0410/astro-ph0410509_arXiv.txt": {
        "abstract": "For the first time we have directly detected  magnetic fields in  central stars of planetary nebulae by means of  spectro-polarimetry with FORS1 at the VLT. In all four objects of our sample we found kilogauss magnetic fields, in NGC\\,1360 and LSS\\,1362  with very high significance, while in Abell\\,36 and EGB\\,5 the existence of a  magnetic field is probable but with less certainty. This discovery supports the hypothesis that the non-spherical symmetry of most planetary nebulae is caused by  magnetic fields in AGB stars. Our high discovery rate demands mechanisms to prevent full conservation of magnetic flux during the transition to white dwarfs. ",
        "introduction": "The reason why more than 80\\%\\ of the known planetary nebulae (PNe) are mostly bipolar and not spherically symmetric (Zuckerman \\&\\ Aller 1986, Stanghellini et al. 1993) is barely understood. A review on observational and theoretical studies of the shaping of planetary nebulae is given by Balick \\&\\ Frank (2002). It is possible that magnetic fields from the stellar surface are wrapped up by differential rotation so that the later post-AGB wind will be collimated into two lobes Garc{\\'i}a-Segura et al.\\ (1999). Another scenario says that  magnetic pressure at the stellar surface plays an important role driving the stellar wind on the AGB (Pascoli 1997). The idea that  magnetic fields are important has been supported by the detection of polarization in radio data of circumstellar envelopes of AGB stars (Kemball \\&\\ Diamond 1997, Szymczak \\&\\ Cohen 1997, Vlemmings et al. 2002). However, until now no magnetic fields have ever directly been detected in central stars of PNe. ",
        "conclusions": "We have detected magnetic fields in 50\\%-100\\%\\ of our small survey for magnetic fields in central stars of planetary nebulae, depending on how conservatively  the criteria for statistical significance are set. This provides very strong support for theories which explain the non-spherical symmetry (bipolarity) of the majority of planetary nebulae by magnetic fields. In this first survey we have not performed a cross check with any spherically-symmetric nebulae, although we have submitted a proposal for follow-up observations.  Although based on only four objects, our extremely high discovery rate demands that magnetic flux must be lost during the transition phase between central stars and white dwarfs: if the magnetic flux was fully conserved, our four central stars will have fields between 0.35 and 2\\,MG when they become white dwarfs, deduced from the atmospheric parameters and the mass-radius relation of Wood (1994). Although the number of white dwarfs with magnetic fields is still a matter of debate, with a range between about 3 and 30\\%, even the latter value, which includes objects with kG field strengths (Aznar Cuadrado et al. 2004), is far off our high number. Liebert et al. (2003) quantified the incidence of magnetism at the level of $\\sim$ 2\\,MG or greater to be of the order of $\\sim$10\\%. This argument would not change by much if we consider that we have so far only looked at central stars with non-spherical symmetric nebulae. An almost 100\\%\\ probability of magnetic fields larger that 100\\,kG can be excluded by the data from the SPY survey (Napiwotzki et al. 2003) as well as the sample from Aznar Cuadrado et al. 2004). It is also worth mentioning that our central stars have typical white dwarf masses (0.48-0.65\\mbox{\\,$\\rm M_{\\odot}$}) and are not particularly massive.  White dwarfs with MG fields tend to be more massive than non-magnetic objects (Liebert 1988).  If the magnetic field is located deep in the degenerate core of the central star, it is very difficult to imagine a mechanism to destroy the ordered magnetic fields. Therefore, it would be more plausible to argue that the magnetic field in the central stars is present mostly in the envelope where it can be affected by convection and mass-loss. For central stars hotter than 100\\,000\\,K we do, however, not expect convection; only in the central star of EGB\\,5 we cannot exclude such a mechanism.  If we assume that the magnetic fields are fossil and magnetic flux was conserved until the central-star phase, we estimate that the field strengths on the main sequence were 9-50\\,G, which are not directly detectable. Therefore, our measurement may indirectly provide evidence for such low magnetic fields on the main sequence.  Polarimetry with the VLT has led to discovery of magnetic fields in a large number of objects in the final stage of stellar evolution: white dwarfs (Aznar Cuadrado et al. 2004), hot subdwarf stars O'Toole et al. (these proceedings), and now in central stars of planetary nebulae. Although we have now provided a good basis for the theoretical explanation of the planetary nebula morphology -- which can more quantitatively be correlated with additional observations in the future --  new questions about the number statistics of magnetic fields in the late stages of stellar evolution have been raised. The full details of our analysis can be found in Jordan et al. (2005)."
    },
    "0410/hep-th0410234_arXiv.txt": {
        "abstract": "In the context of brane world scenario, cosmic superstrings can be formed in D-brane annihilation at the end of the brane inflationary era. The cosmic superstring network has a scaling solution and the characteristic scale of the network is proportional to the square root of the reconnection probability. ",
        "introduction": "Cosmic superstrings have recently gained a lot of interest. They are expected to be formed in the context of a brane world scenario and it was advertised that they lead to observational predictions which might distinguish them from the gauge theory solitons, we are most familiar with.  Cosmic strings are formed in a large number of Spontaneous Symmetry Breaking (SSB) phase transitions. Suppose we have the symmetry breaking of a group $G$ down to a sub-group $H$ of $G$. The formation of cosmic strings is related to the first homotopy group $\\pi_1(G/H)$ of the vacuum manifold ${\\cal M}={G/ H}$, i.e., if ${\\cal M}$ is not simply connected then cosmic strings form. In an exhaustive study of symmetry breaking schemes from a large gauge group down to the standard model gauge group, Jeannerot \\etal~\\cite{jrs} have shown that in the context of supersymmetric grand unified theories, cosmic strings generically appear at the end of a hybrid inflationary era. However, the measurements of the Cosmic Microwave Background (CMB) temperature anisotropies impose severe constraints on the cosmic strings contribution to the CMB data, as it was shown by Bouchet \\etal~\\cite{bouchet}, and more recently by Pogosian \\etal~\\cite{p1,p2}.  Thus, CMB measurements can impose upper limits to the value of the dimensionless parameter $G\\mu\\sim \\eta^2/m_{\\rm Pl}^2$, where the SSB scale $\\eta$ determines the critical temperature $T_{\\rm cr}$ for the phase transition leading to string formation, and $G=1/m_{\\rm Pl}^2$ is Newton's constant. Recently, Rocher and Sakellariadou~\\cite{jm1,jm2} have shown that F-term as well as D-term supersymmetric hybrid inflation, accompanied by cosmic string formation at the end of the inflationary era, are still compatible with CMB measurements, provided the couplings and the mass scales are tuned within acceptable limits.  As the string network evolves, a number of string intersections take place, and this, with a given probability, can lead to an exchange of partners. String intercommutations lead to the creation of kinks along the long strings, which implies that long strings are not straight but they develop wiggles, and thus they emit gravitational radiation, as it was calculated by Sakellariadou~\\cite{ms1}. In addition, the loops which are chopped off from certain string intersections and self-string intercommutations also emit gravitational waves. Gravitational radiation is conventionally believed to be the most efficient mechanism for string decay. This issue is not settled since Vincent \\etal~\\cite{v1,v2} have shown that the string network looses energy directly into scalar and gauge radiation. Numerical studies by Vincent \\etal~\\cite{v1} have for the first time shown that particle production, and not gravitational radiation, is the dominant energy loss mechanism for cosmic string networks. Whether or not gravitational radiation is the most efficient mechanism for cosmic string decay is a very important issue, which however we will not address here and we will come to this question later, in Ringeval and Sakellariadou~\\cite{rs}.  In the context of brane world scenario one naturally expects an era of brane inflation, as proposed by Dvali and Tye~\\cite{dt}. The end of brane inflation is taken place when, either a brane and an anti-brane, or a pair of branes oriented at angles, collide. The inter-brane separation, which is an open string mode, plays the r\\^{o}le of the inflaton field, while the inflaton potential emerges from the exchange of closed string modes between the branes.  As the inter-brane separation becomes smaller than some critical value, an open string mode stretching between the branes becomes tachyonic and the rolling of the tachyon field describes the decay of the pair of branes. This leads to heating of the universe and the beginning of the big bang. Brane annihilation can lead to the formation of lower dimensional D-branes\\footnote{Due to brane intersections leading to reconnection and unwinding, Durrer \\etal~\\cite{rmm} argue that all D$p$-branes of dimension $p>3$ disappear and that one of the stable  3-dimensional branes plays the r\\^ole of our Universe.} which are seen as topological defects by brane observers. As it was shown by Sarangi and Tye~\\cite{st}, cosmic superstrings are copiously produced during brane collisions.  These cosmic superstrings are D$p$-branes with $(p-1)$ dimensions compactified. It is important that neither monopoles, nor domain walls, nor textures, are produced during this process.  Cosmic superstrings share a number of properties in common with the familiar cosmic strings. It is however important to identify their differences which may lead to distinctive observational signatures.  In what follows we examine the behavior of the characteristic scale of the string network with respect to the string reconnection probability. We reach conclusions which differ from results by Dvali and Vilenkin~\\cite{dv} and Jones \\etal~\\cite{jst}.  This may lead to important consequences regarding the observational signatures of cosmic superstrings and their differences with respect to their solitonic analogues. ",
        "conclusions": "Cosmic superstrings are expected to be the outcome of brane inflation in the context of a brane world scenario. These objects are quite similar to their solitonic analogues, even though they have some distinguishable features which may lead to different observable signatures. The main difference between cosmic superstrings and cosmic strings is the fact that the former live in extra dimensions. This results to a smaller reconnection probability, which will alter the subsequent evolution of a superstring network as compared to the most familiar string network evolution.  We have studied numerically the evolution of a cosmic superstring network in a Minkowski space. We have found that the network reaches a scaling solution, which means that the characteristic length scale of the string network, which gives the typical curvature radius of long strings and the characteristic distance between long strings, grow both proportional to the Hubble radius $1/H$, or the age of the universe $t$.  In addition, we found that the evolution rate $\\zeta$ is proportional to the square root of the reconnection probability. This will effect the string energy density and may lead to distinguishable observational signatures."
    },
    "0410/astro-ph0410059_arXiv.txt": {
        "abstract": "%  We present a deep optical survey of Uranus' Hill sphere for small satellites.  The Subaru 8-m telescope was used to survey about 3.5 square degrees with a $50 \\%$ detection efficiency at limiting red magnitude $m_R$ = 26.1 mag.  This magnitude corresponds to objects that are about 7 km in radius (assuming an albedo of 0.04).  We detected (without prior knowledge of their positions) all previously known outer satellites and discovered two new irregular satellites (S/2001 U2 and S/2003 U3).  The two inner satellites Titania and Oberon were also detected.  One of the newly discovered satellites (S/2003 U3) is the first known irregular prograde of the planet.  The population, size distribution and orbital parameters of Uranus' irregular satellites are remarkably similar to the irregular satellites of gas giant Jupiter.  Both have shallow size distributions (power law indices $q \\sim 2$ for radii $> 7$ km) with no correlation between the sizes of the satellites and their orbital parameters. However, unlike those of Jupiter, Uranus' irregular satellites do not appear to occupy tight distinct dynamical groups in semi-major axis versus inclination phase space.  Two groupings in semi-major axis versus eccentricity phase space appear to be statistically significant. ",
        "introduction": "Planetary satellites are confined to the space in which the planet's gravitational force dominates over the Sun's.  This region is known as the Hill sphere, the radius of which, $r_{H}$, is given by  \\begin{equation} r_H = a_p \\left[\\frac{m_p}{3M_{\\odot}}\\right]^{1/3} \\label{eq:hill} \\end{equation}  \\noindent where $a_p$ and $m_p$ are the semi-major axis and mass of the planet and $M_{\\odot}$ is the mass of the Sun.  Table 1 lists the Hill sphere radii and projected areas for each of the giant planets.  The giant planets possess two distinct types of satellite (Peale 1999).  Regular satellites are found within about $0.05 r_{H}$ and are tightly bound to their planet.  They have nearly circular, prograde orbits with low inclinations.  Regular satellites likely formed within a circumplanetary disk of gas and dust around the giant planets as part of the planetary formation process itself.  In contrast, irregular satellites are found up to $0.65 r_{H}$ from their host planets and have moderate to high eccentricities and inclinations with prograde or retrograde orbits.  Irregular satellites can not have formed in their present orbits and are likely products of early capture from heliocentric orbit (Kuiper 1956; Pollack, Burns \\& Tauber 1979).  Table 1 lists the currently known populations of irregular satellites for each giant planet as of August 1, 2004.  Burns (1986) offered a definition of irregular satellites as those satellites which are far enough from their parent planet that the precession of their orbital plane is primarily controlled by the Sun instead of the planet's oblateness.  In other words, the satellite's inclination is fixed relative to the planet's orbit plane instead of the planet's equator.  By this definition, any satellite with a semi-major axis larger than the critical value, $a_{\\mbox{crit}} \\sim (2 J_{2} r_{p}^{2} a_{p}^{3} m_{p} / M_{\\odot} )^{1/5}$, is an irregular satellite (Burns 1986; see Table 1).  Here $J_{2}$ is the planet's second gravitational harmonic coefficient and $r_{p}$ is the planet's equatorial radius.  Because of the reversibility of Newton's equations of motion some sort of energy dissipation is required for permanent satellite capture. The giant planets currently have no efficient mechanism of energy dissipation for satellite capture.  During the planet formation epoch several mechanisms may have operated to capture satellites: 1) gas drag in an extended, primordial planetary atmosphere (Pollack, Burns \\& Tauber 1979) 2) pull-down capture caused by the mass growth of the planet and consequent expansion of the Hill sphere (Heppenheimer \\& Porco 1977) and 3) orbital energy dissipation from collisions or collisionless interactions between asteroids and/or satellites passing near the planet (Colombo \\& Franklin 1971; Tsui 2000).  Study of the irregular satellites is important for the insight these objects might provide into the planet formation process.  Core accretion models of planet formation struggle to form Uranus and Neptune within the age of the solar system (Lissauer et al. 1995; Pollack et al. 1996; Boss 2001).  Disk instability models do not readily provide the hydrogen- and helium- depleted compositions of the two ice giants (Boss 2001).  The massive hydrogen and helium gas giants Jupiter and Saturn likely formed quickly in the protoplanetary disk ($\\le 10^{6}$ years).  The less massive, deficient in hydrogen and helium, more distant ice giants Uranus and Neptune appear to have taken much longer or an altogether different route of formation (Thommes et al. 2002; Boss 2002).  Uranus is noteworthy in the sense that its obliquity exceeds 90 degrees.  This compares to the modest obliquities of Jupiter, Saturn and Neptune at 3, 27 and 30 degrees respectively.  A possible cause is that a protoplanet of about one Earth mass collided with Uranus near the end of its growth phase (Korycansky et al. 1990; Slattery et al. 1992).  Greenberg (1974) argued that the current regular satellites must have formed after Uranus' obliquity reached 98 degrees because their low inclination prograde orbits would not have adjusted to their current configurations with Uranus.  Recently Brunini et al. (2002) suggested that if Uranus' tilt was created by a giant impact any satellites beyond about $2 \\times 10^{6}$ km (i.e. all known irregular satellites of Uranus) would have likely been lost owing to the orbital impulse imparted to Uranus by the impactor.  In addition, Beauge et al. (2002) show that any significant migration by Uranus through a residual planetary disk would have caused its outer satellites to become unstable.  By virtue of its proximity, Jupiter has the best-studied irregular satellite system (Figure \\ref{fig:distance26uranus}), with 55 irregular satellites currently known (Sheppard \\& Jewitt 2003).  In this paper we ask the question ``does the ice-giant Uranus have a population of irregular satellites similar to that of gas-giant Jupiter?''.  The greater distance of Uranus requires the use of very deep surveys in order to meaningfully probe the smaller Uranian satellites.  Previous surveys near Uranus were conducted using 4-meter class telescopes, including a search of $\\sim$5 deg$^2$ to limiting red magnitude $m_R \\sim$ 24.3 by Gladman et al. (2000) and of $\\sim$1 deg$^2$ to limiting magnitudes in the $m_R \\sim$ 25.0 to $m_R \\sim$ 25.4 range by Kavelaars et al. (2004).  In the present work, we used the 8-meter Subaru telescope and its prime-focus survey camera to survey most of the Hill sphere to a limiting red magnitude $m_R$ = 26.1.  Our primary goal was to cover the dynamically stable inner Hill sphere of the planet (radial extent $\\sim 0.7 r_{H}$; see Hamilton \\& Krivov 1997) in an unbiassed, deep and uniform survey. ",
        "conclusions": "Our survey detected, without prior knowledge of their positions, all of the then-known six irregular satellites of Uranus as well as the two outer-most regular satellites Titania and Oberon.  We also discovered the two new Uranian satellites S/2001 U2 and S/2003 U3 (Sheppard et al. 2003).  S/2001 U2 has a 2001 designation because it had been detected but not confirmed as a Uranian satellite in 2001 (Holman et al. 2003).  S/2001 U3 was confirmed as a Uranian irregular satellite after our survey (Marsden et al. 2003).  This object went undetected in our Subaru data because it fell on a bright star during observations.  In addition, six other objects near Uranus which turned out to be Centaurs, as well as hundreds of Kuiper Belt objects, were discovered.  A detailed report about the discovery of these other objects will be given in a future paper.  To determine the size limit of satellites detectable by the survey and the approximate sizes of the new satellites we relate the apparent red magnitude, $m_{R}$, to the radius, $r$, through  \\begin{equation} r = \\left[ \\frac{2.25\\times 10^{16}R^{2}\\Delta ^{2}}{p_{R}\\phi (\\alpha)} \\right]^{1/2} 10^{0.2(m_{\\odot} - m_{R})} \\label{eq:appmaguranus} \\end{equation}  \\noindent in which $r$ is in km, $R$ is the heliocentric distance in AU, $\\Delta$ is the geocentric distance in AU, $m_{\\odot}$ is the apparent red magnitude of the sun ($-27.1$), $p_{R}$ is the geometric red albedo, $\\phi (\\alpha)$ is the phase function and $\\alpha$ is the phase angle ($\\alpha=0$ deg at opposition).  For linear phase functions we use the notation $\\phi (\\alpha) = 10^{-0.4 \\beta \\alpha}$, where $\\beta$ is the ``linear'' phase coefficient.  Using data from Table~2 and an albedo of 0.04 we find that 26.1 magnitudes corresponds to a satellite with radius of about 7 km.  \\subsection{Size and Population Distribution}  The Cumulative Luminosity Function (CLF) describes the sky-plane number density of objects brighter than a given magnitude.  The CLF is conveniently described by  \\begin{equation} \\mbox{log}[\\Sigma (m_{R})]=\\alpha (m_{R}-m_{o})  \\label{eq:slope} \\end{equation}  \\noindent where $\\Sigma (m_{R})$ is the number of objects brighter than $m_{R}$, $m_{o}$ is the magnitude zero point, and $\\alpha$ describes the slope of the luminosity function.  The CLF for Uranus' irregular satellites is shown in Figure~\\ref{fig:cumuran}.  Our survey is complete to $m_R$ = 26.1 mag ($r > 7$ km).  For Uranus we find the best fit to the CLF for $m_{R} < 26$ mag is $\\alpha = 0.20 \\pm 0.04$ and $m_{o} = 20.73 \\pm 0.2$ (Figure \\ref{fig:cumuran}).  The points in a CLF are heavily correlated with one another, tending to give excess weight to the faint end of the distribution.  The Differential Luminosity Function (DLF) does not suffer from this problem.  We plot the DLF using a bin size of 2 mag for all Uranian irregular satellites in Figure \\ref{fig:diffuran}.  We find that the DLF has a slope of $\\alpha = 0.16 \\pm 0.05$ with $m_{o} = 21.0 \\pm 0.4$.  Bin sizes of 1 mag and 1.5 mag give similar results.  The uncertainty of the fit for the 2 mag bin size is estimated from the fits of these other bin sizes which were larger than the uncertainty of the least square fit for the 2 mag bin size alone.  The luminosity functions represent the combined effects of the albedos, heliocentric distances and size distributions of the objects. If we assume that all Uranian satellites are at the same heliocentric distance and that their albedos are similar, the DLF and CLF simply reflect the size distribution of the satellites.  The objects appear to follow a single power-law size distribution for $r > 7$ km.  In order to model the irregular satellite size distribution we use a differential power-law radius distribution of the form $n(r)dr=\\Gamma r^{-q}dr$, where $\\Gamma$ and $q$ are constants, $r$ is the radius of the satellite, and $n(r)dr$ is the number of satellites with radii in the range $r$ to $r+dr$.  The slope of the DLF ($\\alpha$) and exponent of the size distribution ($q$) are simply related as $q = 5 \\alpha + 1$ when assuming similar heliocentric distance and albedos for all satellites (Irwin et al. 1995).  Using $\\alpha = 0.16 \\pm 0.05$ for Uranus' irregular satellites we find $q = 1.8 \\pm 0.3$.  This is similar to the value found by Kavelaars et al. (2004) and identical to $q = 1.9 \\pm 0.3$ found for the irregular satellites of Jupiter in the same size range (Sheppard \\& Jewitt 2003).  In comparison, collisional equilibrium gives $q \\sim 3.5$ (Dohnanyi 1969), nonfamily small asteroids have $q \\sim 2.0$ to $2.5$ (Cellino et al. 1991), while large Kuiper Belt Objects (KBOs) have $q=4.2\\pm 0.5$, $\\alpha = 0.64 \\pm 0.1$, $m_{o} = 23.23 \\pm 0.15$ (Trujillo et al. 2001) and Centaurs have KBO-like slope with $m_{o} = 24.6 \\pm 0.3$ (Sheppard et al. 2000).  The small Jovian Trojans ($r<30$ km) have $q=3.0 \\pm 0.3$ while the larger Trojans show a steeper slope of $q=5.5 \\pm 0.9$ (Jewitt, Trujillo \\& Luu 2000).  Large members of asteroid families have been found to usually have $q \\ge 4$ (Tanga et al. 1999) while the smaller members ($r < 5$ km) may have shallower distributions (Morbidelli et al. 2003).  If $q>3$, most of the collisional cross-section lies in the smallest objects while for $q>4$ most of the mass lies in small bodies.  The mass and cross-section of Uranus' irregular satellites are both dominated by the few largest objects.  \\subsection{Comparison of the Uranian and Jovian Systems}  The Jovian and Uranus satellite DLFs are compared in Figure \\ref{fig:diffuran}.  Fitting the Jovian DLF (which is complete to around $r \\sim 3$ km; Sheppard \\& Jewitt 2003) over the same size range as that used previously for Uranus ($r > 7$ km) we find a slope of $\\alpha = 0.18 \\pm 0.05$ with a zero point magnitude $m_{o} = 14.0 \\pm 0.4$.  The measured slope is compatible with $\\alpha = 0.16 \\pm 0.05$, as found for the Uranian satellites, within the statistical uncertainties.  We conclude that the size distribution indices of the irregular satellites at Uranus and Jupiter are remarkably similar for $r > 7$ km, and both are quite different from the (steeper) distribution that would be measured, for example, amongst the main-belt asteroids.  Jupiter's irregular satellites appear depleted in the 4 to 10 km size range relative to an extrapolation of a power-law fitted at larger sizes (Sheppard and Jewitt 2003).  We would need deeper survey observations to $m_R \\sim$ 27.2 mag to determine whether the Uranian population shows a similar depletion.  The difference between the Jovian and Uranian DLF magnitude zero points is $\\Delta m_0$ = 7.0$\\pm$0.6 mag.  This is to be compared with the magnitude difference expected from the inverse square law and the different distances to the two planets.  The expected magnitude difference is $\\Delta m$ = 5 $log_{10}\\left[\\frac{R_U(R_U - 1)}{R_J(R_J - 1)}\\right]$, where $R_J$ = 5.3 AU and $R_U$ = 20.0 AU, are the heliocentric distances to the two planets and opposition geometry is assumed.  Substituting, we obtain $\\Delta m$ = 6.1 mag., which is different from $\\Delta m_0$ by only about 1.5$\\sigma$.  In this sense, the smaller number of known irregular satellites at Uranus seems to be an artifact of the greater distance.  To emphasize these points, we show the cumulative size distributions of the Jovian and Uranian irregular satellites in Figure \\ref{fig:cumuransize}. Uranus has nine satellites with r$ \\geq$ 7 km (about the completeness level) while Jupiter has eight.  In terms of the size distributions and total populations, the irregular satellite systems of Uranus and Jupiter are very similar.  If we assume the size distribution of Uranus' irregular satellites extends down to radii of about 1 km, we would expect about $75 \\pm 30$ irregular satellites of this size or larger.  Several competing processes could influence the size distribution of the satellites.  The larger objects may retain some memory of the production function, as is apparently the case with the nearby Jovian Trojans, in which the size distribution is steeper below $r \\sim$ 30 km than above it (Jewitt et al. 2000).  Small satellites could be lost to gas drag, leading to a flattening of the size distribution.  This is more likely at gas-giant Jupiter than at ice-giant Uranus, where much less gas is thought to have been available during the accretion epoch.  Collisions between satellites and with interplanetary projectiles would lead to the production of many small fragments. Given that fragment velocity and size are inversely related, it is natural to expect that the smaller objects produced collisionally would be lost (the escape velocity from Uranus at the semi-major axis of S/2001 U2 is only 0.8 km s$^{-1}$), again leading to a flattened size distribution.  Collisional scenarios, in general, require higher collision rates than now prevail in the solar system.  Perhaps the irregular satellites were originally much more numerous than now. Separately, we know that the flux of planet-crossing projectiles was much higher between the epochs of planet formation and the end of the terminal bombardment phase at about 3.9 Gyr.  \\subsection{Orbital Element Distribution}  Table 3 lists some of the properties of the known irregular satellites of Uranus.  Figures \\ref{fig:irrsatsuranusecc} and \\ref{fig:irrsatsuranus} compare the semi-major axes with inclinations and eccentricities, respectively, for all known irregular satellites of the planets.  The figures show that the ice giants Uranus and Neptune have the smallest known irregular satellite systems, in units of Hill radii.  In the case of Uranus, one contributing factor may be that its regular satellite system is much less massive and does not extend as far from the planet as those of the gas giant planets of Jupiter and Saturn.  Thus, interactions between Uranus' regular and irregular satellites are less important as a clearing mechanism.  We also note that Uranus has the smallest $a_{\\mbox{crit}}$ of any of the giant planets (Table 1).  Neptune's satellite system may have been severely disrupted by the unusually large retrograde satellite Triton (Goldreich et al. 1989).  Figure \\ref{fig:irrsatsuranusecc} suggests that the Uranian irregular satellites may be grouped in semi-major axis vs. eccentricity phase space.  The four retrograde irregular satellites closest to Uranus (S/2001 U3, Caliban, Stephano and Trinculo; semi-major axes of $a < 0.15 r_{H}$) have eccentricities $e \\sim$0.2 while the four retrograde irregular satellites (Sycorax, Prospero, Setebos and S/2001 U2) with $a > 0.15 r_{H}$ have $e \\sim$0.5.  To judge the significance of these two retrograde groups we performed several statistical tests.  The retrograde low eccentricity, low semi-major axis group (the Caliban group) has a mean eccentricity of $0.19 \\pm 0.02$ and mean semi-major axis of $7.0\\pm 0.9 \\times 10^{6}$ km while the retrograde high semi-major axis and eccentricity group (the Sycorax group) has $0.52 \\pm 0.05$ and $16.7 \\pm 1.8 \\times 10^{6}$ km, respectively.  The Student's t-test (with 7 degrees of freedom) gave a t-statistic of 6.0 with a significance of 99.8\\% for the difference in their mean eccentricities and a t-statistic of 4.8 with a significance of 99.4\\% for the difference in their mean semi-major axes.  The student's t-test suggests that the two groups are significant but makes the unjustified assumption that the eccentricity and semi-major axes are normally distributed.  Therefore, we used the more stringent nonparametric Mann-Whitney U-test and a permutation test (see Siegel \\& Castellan 1988) to assess the significance of the two groups.  Both found that the groupings of the eccentricity were statistically significant at or above the $\\geq 3\\sigma$ (99.7\\%) level of confidence.  Figure \\ref{fig:ecccorruran} shows the two retrograde groups in semi-major axis vs. eccentricity space.  It is clearly seen that the more distant satellites have larger eccentricities as was also noted by Kavelaars et al. (2004).  The significance of a linear fit to the data is only at the $\\sim$97 \\% ($< 3 \\sigma$) level and thus less significant than the two groupings.  Possible reasons for higher eccentricities more distant from the planet is that the closer satellites would be unstable to perturbations by the much larger regular satellites of Uranus if they had large eccentricities or that the more distant satellites are susceptible to solar and planetary perturbations.  In Figure \\ref{fig:ecccorruran} we plot several lines of constant periapse from Uranus.  The orbital velocities of the irregular satellites around Uranus range from 300 to 1100 m/s (Kessler 1981).  The relative velocities amongst the satellites are typically much smaller.  Velocity differences among Caliban, Stephano and Trinculo  are comparable to Caliban's escape velocity $\\sim$40 m/s.  The retrograde high eccentricity objects have relative velocities comparable to Sycorax's escape velocity of $\\sim$80 m/s.  For comparison, Jupiter's irregular satellites were found to be grouped in semi-major axis and inclination phase space with their relative velocities within a group about 30 m/s while group velocities relative to each other were over 200 m/s (Sheppard \\& Jewitt 2003).  As at Jupiter, Uranus' possible two retrograde groups have relative velocities of over 100 m/s while members within a group have velocities comparable to the largest members escape velocity.  As at Jupiter, the dynamical groupings suggest formation from parent objects which were collisionally shattered.  A simple particle in a box calculation shows that collisions between the currently known or predicted outer satellite population of Uranus would occur on timescales ($\\sim 10^{10}$ years) longer than the age of the solar system.  Fragmentation could have occurred from collisions with objects in heliocentric orbits (principally comets around the heavy bombardment period (Sheppard \\& Jewitt 2003)) or other now defunct satellites of the planet (Nesvorny et al. 2004).  Jupiter's irregular satellite orbital velocities ($> 2200$ m/s) are much greater than Uranus' and thus any collisional processing would have been much more violent compared to collisions around Uranus.  Except for the distinct prograde irregular S/2003 U3 with an orbital inclination of 57 degrees to the ecliptic compared to the eight irregulars between 140 and 170 degrees there are no obvious tight groupings in semi-major axis vs. inclination phase space as was found around Jupiter (Figure \\ref{fig:irrsatsuranus}).  We do note that the majority of the low eccentricity retrograde satellites have inclinations near 142 degrees while the retrograde high eccentricity satellites are closer to 156 degrees in inclination (Table 3).  The intermediate inclinations $60 < i < 140$ degrees are devoid of known satellites, consistent with instabilities caused by the Kozai instability (Kozai 1962; Carruba et al. 2002; Nesvorny et al. 2003). In this instability, solar perturbations at apoapse cause the satellites at high inclinations to acquire large eccentricities which eventually lead to collisions with the planet, or a regular satellite or loss from the Hill sphere in $10^{7} - 10^{9}$ years (Carruba et al. 2002; Nesvorny et al. 2003).  We find no clear size vs. semi-major axis, inclination or eccentricity correlations for Uranus' irregular satellites, as may be expected if significant gas drag was present in the past.  Of the two largest irregular satellites around Uranus one of each is in the two possible eccentricity groups.  Caliban is relatively close to the planet with a low eccentricity while Sycorax is with the distant, higher eccentricity irregular satellites (Table 3).  At Uranus there are many more known retrograde (8) than prograde (1) outer satellites.  Jupiter also has an over abundance of known retrograde outer satellites (48 retrograde versus 7 prograde).  These asymmetries are greatly diminished if we compare numbers of satellite groups ($\\sim 3-4$ retrograde versus 3 prograde groups at Jupiter and possibly 1 or 2 retrograde and 1 prograde group at Uranus).  Given this, and the fact that the statistics of the groups remain poor (especially at Uranus), we cannot currently use the relative numbers of retrograde and prograde objects to constrain the mode of capture."
    },
    "0410/astro-ph0410090_arXiv.txt": {
        "abstract": "High velocity outflows from supermassive black holes have been invoked to explain the recent identification of strong absorption features in the hard X-ray spectra of several quasars. Here, Monte Carlo radiative transfer calculations are performed to synthesise X-ray spectra from models of such flows. It is found that simple, parametric bi-conical outflow models with plausible choices for the wind parameters predict spectra that are in good qualitative agreement with observations in the 2 -- 10~keV band. The influence on the spectrum of both the mass-loss rate and opening angle of the flow are considered: the latter is important since photon leakage plays a significant role in establishing an ionization gradient within the flow, a useful discriminant between spherical and conical outflow for this and other applications. Particular attention is given to the bright quasar PG1211+143 for which constraints on the outflow geometry and mass-loss rate are discussed subject to the limitations of the currently available observational data. ",
        "introduction": "There is a growing body of evidence for high-velocity, blueshifted absorption features in the hard (2 -- 10~keV) X-ray spectra of quasars. Chartas et al. (2002) first identified a pair of absorption features in the X-ray spectrum of the relatively high redshift quasar APM08279+5255 ($z=3.91$) and proposed that they were due to highly ionized iron. They were unable, however, to identify the ionization stage with certainty. More recently, Pounds et al. (2003a) have identified narrow absorption features in the spectrum of the nearby quasar PG1211+143. In this case, the identification of an X-ray line at $\\sim 7.5$~keV with the significantly blueshifted ($v \\sim 0.1c$) K~$\\alpha$ transition of Fe~{\\sc xxv} or {\\sc xxvi} was supported by the identification of other, weaker features at softer energies. Similar narrow absorption features have also now been reported in a second high-redshift quasar (PG1115+080, Chartas, Brandt \\& Gallagher 2003) and a second low-redshift quasar (PG0844+349, Pounds et al. 2003b). In addition, high velocity absorption has been reported in yet another quasar (PDS456) by Reeves, O'Brien \\& Ward (2003), although in this case the absorption is broader and the authors favour an interpretation involving absorption edges of Fe~{\\sc xvii}-{\\sc xxiv} rather than lines of more highly ionized material.  It has been suggested that these absorption features form in a fast, massive outflow from the accreting supermassive black hole in the quasar nucleus (Chartas et al. 2002, Pounds et al. 2003a, King \\& Pounds 2003). To account for the observed absorption features requires such flows to be both more highly ionized and have significantly greater column densities than necessary for the outflows detected via blueshifted ultraviolet absorption lines in broad absorption line (BAL) quasars (Pounds et al. 2003b). It is not yet established whether highly ionized fast outflows might be common in the quasar population, but it is clear that when present, they may carry sufficient energy to be important in the energy budget of accretion by the nuclear black hole. In addition, King \\& Pounds (2003) have presented simple arguments to show that outflows from black holes accreting at around the Eddington limit are likely to be optically thick giving rise to an XUV photosphere, emission from which may be responsible for the ``big blue bump'' (BBB) seen in quasar spectra.  This interpretation, however, is not firmly established: in some instances (e.g. APM08279+5255; Chartas et al. 2002) there remains ambiguity in the line identification while in others (PG1211+143, PDS456; McKernan, Yaqoob \\& Reynolds 2004) it has been suggested that some portion of the absorption features is due to gas in the vicinity of our Galaxy, rather than material intrinsic to the quasar. In addition, Kaspi (2004) has shown that many of the features in the spectrum of PG1211+143 can be explained by an alternative model involving a relatively small outflow velocity ($\\sim 3000$~km~s$^{-1}$). It is noted, however, that the preliminary study presented by Kaspi (2004) addresses only the {\\it XMM-Newton} Reflection Grating Spectrometer (RGS) data and not the European Imaging Camera (EPIC) data: it is not clear that a low-velocity outflow could explain the absorption features in the EPIC spectrum identified by Pounds et al. (2003a).  To better understand the hard X-ray absorption features and the potentially important outflows with which they may be associated requires synthesis of the X-ray spectrum predicted by plausible physical models. The primary goal of this paper is to perform realistic radiative transfer calculations and synthesise spectra in order to determine whether simple physical outflow models, as discussed by Pounds et al. (2003a) and King \\& Pounds (2003), readily predict strong, narrow absorption features as required by the observations.  A secondary objective is to use model spectra to constrain the plausible range of flow parameters by comparison with the data for a particular quasar. This investigation is focused on the narrow emission line quasar PG1211+143. Of the two nearby quasars in which narrow hard X-ray absorption lines have been detected by Pounds et al. (2003a,b), this object has been chosen over the other candidate (PG0844+349) since the observational data is of higher quality and a wider variety of absorption features have been identified. PG1211+143 is bright, nearby ($z = 0.0809$, Marziani et al. 1996) and known to be radiating at around the Eddington luminosity (Boroson 2002, Gierli\\'{n}ski \\& Done 2004). Based on simple ionized absorber fits to the data, Pounds et al. (2003a) suggest the high-velocity outflow in PG1211+143 has a mass-loss rate $\\sim 3$~M$_{\\odot}$~yr$^{-1}$, and they suggest that the outflow is likely to subtend a wide opening angle as viewed from the central black hole.  This work is also of general interest for the study of radiative transfer in non-spherical outflows. It will be shown that calculations for conical geometries are significantly different from spherical outflow, primarily owing to the influence of photon leakage through the conical boundary.  During the refereeing of this article, Everett \\& Ballantyne (2004) presented a related study on radiatively driven outflows from black holes. They investigated continuum driven flows and conclude that, while such flows are possible, they are likely to be too highly ionized to produce absorption line features. They did not, however, investigate line or MHD driven outflows nor consider multi-dimensional radiative transfer effects.  In Section 2, the outflow model that will be considered is discussed. The calculations presented in this paper have been performed with a Monte Carlo radiative transfer code which has been adapted from the code written by Sim (2004). The code uses the Macro Atom formalism developed by Lucy (2002, 2003) and is discussed in Section 3. Section 4 identifies the atomic data used in the Monte Carlo simulations. The spectra computed from various models are discussed in Section 5 and conclusions drawn in Section 6. ",
        "conclusions": "It has been shown that realistic radiative transfer calculations support the proposal that the recently identified X-ray absorption features in PG1211+143 (in particular the strong absorption line observed at $\\sim 7$~keV) can be explained in terms of a simply parameterised bi-conical outflow model.  Spectra have been computed for outflows with a range of both mass-loss rate and opening angle. The opening angle determines the extent of photon leakage from the flow which plays an important role in establishing a gradient of ionization in the wind. This affects the strengths of the spectral features providing a useful discriminant between spherical and conical outflow. The calculations presented here favour models in which the flow has a moderately small opening angle $\\leq 30$\\textdegree~($b \\leq 0.13$) and mass-loss rate a little larger than that suggested by Pounds et al. (2003a)~($\\Phi/b \\sim $ 6~M$_{\\odot}$ yr$^{-1}$ is proposed here).  A narrow opening angle suggests interpretation in terms of a collimated disk wind or jet originating from the accretion disk. Such outflows are certainly favoured as elementary components of the structure of quasars and a relatively narrow opening angle bi-conical outflow is consistent with existing models for quasars (e.g. Antonucci \\& Miller 1985; Elvis 2000, 2004). In particular, in the context of the Elvis (2000) model, one may speculate that the highly ionized outflow modelled here lies in some portion of the polar region bounded by the WHIM. Alternatively the X-ray features may form in the most extremely ionized parts of the WHIM itself.  Given that high velocity outflows may be important for understanding the processes which operate in the cores of quasars, there is a real need for higher quality observational data (such as may be provided by the next generation of X-ray observatories as discussed by Chartas et al. 2003). In particular, it is to be hoped that an increase in spectral resolution such as will be provided by the forthcoming {\\it Astro-E2} mission, will allow the study of line profiles in much greater detail than is currently feasible.  As better observational data becomes available, future theoretical work will include a detailed investigation of the various absorption and emission features observed at softer energies. Although the O~{\\sc viii} Lyman~$\\alpha$ line has been discussed here, several of the lower ionization features that are observed (e.g. those of O~{\\sc vii}) are not predicted by the models at present. An important step in modelling these features is likely to be self-consistent calculation of the temperature in the models since variation of the temperature through the flow may lead to quantitative changes in the ionization balance owing to its effect on recombination rates. Temperature calculations will also provide insight into the formation of a photosphere in the outflow, which might explain the BBB (Pounds et al. 2003a, King \\& Pounds), and help constrain the launching radius of the wind ($R_c$). Such calculations go beyond the scope of the simple modelling presented here but the influence of the choice of temperature on the hard X-ray spectrum is briefly discussed in the Appendix. It may also become necessary to consider angular ($\\theta$) variation of the ionization fractions in the flow, departures from ionization equilibrium, rotation and the influence of viewing angle if the basic bi-conical geometry stands up to tighter observational constraints.  This paper has not addressed the important issue of how high velocity flows may be accelerated. It has been suggested that radiation pressure may be primarily responsible (King \\& Pounds 2003, Pounds et al. 2003a). Since the source here has been assumed to radiate at the Eddington luminosity it is, by definition, the case that radiation pressure is sufficient to overcome gravity but details of how the flow is accelerated over and above gravity to the terminal velocity warrant investigation. Everett \\& Ballantyne (2004) have concluded that continuum driving alone cannot account for outflows such as have been proposed by Pounds et al. (2003a). However, the part played by spectral lines in driving an outflow has yet to be quantified. If the flow is radiatively driven by lines and bound-free continua then understanding the driving is closely coupled to modelling the softer parts of the spectrum than have been considered here: the X-ray region is relatively sparse of spectral lines and very preliminary indications from the models presented here suggest that the lines in this region are unlikely to provide a significant fraction of the required outward force.  In conclusion, the spectral synthesis presented in this paper suggests that simple flow models can reproduce the absorption features observed in the hard X-ray spectrum of PG1211+143. However, significant further study, both observational and theoretical, is needed in order that these high velocity outflows and their relationship to the black hole accretion process which lies at the heart of the quasar phenomenon can be understood."
    },
    "0410/astro-ph0410573_arXiv.txt": {
        "abstract": "\\noindent The collapse of a rotating 3-axis ellipsoid is approximated by a system of ordinary differential equations. Violent relaxation, mass and angular momentum losses are taken into account phenomenologically. The formation of the equilibrium configuration and different types of instability are investigated.\\\\  \\noindent {\\bf Keywords:} dark matter - large-scale structure - instability ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410329_arXiv.txt": {
        "abstract": "The MIGALE project provides databases and data analysis tools to study the evolution of galaxies from z=1 to z=0. It develops and maintain a general database, HyperLeda, to give a homogenized parameterization for 3 million objects, and several archives or specialized databases. It also develops tools to analyse on-the-fly data extracted from the database or obtained through the Virtual Observatory (Virtual Instruments). The package made for this project, Pleinpot, is distributed as open source. ",
        "introduction": "The goal of the Virtual Observatory (VO, see http://www.ivoa.net/) is to enhance the scientific efficiency by providing a transparent access to the archive data stored throughout the world, so that the end user will have the filling that all these resources are located on its own computer. To make this enormous quantity of data useful (volumes are counted in Pbytes), high performance discovery tools are needed, and a condition is to provide an accurate description of any individual dataset.  A precise description also enables to change the strategy for the data analysis. For example, it will not be required anymore to tell to the analysis program what is the spectral resolution and how it changes with wavelength: This information will come with the data, as part of the so-called {\\it meta-data} or more familiarly FITS keywords. It will be possible soon to push one step forward the automation of the data analysis.  Why is it desirable to automate analysis pipelines? Naturally, the main reason is that it is expected to be faster, and possibly more reliable since it eliminates some risks of errors during the preparation of the reduction procedure. The complexity of data analysis also increases as it allows to put finer constraints to the details of the physical mechanisms. So, astronomers interested in a specific phenomenon and using multi-wavelength data can hardly be experts in all the aspects of the data analysis; more automatic procedures are simpler to use and enable non-specialists to interpret the data. There is obviously the risk to miss-use a procedure, as for example give a wrong combination of parameters to a model, but astronomers are highly aware of this danger which anyway already exists when we are using the present software. The VO does certainly not worsen the situation.  Actually, not only it is desirable to automate processing, but it is necessary. Some spectrographs produce millions of spectra over their lifetime, and it is absolutely impossible to make a detailed interactive analysis of each of them (see for example the analysis of SDSS data in Mathis et al. 2004).  Analysis pipelines, to compare models to observations, will be developed either as client packages to be run on the user's machine or as online services run on a remote server fed by data provided by the user or extracted from the VO and returning results in the formats defined by the VO so they can be used by other VO tools. Such services already exist. For example, a VO implementation of sextractor was used at the 2004 AVO science demo (Padovani et al. 2004). ",
        "conclusions": ""
    },
    "0410/astro-ph0410603_arXiv.txt": {
        "abstract": "We use a sample of 69726 galaxies from the SDSS to study the variation of the bimodal color-magnitude (CM) distribution with environment.  Dividing the galaxy population by environment ($\\Sigma_5$) and luminosity ($-23<M_r<-17$), the $u-r$ color functions are modeled using double-Gaussian functions.  This enables a deconvolution of the CM distributions into two populations: red and blue sequences.  The changes with increasing environmental density can be separated into two effects: a large increase in the fraction of galaxies in the red distribution, and a small color shift in the CM relations of each distribution. The average color shifts are $0.05\\pm0.01$ and $0.11\\pm0.02$ for the red and blue distributions, respectively, over a factor of 100 in projected neighbor density.  The red fraction varies between about 0\\% and 70\\% for low-luminosity galaxies and between about 50\\% and 90\\% for high-luminosity galaxies.  This difference is also shown by the variation of the luminosity functions with environment. We demonstrate that the effects of environment and luminosity can be unified.  A combined quantity, $\\Sigma_{\\rm mod} = (\\Sigma_5/{\\rm Mpc}^{-2}) \\, + \\, (L_r / L_{-20.2}) \\,$, predicts the fraction of red galaxies, which may be related to the probability of transformation events. Our results are consistent with major interactions (mergers and/or harassment) causing galaxies to transform from the blue to the red distribution. We discuss this and other implications for galaxy evolution from earlier results and model the effect of slow transformations on the color functions. \\footnote{Article written 2004 August 16th.  Refs.\\ updated 2004 October 25th.  To appear in AIP Proc., The New Cosmology, eds.\\ R.~E.\\ Allen et al.} ",
        "introduction": "\\label{sec-baldry:intro}  The study of galaxy properties in cosmology is important both for understanding the formation and evolution of galaxies and for interpreting measurements of large-scale structure.  For example, it is necessary to determine the clustering bias of different types of galaxies relative to dark matter in order to accurately quantify cosmological parameters.  Galaxies were first classified based on their single-color morphological properties by Hubble in the 1920's \\citep{Hubble26}. The classification followed a sequence of increasing complexity from ellipticals (E0-E7) to lenticulars (S0) to spirals (Sa-Sc) to irregulars (Irr).\\footnote{The adjectives ``early'' and ``late'' were used to describe the relative positions in the morphological sequence \\citep{Hubble26}.  The direction from simple to complex forms was chosen from the precedent of stellar spectral types, where early types (O+B stars) generally have more simple spectral features than late types (K+M stars).  Ironically, the spectra of early-type galaxies are generally dominated by late-type stellar spectra and vice versa.} With the advent of color measurements, morphology-color relations were established, with early types being generally redder in optical colors than late types (for a review, see ref.~\\citep{RH94}).  While there are many relationships between properties for various types of galaxies, color and absolute magnitude are two of the most useful variables; and unlike structural and spectral properties they are less dependent on imaging resolution and aperture effects, respectively.  A color-magnitude (CM) relationship for E+S0 galaxies was shown to have a shallow slope with a small intrinsic scatter \\citep{Faber73,VS77,APF81}.  This was related to a metallicity-luminosity correlation, with more luminous galaxies having a higher luminosity-weighted metallicity \\citep{Faber73,Larson74,KA97}.  Spirals also follow CM relationships but with larger intrinsic scatter \\citep{CR64,VG77,Griersmith80,Visvanathan81,TMA82}.  For spirals, the luminosity correlations can be attributed to changes in star formation history (SFH) \\citep{PdG98}, dust attenuation \\citep{tully98}, and/or metallicity \\citep{ZKH94}.  When all types are considered together, the color function of galaxies can be approximated by the sum of two Gaussian functions that is a bimodal function \\citep{strateva01,baldry04}. This argues that the natural division in the galaxy population is into two distributions; at least out to $z\\sim1$ \\citep{bell04red}.  Related spectral quantities such as from H$\\alpha$ emission \\citep{balogh04}, H$\\delta$ absorption and the 4000\\AA\\ break \\citep{kauffmann03B}, and the derived star formation rate (SFR) \\citep{brinchmann04}, also produce a bimodal distribution. In \\S~\\ref{sec-baldry:cm-distribution}, we review a quantitative study of the CM distribution of galaxies, over all environments, using a two population model; in \\S~\\ref{sec-baldry:enviro}, we analyze the environmental dependence of the color bimodality; and in \\S~\\ref{sec-baldry:discussion}, we discuss the results. ",
        "conclusions": "\\label{sec-baldry:discussion}  Many of the discussions in the literature have focused on quite specific mechanisms for different types of galaxies (e.g., see review by \\citet{Fritze04} on S0 galaxies). Here, the analysis shows that there are essentially two phase spaces for galaxies (in terms of color and absolute magnitude). Even with the addition of other parameters such as surface brightness and concentration, the multi-variate distribution is bimodal \\citep{hogg02red,blanton03broadband}; while bimodality is not a ubiquitous feature of semi-analytic models of galaxy formation \\citep{WF91,kauffmann99,SP99,cole00}. Perhaps the evolution of galaxies is in some sense simple.  Regardless of the processes (accretion, merging, harassment, ram-pressure stripping), galaxies will populate one of two regions in parameter space and while the transformation from the blue to the red distribution may be deterministic (e.g.\\ major merger), the properties within a distribution depend most strongly on the mass/luminosity of the galaxy and the effects of environment are mostly chaotic (e.g.\\ the effects of minor mergers depending on impact parameters).  What about morphology? It could be argued that the CM distributions defined here are reproducing the morphology-density and CM relations of E+S0 and spiral galaxies. It is not possible to obtain Hubble types for all galaxies in the SDSS main galaxy sample because the imaging resolution is not sufficient. Nevertheless, it is unlikely that the red/blue distributions correspond precisely to these classes. Observers do not always agree on Hubble type even with many resolution elements and, in the spirit of our interpretation, whatever processes give rise to the blue/red distribution should also give rise to {\\em distributions} in morphology. Thus, S0 or Sa galaxies could have a probability of belonging to one or the other distributions and should not be considered as classes. The method presented in this article could be extended to include quantified morphological parameters to determine the minor/major environmental effects in terms of morphology.  What about other populations? This analysis only shows that there are two dominant populations. There are other distinct populations or distinct sub-categories. It is important to distinguish between extremes of a population and populations with separate identities. For example, passive spirals \\citep{couch98,goto03} could be regarded as the extremes of one or both of the populations, whereas post-starburst galaxies \\citep{zabludoff96,goto03hd} could represent a separate, transforming, population. In the next section, we discuss a model that includes a significant fraction of slowly transforming galaxies.  \\subsection{Transformation modeling}  One possibility for the trend in population abundance with environment could be that the mass function varies with local density.  However, observations of the near-infrared luminosity function show that any change with environment is likely to be small \\citep{depropris98,AP00,balogh01}.  On the other hand, the strong redshift evolution observed in the colors of galaxies in both clusters \\citep{BO84,fairley02,FZM04} and the field \\citep{lilly96,connolly97} suggests that some galaxies transform from one population to another. In particular, \\citet{bell04red} noted a build up of stellar mass on the red sequence by a factor of about two between $z\\sim1$ and 0, averaged over all environments.  Various mechanisms for these transformations have been proposed.  Some, like galaxy mergers and ram-pressure stripping by the intracluster medium, occur on timescales that are short compared with the lifetime of galaxies \\citep{GG72,FN99,moore99}; other processes, such as the gradual starvation of a system through the removal of gas, result in a slow decline in star formation \\citep{LTC80,BNM00}.  In Fig.~\\ref{fig-baldry:modelcomp} we show the $u-r$ color evolution of two model galaxies, generated using the {\\sc GALEV} population-synthesis code \\citep{BC03}.  We start with a 7\\,Gyr old galaxy that has been forming stars at a rate that has been slowly declining exponentially, with a timescale $\\tau=4$\\,Gyr.  Assuming a Salpeter IMF and no dust extinction, this galaxy has a color $u-r\\sim1.4$, which is close to the peak of the observed blue distribution for the fainter galaxies in our sample.\\footnote{The dispersion around the color peak could be for a number of reasons: (i) stochastic variations in star formation (while the general trend for a population of galaxies could be similar, the measured color could vary because of recent bursts or quiescent periods); (ii) variations in dust attenuation because of disk orientation or intrinsic levels of dust; (iii) variations in metallicity; and (iv) photometric errors.} In the first model, we assume all star formation activity ceases after 7\\,Gyr (designated $\\Delta t=0$).  In this case, the galaxy rapidly becomes very red, reaching well within the observed red distribution ($u-r\\sim2$) in less than 0.2\\,Gyr.  Therefore, if galaxies have been transforming at a uniform rate for the last 13.7\\,Gyr, we would only expect to see $<1$\\% of them with intermediate colors $1.4<u-r<2$ (in addition to those from the two normal distributions).  This would not have a noticeable effect on the simplicity of the observed bimodal population. Thus, it is possible that most or all of the red galaxies have been formed through short-timescale transformation from the star-forming blue population. This is consistent with the observed existence of galaxies with short-lived spectral features indicative of recent changes in star formation history \\citep{DG92,goto03hd,poggianti04,quintero04}.  In this interpretation, the trend for galaxies to become slightly redder with increasing density is due to mechanisms (e.g.\\ metallicity, dust or previous SFH) that are independent of this transformation process.  \\begin{figure} \\includegraphics[width=0.7\\textwidth]{baldryF7.ps} \\caption{The $u-r$ color evolution of two models in which the SFR is reduced at time $\\Delta t=0$.  Both models are initially evolved for 7\\,Gyr ($\\Delta t<0$) using a Salpeter IMF and an exponentially declining SFR with a timescale $\\tau=4$\\,Gyr.  The {\\it solid line} represents a model in which star formation completely ceases at $\\Delta t=0$, while the {\\it dashed line} represents a model in which the SFR decreases exponentially with $\\tau=2$\\,Gyr.} \\label{fig-baldry:modelcomp} \\end{figure}  An alternative interpretation is that the transformation is more gradual, as expected in some models \\citep{LTC80,BNM00} and suggested by the lack of environmental dependence of correlations between different SFH indicators \\citep{kauffmann04}.  The second model in Fig.~\\ref{fig-baldry:modelcomp} shows the color evolution of a galaxy in which, for $\\Delta t>0$, the SFR declines exponentially with a timescale $\\tau=2$\\,Gyr.  This decline is still faster than that prior to $\\Delta t=0$, but is long enough that the system is observed with intermediate colors $1.4<u-r<2$ for a substantial amount of time ($\\sim 3$\\,Gyr).  This could produce a significant distortion on the bimodal distribution, and might be the cause of the apparent redward shift of the mean of the blue galaxy distribution with increasing density.  To demonstrate this, we show in Fig.~\\ref{fig-baldry:transform} a series of model fits to the galaxies in moderately dense environments ($0.9<\\Sigma_5<4.4$), where transformations from the star forming population might be expected to be most common.  We restrict the fits to the population fainter than $M_r=-21$, where the blue and red populations are most distinct.  The solid lines show the default, double-Gaussian model fits presented in Fig.~\\ref{fig-baldry:color-enviro}.  The $\\chi^2$ value for this fit, and the fraction of galaxies in the red distribution, are shown in the legend of each panel.  We now assume that the mean color of the blue population is independent of density, and equal to the mean that we compute in the lowest density bin ($\\Sigma_5<0.13$).  This model is shown as the dotted line, and is a poor fit to the data with a much larger $\\chi^2$ value in all luminosity bins.  To model the effect of including a transforming population, we introduce another parameter, which is the proportion of galaxies in a third population existing strictly between the means of the blue and red peaks.  The color distribution of this population over this range is determined by the amount of time spent at each color, as given by the slope of the dotted line in Fig.~\\ref{fig-baldry:modelcomp}.  For the simple model considered here, this transforming distribution is nearly uniform in $u-r$ color.  We then fit the amplitude of this population, as well as those of the red and blue Gaussian distributions, to minimize the $\\chi^2$ of the model. This fit is shown as the long-dashed line in Fig.~\\ref{fig-baldry:transform}.  Although the $\\chi^2$ value is larger than the default two-component model, it still provides an acceptably good fit.  The fraction of galaxies required to be in the transforming population is about 15\\%--20\\%; and when we account for the duty cycle of this population, up to 50\\% of the galaxies may have undergone such a transformation over the past 13.7\\,Gyr, assuming the rate has remained constant over that time.  \\begin{figure} \\includegraphics[width=1.0\\textwidth]{baldryF8.ps} \\caption{Modeling the bimodality. Observed color distributions ({\\it points with error bars}) are shown for four luminosity bins, restricted to the second-highest density bin ($0.9<\\Sigma_5<4.4$).  The {\\it solid line} shows the double-Gaussian model (Fig.~\\ref{fig-baldry:color-enviro}), where the mean and amplitude of each Gaussian are fit as free parameters.  The {\\it dotted line} shows the best fit obtained with the mean of the blue distribution constrained to be the same as in the fit to the lowest density bin ($\\Sigma_5<0.13$).  The {\\it dashed line} shows the best fit obtained by adding a third component, consisting of a population of galaxies with colors intermediate between the two peaks, with a distribution given by the rate of color evolution in the $\\tau=2$\\,Gyr model (Fig.~\\ref{fig-baldry:modelcomp}).  The $\\chi^2$ values and the fraction of galaxies in the red distribution ($f_r$) are shown in the legend for all of these models.  For the third model, the fraction of the total galaxy population comprised by these transforming galaxies at the present day ($f_{\\rm trans}$) and the fraction of galaxies that would have passed through this phase in the last 13.7\\,Gyr ($f_{\\rm hubble}$), assuming the rate of transformation has been uniform, are also shown.} \\label{fig-baldry:transform} \\end{figure}  We conclude that the best fit to the color distribution in each environment is provided by the two-population model presented in refs.~\\citep{baldry04,balogh04bimodal}, with the possibility that the entire population of red galaxies has been built out of transformations from the bluer population that are ongoing today. These data even accommodate a relatively slow timescale for this transformation, if it is assumed that the bulk of the blue galaxy population has a mean color that is independent of environment. However, a dominance of slow transformations is likely ruled out by observations of bimodality at $z>1.5$ \\citep{somerville04} and by morphology-color relations, which imply a more violent origin because of bulge formation.  \\clearpage \\begin{theacknowledgments} The results presented here made use of the CMU-PITT SDSS Value Added Catalog\\footnote{\\url{http://astrophysics.phys.cmu.edu/dr2_value_added/}} created and maintained by K.~Simon Krughoff and Christopher J.~Miller. I.\\,K.\\,B.\\ and K.\\,G.\\ acknowledge generous funding from the David and Lucille Packard Foundation.  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. \\end{theacknowledgments}"
    },
    "0410/astro-ph0410435_arXiv.txt": {
        "abstract": "We present metallicities for 487 red giants in the Carina dwarf spheroidal (dSph) galaxy that were obtained from FLAMES low-resolution Ca triplet (CaT) spectroscopy. We find a mean [Fe/H] of $-1.91$\\,dex with an intrinsic dispersion of 0.25\\,dex, whereas the full spread in metallicities is at least one dex. The analysis of the radial distribution of metallicities reveals that an excess of metal poor stars resides in a region of larger axis distances. These results can constrain evolutionary models and are discussed in the context of chemical evolution in the Carina dSph. ",
        "introduction": "Analyses of the faint Carina dSph have revealed that it contains a % variety of stellar populations (e.g.,~[4]), exhibiting prominent old ($>$11\\,Gyr) and intermediate-age (5--6 and 3\\,Gyr) populations. This implies that Carina must have undergone several star forming (SF) episodes with at least three significant pulses. Despite this wide spread in ages, its colour-magnitude diagram features a remarkably narrow RGB. The reason for that can be an age counteracting spread in metallicities, where metal rich, young stars have colours comparable to the older, more metal poor ones. Such a possible age-metallicity degeneracy can be overcome if accurate and independent [Fe/H] measurements are obtained so that the remaining parameter of age can be estimated from isochrones. Moreover, the overall shape and spatial variations of the metallicity distribution function (MDF) itself contain valuable implications for analysing Carina's unusual SF history. ",
        "conclusions": ""
    },
    "0410/astro-ph0410144_arXiv.txt": {
        "abstract": "We have started a systematic study of the field topologies of magnetic single and accreting white dwarfs using Zeeman tomography. Here we report on our analysis of phase-resolved flux and circular polarization spectra of the magnetic cataclysmic variables BL\\,Hyi and MR\\,Ser obtained with FORS1 at the ESO VLT. For both systems we find that the field topologies are more complex than a dipole or an offset dipole and require at least multipole expansions up to order $l = 3$ to adequately describe the observed Zeeman features and their variations with rotational phase. Overall our model fits are in excellent agreement with observations. Remaining residuals indicate that the field topologies might even be more complex. It is, however, assuring that the global characteristics of our solutions are consistent with the average effective field strengths and the halo field strengths derived from intensity spectra in the past. ",
        "introduction": "Magnetism at a detectable level is a common phenomenon among white dwarfs. According to recent studies the incidence of objects with surface field strengths exceeding 2\\,MG is at least $\\sim$10\\% of all single white dwarfs, and could even be higher (Liebert et al. 2003). A similar fraction of $\\sim$20\\% magnetic systems has been found among accreting white dwarfs in cataclysmic variables (G\\\"ansicke, this volume). Observed field strengths range from $\\sim$\\,1\\,kG--1000\\,MG with a peak around 16\\,MG in single white dwarfs (Wickramasinghe \\& Ferrario 2000, G\\\"ansicke et al. 2002, Schmidt et al. 2003) and 7--230 MG in magnetic CVs (Beuermann 1998). The frequency distributions of field strengths for both populations are compared in Fig.\\,1.  The origin of the magnetic fields is not well understood. While it is reasonable to assume that the white dwarfs with the highest magnetic fields evolve from main-sequence Ap and Bp stars, low- and moderate-field magnetic white dwarfs appear to imply another origin. The decay times of the lowest multipole components are predicted to be long compared to the evolutionary ages of the white dwarfs. The magnetic field topologies, at least of isolated white dwarfs, are, therefore, likely to be relics of previous evolutionary phases. In accreting systems, the field structure in the outer layers of the white dwarf, however, may have been significantly changed if the accretion rate is high enough that accretion occurs more rapidly than ohmic diffusion (Cumming 2002) and Rayleigh-Taylor instabilities (Romani 1990).  \\begin{figure}[!ht] \\plotone{reinsch_fig01.eps} \\caption{Frequency distributions of the field strengths of accreting {\\it (top)} and single {\\it (bottom)} white dwarfs. For the former, average field strengths in the main accretion region have been used of all 53 polars for which reliable measurements are available (cf. Beuermann 1998). The statistics for single white dwarfs include data from the compilation by Wickramasinghe \\& Ferrario (2000) and  white dwarfs from the SDSS (Schmidt et al. 2003). The distribution for magnetic CVs is biased by the omission of low-field systems (intermediate polars) for which only field strength estimates exist and probably also by selection effects which hamper the discovery of high-field polars. } \\end{figure} ",
        "conclusions": "The Zeeman tomographical analysis of phase-resolved spectropolarimetry provides for the first time detailed information about the range of field strengths and the field topology of accreting white dwarfs. For both systems discussed here we find that at least multipole expansions up to order $l = 3$ are required to describe the field topologies. Remaining residuals indicate that the field topologies might even be more complex."
    },
    "0410/astro-ph0410658_arXiv.txt": {
        "abstract": "We construct analytically stationary global configurations for both aligned and logarithmic spiral coplanar magnetohydrodynamic (MHD) perturbations in an axisymmetric background MHD disc with a power-law surface mass density $\\Sigma_0\\propto r^{-\\alpha}$, a coplanar azimuthal magnetic field $B_0\\propto r^{-\\gamma}$, a consistent self-gravity and a power-law rotation curve $v_0\\propto r^{-\\beta}$ where $v_0$ is the linear azimuthal gas rotation speed. The barotropic equation of state $\\Pi\\propto\\Sigma^{n}$ is adopted for both MHD background equilibrium and coplanar MHD perturbations where $\\Pi$ is the vertically integrated pressure and $n$ is the barotropic index. For a scale-free background MHD equilibrium, a relation exists among $\\alpha$, $\\beta$, $\\gamma$ and $n$ such that only one parameter (e.g., $\\beta$) is independent. For a linear axisymmetric stability analysis, we provide global criteria in various parameter regimes. For nonaxisymmetric aligned and logarithmic spiral cases, two branches of perturbation modes (i.e., fast and slow MHD density waves) can be derived once $\\beta$ is specified. To complement the magnetized singular isothermal disc (MSID) analysis of Lou, we extend the analysis to a wider range of $-1/4<\\beta<1/2$. As an example of illustration, we discuss specifically the $\\beta=1/4$ case when the background magnetic field is force-free. Angular momentum conservation for coplanar MHD perturbations and other relevant aspects of our approach are discussed. ",
        "introduction": "Magnetic fields are ubiquitous in diverse astrophysical settings on various scales ranging from proto-stellar discs, young stellar objects (YSOs), microquasars, quasars, galaxies and grandiose astrophysical jets to clusters of galaxies. In some cases, magnetic fields are relatively weak so that dynamical processes are almost unaffected by their presence. However, there do exist numerous cases where magnetic fields are necessary and important for both dynamics and diagnostics, especially in spiral galaxies and accretion disc systems (e.g., Sofue et al. 1986; Beck et al. 1996; Balbus \\& Hawley 1998; Tagger \\& Pellat 1999; Widrow 2002). In general, it is challenging to model magnetic fields realistically because of their complexities (Sofue et al. 1986; Kronberg 1994; Caunt \\& Tagger 2001; Widrow 2002). In the problems involving protostar formation and disc galaxies, Shu \\& Li (1997) presumed the so-called `isopedic' magnetic field configuration where the mass-to-magnetic flux ratio in a razor thin disc remains constant. Using a singular isothermal disc (SID) model (Mestel 1963) that is isopedically magnetized, Shu et al. (2000) studied stationary coplanar perturbations and proposed that these global modes are bifurcations to either secularly or dynamically unstable configurations. Along a separate yet complementary line, Lou (2002) studied an azimuthally magnetized singular isothermal disc (MSID) and derived two different global stationary magnetohydrodynamic (MHD) perturbation modes. Based on the MSID model, Lou (2002) explored the manifestation of interlaced optical and magnetic field spiral arms in the outer portion of a disc with a nearly flat rotation curve such as the case of the nearby spiral galaxy NGC 6946 (Beck \\& Hoernes 1996; Fan \\& Lou 1996; Lou \\& Fan 1998a, 2002; Frick et al. 2000, 2001; Lou et al. 2002).  In contexts of spiral galaxies, it is natural and important to consider a composite disc system consisting of gravitationally coupled gaseous and stellar discs. This is because various physical processes on different scales occur in the gaseous disc yet large-scale gas dynamics and environment are significantly affected by large-scale structures in the stellar disc (Lou \\& Fan 1998b). To construct global coplanar perturbation structures in a systematic manner, we started with a composite SID system of two fluid discs without magnetic field (Lou \\& Shen 2003) to generalize the work of Shu et al. (2000) for a single SID. In terms of the global axisymmetric stability for such a composite SID system, we provided a more straightforward $D-$criterion (Shen \\& Lou 2003) in contrast to the axisymmetric stability criteria of Elmegreen (1995) and Jog (1996). To be more general than a background SID profile, we recently constructed global stationary perturbation patterns in a composite system of two fluid scale-free discs without magnetic field (Shen \\& Lou 2004a, b) to further generalize the work of Syer \\& Tremaine (1996). With proper adaptations, these global stationary perturbation solutions may be utilized to model large-scale structures spiral galaxies.  While a scale-free disc has no characteristic scales, one may impose boundary conditions to describe a modified scale-free disc, e.g., by cutting a central hole in a disc (e.g., Zang 1976; Evans \\& Read 1998a, b). With such boundary conditions, discrete eigenfunctions for perturbations may be constructed as growing normal modes. On the other hand, by specifying a phase relation for a postulated reflection of spiral waves from the origin $r=0$, Goodman \\& Evans (1999) could also define discrete normal modes even for an unmodified gaseous SID. Shu et al. (2000) speculated that the swing amplification process (e.g., Goldreich \\& Lynden-Bell 1965; Julian \\& Toomre 1966; Toomre 1977; Fan \\& Lou 1997) across corotation allows a continuum of normal modes (Lynden-Bell \\& Lemos 1993) while proper `boundary conditions' may select from this continuum a discrete spectrum of unstable normal modes; they also suggested that zero-frequency (stationary) disturbances do signal the onset of instability.  On the basis of the work of Shu et al. (2000) for a single isopedically magnetized SID and of Lou \\& Shen (2003) for a composite SID system without magnetic field, Lou \\& Wu (2004) analyzed coplanar MHD perturbation structures in a composite SID system with one of the disc being isopedically magnetized (Shu \\& Li 1997). As X-ray emitting hot gases can freely flow along open magnetic field lines anchored at the gaseous disc, global stationary solutions constructed by Lou \\& Wu (2004) set an important preliminary stage for modelling spiral galactic MHD winds. In particular, if activities of star formation, starbursts and supernovae etc. are mainly responsible for producing X-ray emitting hot gases streaming out of the galactic plane from both sides, we would expect stronger winds from the circumnuclear starburst region (e.g. Lou et al. 2001) and along spiral arms than those winds from the remainder of the galactic disc plane.  Both SID and MSID models offer valuable physical insight for disc dynamics and belong to a wider family of scale-free discs (Syer \\& Tremaine 1996) where physical variables in the axisymmetric background equilibrium scale as powers of radius $r$ (e.g., the rotation curve $\\propto r^{-\\beta}$). It turns out that coplanar perturbations in such scale-free discs in a proper range of $\\beta$ can be treated globally and analytically (Lemos et al. 1991; Syer \\& Tremaine 1996) without invoking the usual WKBJ or tight-winding approximation to solve the Poisson equation for density wave perturbations\\footnote{By numerically solving the perturbed Poisson equation using the Fourier-Bessel transform, R\\\"udiger \\& Kitchatinov (2000) approached the linear stability of the disc in terms of an eigenvalue problem and could also treat the density wave problem globally. Their formulation differs from ours in two major aspects. Firstly, they include a central mass that mainly determines the rotation curve (i.e., an approximate Keplerian disc rotation for a disk mass being much less than the central mass). In our model, the disc rotation curve is completely controlled by the disc self-gravity and our approach is semi-analytical. Secondly, for axisymmetric instabilities, they only revealed ring fragmentation instabilities when the rotational Mach number becomes sufficiently high (which is similar to our result). However, our model analysis further indicates another type of collapse instabilities (i.e., magneto-rotational Jeans instability) in the perturbation regime of large radial spatial scales (see Section 3.2.1 for more details).} (Lin \\& Shu 1964; Binney \\& Tremaine 1987; Bertin \\& Lin 1996). In other words, it becomes feasible to include azimuthal magnetic fields in a scale-free disc and to construct global coplanar MHD perturbation configurations that are stationary in an inertial frame of reference. Such analytical solutions, still idealized and limited, are making important further steps and are extremely valuable for bench-marking numerical codes and for initializing numerical MHD simulations.  For an azimuthally magnetized scale-free disc with an infinitesimal thickness, the axisymmetric background equilibrium without radial flow is characterized by following features: a surface mass density $\\Sigma_0\\propto r^{-\\alpha}$, a rotation curve $v_0\\propto r^{-\\beta}$, an azimuthal magnetic field $B_0\\propto r^{-\\gamma}$ and a barotropic equation of state $\\Pi=K\\Sigma^n$ where $\\Pi$ is the vertically integrated gas pressure with constant $K>0$ for a warm disc and $n>0$ as the barotropic index. To maintain a radial force balance at all radii, there exists an explicit relationship among $\\alpha$, $\\beta$, $\\gamma$ and $n$ such that only one parameter is independent. As will be seen in Section 2, it is convenient to use $\\beta$ as this independent parameter (Syer \\& Tremaine 1996; Shen \\& Lou 2004b). In order to satisfy relevant constraints, the prescribed range of $\\beta$ turns out to be $-1/4<\\beta<1/2$ that includes the previous special case of $\\beta=0$ (i.e., the MSID case) studied by Lou (2002), Lou \\& Zou (2004a) and Lou \\& Wu (2004). In Section 2, we introduce coplanar MHD perturbations and present the linearized perturbation equations. We construct and analyze global stationary aligned and logarithmic spiral perturbation configurations in Section 3 and as an example of illustration, we discuss specifically the case of $\\beta=1/4$ with the background azimuthal magnetic field being force-free. In Section 4, we derive phase relationships between perturbation enhancements of surface mass density and magnetic field for both aligned and logarithmic spiral cases, and discuss the problem of angular momentum conservation. Finally, we summarize our results in Section 5. For the convenience of reference, relevant technical details are collected in the Appendices. ",
        "conclusions": "In this paper, we have explored and analyzed stationary perturbation structures in coplanarly magnetized razor-thin scale-free discs in a more general manner. To be specific, we start from a rotational and magnetized background equilibrium of axisymmetry that is dynamically self-consistent with a surface mass density $\\Sigma_0\\propto r^{-\\alpha}$, a rotation curve $v_0\\propto r^{-\\beta}$, a purely azimuthal (ring) magnetic field $B_0\\propto r^{-\\gamma}$ with a vertically integrated barotropic equation of state in the form of $\\Pi=K\\Sigma^n$. The radial force balance at all radii in an MHD disc implies several simple relationships among these power indices $\\alpha$, $\\beta$, $\\gamma$ and $n$, as stated explicitly in equation (\\ref{ParaRelation}). Without loss of generality, we can simply use $\\beta$ as an independent parameter to specify properties of an important class of radial variations for rotational MHD background equilibria. The allowed range of $\\beta$ falls within the interval $(-1/4,1/2)$ (Syer \\& Tremaine 1996; Shen \\& Lou 2004a, b) with the special case of $\\beta=0$ corresponding to a magnetized singular isothermal disc (MSID) system (Shu et al. 2000; Lou 2002; Lou \\& Fan 2002; Lou \\& Zou 2004a, b; Lou \\& Wu 2004). For clarity and convenience of our analysis, we introduced several dimensionless parameters. The first parameter is the partial disc parameter $F\\equiv\\phi/\\phi_T$ for the ratio of the gravitational potential ($\\phi$) arising from the disc to that ($\\phi_T$) arising from the entire system (Syer \\& Tremaine 1996; Shu et al. 2000; Lou 2002) including an axisymmetric halo mass distribution that is involved in the background equilibrium but is presumed to be unresponsive to coplanar MHD disc disturbances (a massive dark matter halo is an example in mind); for a partial disc system, we have $0<{F}<1$, while for a full disc system, we have ${F}=1$ by conventional definitions. The second parameter is an effective rotational Mach number $D$ defined as the ratio of the disc rotation speed to the sound speed yet with an additional scaling factor $(1+2\\beta)^{-1/2}$; $D$ is constant at all radii. The third parameter $q$ is a measure for the background ring magnetic field strength defined as the ratio of the azimuthal Alfv\\'en speed to the sound speed; and $q$ remains constant at all radii. For $1/4<\\beta<1/2$, there is no constraint on $q$ and the net Lorentz force of the background azimuthal magnetic field is radially outward (the magnetic pressure force is stronger than the magnetic tension force) to resist the self-gravitation; for $\\beta=1/4$, the background azimuthal magnetic field is force-free (the magnetic pressure and tension forces cancel each other exactly); for $-1/4<\\beta<1/4$, the net Lorentz force of the background azimuthal magnetic field points radially inward (the magnetic pressure force is weaker than the magnetic tension force) to aggravate the self-gravity and thus a restriction [i.e., inequality (\\ref{restrt})] is imposed on the magnetic field strength in order to guarantee that the surface mass density be positive in equation (\\ref{PropEquilib}). Inequality (\\ref{restrt}) provides a necessary physical criterion for plausible $D^2$ solutions.  With such a rotational MHD background equilibrium chosen, we introduce coplanar MHD disturbances to construct global stationary perturbation structures as viewed in an inertial frame of reference. Perturbation variables are expressed in terms of Fourier components for either aligned or logarithmic spiral pattern forms (Syer \\& Tremaine 1996; Shu et al. 2000; Lou 2002; Lou \\& Shen 2003; Shen \\& Lou 2003; Lou \\& Zou 2004a). We derive analytical solutions for stationary coplanar MHD perturbation configurations in both aligned and logarithmic spiral cases. We now summarize key results below.   \\begin{enumerate} \\item[(i)]{Cases of Aligned Coplanar MHD Perturbations} \\vskip 0.2cm  For aligned cases, we choose perturbation variables that bear the same power-law dependence in radius $r$ as the background MHD equilibrium does. For the special aligned case of $m=0$, the stationary coplanar MHD perturbations actually describe alternative equilibria to the axisymmetric background equilibrium with proper rescaling factors (Shu et al. 2000; Lou 2002; Lou \\& Zou 2004a) and are therefore somewhat trivial.  For aligned cases of $m\\ge 1$, we derive a quadratic equation (\\ref{aligned}) in terms of $y\\equiv D^2$. The resulting positive $D^2$ is the rotational parameter required for sustaining a stationary nonaxisymmetric MHD perturbation configuration. Other parameters involved in quadratic equation (\\ref{aligned}) include $\\beta$, $m$, ${F}$ and $q$. For a general combination of $\\beta$, $m$, ${F}$ and $q$, the determinant $\\Delta$ of quadratic equation (\\ref{aligned}) may not be always positive definite. On the other hand, for specific values of $\\beta$, equation (\\ref{aligned}) always gives two different real solutions (see Appendix A). Once we prescribe parameters $\\beta$, $m$, ${F}$ and $q$, the two solutions of $D^2$ can be readily determined. For physical solutions, $D^2$ must be positive and satisfy inequality (\\ref{restrt}) simultaneously.  As a specific case study, we examined the case of $\\beta=1/4$. Here the background azimuthal magnetic field remains force-free and the $D^2$ solution is free from constraint (\\ref{restrt}). It has been further shown in Appendix A that for this $\\beta=1/4$ case, equation (\\ref{aligned}) always has two distinct real solutions, the plus-signed solution $y_1$ and the minus-signed solution $y_2$, as contained in equation (\\ref{Solalign}). All these aspects contribute to a significant simplification of the relevant computational procedures. We derive two branches of $D^2$ solutions and need to require only $D^2\\ge 0$. There is yet a qualitative difference between the $m=1$ case and the $m\\ge 2$ cases. For $m=1$, there exists one critical value ${F}_c\\simeq 0.6842$; for ${F}$ below ${F}_c$, $y_1$ and $y_2$ are negative and positive, respectively, while ${F}$ above ${F}_c$, both branches of $D^2$ solutions are positive and thus physical. For $m\\ge 2$ in comparison, both branches of $D^2$ solutions remain always positive and thus physical. Moreover, $y_1$ remains always to be the upper branch for azimuthal stationary fast MHD density waves and represents the hydrodynamic counterpart of the unmagnetized single disc case; complementarily, $y_2$ remains always to be the lower branch for azimuthal stationary slow MHD density waves (Fan \\& Lou 1996; Lou \\& Fan 1998a) and is caused by the very presence of the azimuthal magnetic field (detailed computations reveal that $y_2$ must be smaller than $2/3$ which seems to be too small for galactic applications). When a disc rotates sufficiently fast, it could only support the upper $y_1$ stationary MHD configurations. An increase of magnetic field parameter $q$ will raise both physical branches of $D^2$ solutions as displayed in Fig. 1. The variation trend of the upper $y_1$ branch appears to be more sensitive. Therefore for a much stronger magnetic field, a sufficiently rapidly rotating disc may sustain stationary coplanar MHD configurations in terms of the upper $y_1=D^2$ branch.  \\vskip 0.2cm \\item[(ii)]{Cases of Logarithmic Spiral Configurations for Coplanar MHD Perturbations } \\vskip 0.2cm  For cases of stationary spiral perturbation structures, we choose coplanar MHD perturbations in terms of Kalnajs logarithmic spirals (Kalnajs 1971; Lemos et al. 1991; Shu et al. 2000; Lou 2002; Lou \\& Fan 2002; Lou \\& Shen 2003; Lou \\& Zou 2004a; Lou \\& Wu 2004). For the special $m=0$ case with radial propagations, it happens that inequality (\\ref{restrt}) is automatically satisfied and the stationary coplanar MHD configurations with $y\\equiv D^2>0$ obtained here represent marginal stability curves (Syer \\& Tremaine 1996; Shu et al. 2000; Shen \\& Lou 2003, 2004b; Lou \\& Zou 2004b). Only when the rotational parameter $D^2$ falls within a specific finite range can the magnetized disc be stable against axisymmetric disturbances at all wavelengths. A disc with too slow a rotation speed will succumb to Jeans' instability in the collapse regime corresponding to long wavelength perturbations, while a disc with too fast a rotation speed will suffer the ring-fragmentation instability corresponding to relatively short wavelength perturbations (Safronov 1960; Toomre 1964; Lemos et al. 1991; Lou \\& Fan 1998a, 2000; Shu et al. 2000; Shen \\& Lou 2003, 2004a; Lou \\& Zou 2004b). In comparison to a hydrodynamics disc system, the enhancement of azimuthal magnetic field tends to suppress ring-fragmentation instabilities as expected from the perspective of the MHD $Q$ parameter (Lou \\& Fan 1998a; Lou 2002; Lou \\& Zou 2004a, b), while an enhancement of azimuthal magnetic field tends to suppress Jeans' collapse instabilities for $-1/4<\\beta<1/4$, to aggravate Jeans' collapse instabilities for $1/4<\\beta<1/2$ and to bear no effect on Jeans' collapse instabilities when $\\beta=1/4$. We have provided intuitive interpretations for the roles of magnetic field earlier.  In parallel with the aligned cases, we derive quadratic equation (\\ref{spiral}) in terms of $y\\equiv D^2$ for logarithmic spiral cases of $m\\ge 1$. In addition to the four parameters $\\beta$, $m$, ${F}$ and $q$, we now have one more parameter $\\xi$ for the radial wavenumber. While the determinant $\\Delta$ of quadratic equation (\\ref{spiral}) may not be always positive in general, we prove in Appendix A that $\\Delta>0$ for the specific $\\beta=1/4$ case. Once the five parameters $\\beta$, $m$, ${F}$, $q$ and $\\xi$ are specified, the $D^2$ solutions can be readily obtained. For physical solutions, $D^2$ must be non-negative and also satisfy inequality (\\ref{restrt}).  As an example of illustration, we study the specific $\\beta=1/4$ case inequality (\\ref{restrt}) being satisfied automatically. Not surprisingly, the aligned cases and the logarithmic spiral cases parallel with each other very well, because of a mere change of wave propagation direction. For the special spiral case of $m=1$, there again exists the same critical ${F}_c\\simeq 0.6842$ for ${F}$ value as in the aligned $m=1$ case; for ${F} <{F}_c$, there is one critical point of $\\xi_c$ where $D^2\\equiv y_1$ diverges. More specifically, we have unphysical $y_1<0$ for $0<\\xi<\\xi_c$, while we have physical $y_1>0$ for $\\xi>\\xi_c$. There is no divergent point for $D^2$ solution branch and we always have $y_2\\ge 0$ ($y_2\\equiv 0$ when $q=0$). The solution structures become simpler for $m\\ge 2$ cases when both $D^2$ branches of solutions are non-negative and there is no divergent point of $D^2\\equiv y_1$. Moreover, $y_1$ remains always to be the upper branch for stationary spiral fast MHD density waves and is the hydrodynamic counterpart for those in the unmagnetized single disc case, while $y_2$ remains always to be the lower branch for stationary spiral slow MHD density waves (Fan \\& Lou 1996) and is caused by the very presence of the magnetic field (detailed computations reveal $y_2$ must be smaller than $2/3$).  \\vskip 0.2cm \\item[(iii)]{Phase Relationships among Coplanar MHD Perturbation Variables} \\vskip 0.2cm  So far, we have constructed global stationary coplanar MHD configurations for both aligned and logarithmic spiral cases in magnetized discs with a range of rotation curves. As expected, the inclusion of azimuthal magnetic field gives rise to one additional distinct physical solution for stationary slow MHD density waves to the problem (Fan \\& Lou 1996; Lou \\& Fan 1998a; Lou 2002; Lou \\& Zou 2004a); the Alfv\\'enic modes are excluded by our restricted consideration for coplanar MHD perturbations. The two $D^2$ branches of solutions are at least mathematically reasonable and may both have applications to astrophysical disc systems such as magnetized disc galaxies and so forth. These devised MHD disc problems are highly idealized with yet precious global analytical perturbation solutions, and conceptually still belong to a subclass of more general time-dependent MHD perturbation solutions. For applications to magnetized spiral galaxies, it would be of considerable interest to discuss pertinent physical aspects and identify relevant observational diagnostics.  For physical solutions of $D^2>0$, the $y_1$ branch would be applicable to MHD discs with relatively fast rotation while the $y_2$ branch would be applicable to MHD discs with relatively slow rotation in general. For spatial phase relationships, the two branches of $D^2>0$ solutions give different results. We take the specific $\\beta=1/4$ case as an example of illustration. For corresponding aligned cases, the plus-solution $y_1$ will lead to $\\hbox{i}R/S>0$ and the minus-solution $y_2$ will lead to $\\hbox{i}R/S<0$; there are no azimuthal magnetic field perturbations for aligned $\\beta=1/4$ cases. For tightly wound logarithmic spiral cases, perturbations of both radial and azimuthal magnetic field are approximately in phase with the perturbation enhancement of the surface mass density for the $y_1$ branch, while they are approximately out of phase for the $y_2$ branch. For open logarithmic spirals, the radial and azimuthal magnetic field perturbations are either ahead of or lagging behind the enhancement of the surface mass by a significant phase difference. These phase relationships together with those associated with flow perturbations (Visser 1980a, b) are valuable to interpret the spatial phase shifts between optical arms and magnetic arms in lopsided, barred or other spiral galaxies (Mathewson et al. 1972; Neininger 1992; Beck \\& Hoernes 1996; Beck et al. 1996; Fan \\& Lou 1996; Lou \\& Fan 1998, 2000, 2002; Frick et al. 2000, 2001; Lou et al. 2002).  \\vskip 0.2cm \\item[(iv)]{Angular Momentum Transfer in Steady Logarithmic Spirals for Coplanar MHD Perturbation Configurations} \\vskip 0.2cm  Stationary logarithmic spiral configurations of MHD density waves carry constant angular momentum flux either outward or inward associated with the advective transport, the gravity torque and the magnetic torque (Lynden-Bell \\& Kalnajs 1972; Goldreich \\& Tremaine 1978; Fan \\& Lou 1999). Therefore, the net angular momentum flux in the entire disc system is conserved and there must be either a source or a sink at the disc center. For the $y_1$ branch with $D^2>0$, we find that the total angular momentum flux is inward for leading (i.e., $\\xi>0$) spiral MHD density waves and is outward for trailing (i.e., $\\xi<0$) spiral MHD density waves. For the $y_2$ branch in comparison, the total angular momentum flux is outward for leading (i.e., $\\xi>0$) spiral MHD density waves and is inward for trailing (i.e., $\\xi<0$) spiral MHD density waves. As a real MHD disc system typically rotates at a relatively large $D^2$ and inequality (\\ref{restrt}) can also rule out small $D^2$ for $\\beta\\neq1/4$, the lower $y_2$ branch might be rarely useful in galactic applications. In either case, however, the net mass accretion rate is zero since if we focus on materials between two concentric circles, at any time the influent angular momentum from the inner ring is equal to the effluent angular momentum from the outer ring such that materials inside this belt neither gain nor lose any angular momentum and therefore no net mass accretion occurs.  \\end{enumerate}"
    },
    "0410/astro-ph0410691_arXiv.txt": {
        "abstract": "We present high-resolution VLT spectra of a new DO white dwarf and a new PG1159 star, which we identified in the ESO SPY survey. The PG1159 star is a low-gravity, extremely hot ($T_\\mathrm{eff}=160\\,000$\\,K, $\\log g=6$) star, having a C/He dominated atmosphere with considerable amounts of O and Ne (He=38\\%, C=54\\%, O=6\\%, Ne=2\\% by mass). It is located within the planetary nebula nuclei instability strip and pulsations have been discovered. The DO is a unique object. From \\ion{He}{i}/\\ion{He}{ii} line strengths  we found \\mbox{$T_\\mathrm{eff}$}$\\approx$60\\,000\\,K, however, the \\ion{He}{ii} lines are extraordinarily strong and cannot be fitted by any model. ",
        "introduction": "The ESO Supernovae Ia Progenitor Survey (SPY) is aimed at finding binary WDs  to test the double-degenerate scenario for SN~Ia progenitors (Napiwotzki et\\,al.\\ 2001). Here we report on the identification and spectrum analysis of a PG1159 star and a DO WD. HE1314+0018 (B=15.6\\,mag) and HE1429$-$1209 (B=15.8\\,mag)  were identified in the Hamburg ESO survey (HES; Wisotzki et\\,al.\\ 2000, Christlieb et\\,al.\\ 2001) as WD candidates and were therefore included in the SPY project. The spectra presented here were taken between April 2000 and July 2002. Their resolution is about $R=18\\,500$. Data reduction was performed with the ESO MIDAS software package.  Line blanketed non-LTE model atmospheres were computed using our {\\sc PRO2} code. The models assume plane-parallel geometry and hydrostatic and radiative equilibrium. ",
        "conclusions": ""
    },
    "0410/astro-ph0410372_arXiv.txt": {
        "abstract": "{ RX J0852.0-4622 is a supernova remnant discovered in the {\\it{ROSAT}} all-sky survey. Spatially coincident 1.157 MeV $\\gamma$-ray line emission was detected with the {\\it{COMPTEL}} instrument on-board of the CGRO. The analysis combining the X-ray and $\\gamma$-ray data suggests that  RX J0852.0-4622 is a close-by and young supernova remnant. Follow-up observations with {\\it{ASCA}} show that the two brightest sections of the limb have non-thermal spectra, which make an independent assessment of the age and distance using the Sedov equations for the evolution of the remnant almost impossible. We have observed three rim sections of RX J0852.0-4622 with {\\it{XMM-Newton}} and confirm the power law type spectra measured with {\\it{ASCA}}. We also confirm the presence of an emission line like feature at 4.45 $\\pm$ 0.05 keV, which we suggest to be emission from  Ti and Sc excited by atom/ion or ion/ion high velocity collisions. The high velocity is in agreement with the width of the 1.157 MeV $\\gamma$-ray line. The X-ray line flux expected from such an interaction is consistent with the 1.157 MeV $\\gamma$-ray line flux measured by {\\it{COMPTEL}}. This consistency of the X-ray line flux and the $\\gamma$-ray line flux lends further support to the existence and amounts of Ti in RX J0852.0-4622 claimed by Iyudin et al. (1998) and to the suggestion that RX J0852.0-4622 is young and nearby (Aschenbach et al. 1999). Iyudin et al. (1998) quote a very large broadening of the 1.157 MeV $\\gamma$-ray line which would indicate a large velocity of the emitting matter of about 15.000 km/s. Such high ejecta velocity for Ti is found only in explosion models of sub-Chandrasekhar type Ia supernovae (Woosley \\& Weaver 1994, Livne \\& Arnett 1995). In this case no compact remnant is expected. The obvious questions remaining are what the nature and the origin of the central compact source CXOU J085201.4-461753 are and why the absorption column density apparently associated with RX J0852.0-4622 is much greater than the typical column for the Vela SNR. ",
        "introduction": " ",
        "conclusions": "{\\it{XMM-Newton}} observations of three sections on the rim of RX J0852.0-4622 have been carried out. The north-western rim has been resolved in two clearly separated  filament-like structures. The high energy section of the spectra of both the north-western  rim and the southern rim are consistent with a power law shape and a power law which shows a roll-off towards high energies. Interestingly, the roll-off energy of about 1 keV is fairly high and it remains an open question what the spectrum looks like towards lower energies. The difficulty here is the contribution from the Vela SNR, which is fairly bright at energies below 1 keV. The various models seem to verify an absorbing column density towards RX J0852.0-4622 which is significantly higher than typical for the Vela SNR. Whether this implies a greater distance is not clear. The spectral slope measured with {\\it{XMM-Newton}} is basically the same as what has been measured with {\\it{ASCA}}. The western part of the remnant is different as it shows only a small contribution by a power law and if present at all the spectrum is much steeper. The {\\it{ASCA}} measurements indicated the presence of an emission line like feature at around 4.1 - 4.2 keV. The {\\it{XMM-Newton}} data confirm such a feature, and it seems to be present everywhere on the remnant's rim. The line energy averaged over the three observational fields is 4.45 $\\pm$ 0.05 keV. We attribute this line or lines to the emission of Ti and Sc which might be excited by atom/ion or ion/ion collisions. The X-ray line flux expected from such an interaction is consistent with the 1.157 MeV $\\gamma$-ray line flux measured by {\\it{COMPTEL}}. This consistency of the X-ray line flux and the $\\gamma$-ray line flux lends further support to the existence and amounts of Ti in RX J0852.0-4622 claimed by Iyudin et al. (1998) and to the suggestion that RX J0852.0-4622 is young and nearby (Aschenbach et al. 1999). Iyudin et al. (1998) quote a very large broadening of the 1.157 MeV $\\gamma$-ray line which would indicate a large velocity of the emitting matter of about 15.000 km/s. Such a high ejecta velocity for Ti is found only in explosion models of sub-Chandrasekhar type Ia supernovae (Woosley \\& Weaver 1994, Livne \\& Arnett 1995). In this case no compact remnant is expected. The obvious questions remaining are what the nature and the origin of the central compact source CXOU J085201.4-461753 are and why the absorption column density apparently associated with RX J0852.0-4622 is much greater than typical for Vela."
    },
    "0410/astro-ph0410528_arXiv.txt": {
        "abstract": "The extent to which the ISM in galaxies is well mixed is not yet settled. Measured metal abundances in the diffuse neutral gas of star--forming gas--rich dwarf galaxies are deficient with respect to that of the ionized gas. The reasons, if real, are not clear and need to be based on firm grounds. Far-UV spectroscopy of giant H\\2\\ regions such as NGC604 in the spiral galaxy M33 using \\textit{FUSE} allows us to investigate possible systematic errors in the metallicity derivation. We still find underabundances of nitrogen, oxygen, argon, and iron in the neutral phase by a factor of~$\\sim$6. This could either be explained by the presence of less chemically evolved gas pockets in the sightlines or by dense clouds out of which H\\2\\ regions form. Those could be more metallic than the diffuse medium. ",
        "introduction": "The \\textit{FUSE} observations (LWRS and MDRS apertures) allow us to determine the chemical composition while the HST/\\textit{STIS} spectrum gives the possibility to map the neutral gas inhomogeneities and investigate possible multiple line of sight effects.      \\subsection{Data analysis}    Data analysis has been performed using the profile fitting procedure \\texttt{Owens} developed at the Institut d'Astrophysique de Paris by Martin Lemoine and the \\textit{FUSE} French Team. This program returns most likely values of many free parameters such as heliocentric velocities, turbulent velocities, or column densities by a $\\chi^2$ minimization of absorption lines profiles. Errors on parameters include uncertainties on all the free parameters, in particular the position and shape of the continuum.    We checked the shape of the adopted continuum by comparing the observed spectrum with a theoretical model of young stellar populations. No significant difference is found between the model and the continuum we adopted for the profile fitting. Moreover, by comparing the data with the model, we find no significant contamination of neutral interstellar lines  by stellar atmospheres.    By investigating the two \\textit{FUSE} observations, we show that an additional broadening of the lines can account for the spatial distribution of the bright sources within the slit. The column densities we derive account for this extension. For the first time, we could check saturation effects for O\\1\\ and Fe\\2\\ lines by analyzing lines independently. We find no correlation between column densities derived from each line with the oscillator strength, implying that saturated O\\1\\ and Fe\\2\\ lines in those spectra do not give systematic errors.    \\subsection{Neutral gas inhomogeneities}  The high spatial resolution of the HST/\\textit{STIS} spectrum of NGC604 gives the possibility to extract spectra towards individual stars of the ionizing cluster (Bruhweiler et al. 2003).  So far, we have analyzed three sightlines from which we have measured H\\1\\ column density using the Lyman $\\alpha$ line. We find spatial variations (up to 0.4 dex) suggesting inhomogeneities of the diffuse neutral gas. This could be a source of systematic errors when determining global column densities from a spectrum of a whole cluster.  We built the global spectrum (i.e. the mean spectrum of all sightlines, weighted by star magnitudes) of the HST/\\textit{STIS} observation to mimic the spectrum of a cluster in order to compare the real mean column density we want to determine with the weighted mean we actually measure. Preliminary results show that we could tend to underestimate the actual column density when analyzing a global spectrum of several sightlines towards clouds having different physical properties. ",
        "conclusions": ""
    },
    "0410/astro-ph0410002_arXiv.txt": {
        "abstract": "The problem of the late accretion phase of the evolution of an axisymmetric, isothermal magnetic disk surrounding a forming star has been formulated in a companion paper. The ``central sink approximation'' is used to circumvent the problem of describing the evolution inside the opaque central region for densities greater than $10^{11}$ $\\rm{cm}^{-3}$ and radii smaller than a few AUs. Only the electrons are assumed to be attached to the magnetic field lines, and the effects of both negatively and positively charged grains are accounted for.  After a mass of $0.1$ $M_{\\odot}$ accumulates in the central cell (forming star), a series of magnetically driven outflows and associated outward propagating shocks form in a quasi-periodic fashion. As a result, mass accretion onto the protostar occurs in magnetically controlled bursts. We refer to this process as spasmodic accretion. The shocks propagate outward with supermagnetosonic speeds. The period of dissipation and revival of the outflow decreases in time, as the mass accumulated in the central sink increases. We evaluate the contribution of ambipolar diffusion to the resolution of the magnetic flux problem of star formation during the accretion phase, and we find it to be very significant although not sufficient to resolve the entire problem yet. Ohmic dissipation is completely negligible in the disk during this phase of the evolution. The protostellar disk is found to be stable against interchange-like instabilities, despite the fact that the mass-to-flux ratio has temporary local maxima. ",
        "introduction": "In an accompanying paper (Tassis \\& Mouschovias 2004, hereafter Paper I) we formulated the problem of the formation and evolution of a nonrotating protostellar fragment in a manner that allows accurate modeling of the physics of the protostellar disk without the complication introduced by radiative transfer when the central protostar becomes opaque. In the present paper we follow the evolution until the mass of the central protostar grows to $1$ $M_{\\odot}$. The importance and relevance of this problem as well as previous work on the subject were summarized in \\S 1 of Paper I. The evolution of the system is described by the six-fluid MHD equations, in which neutral molecules, ions, electrons, neutral grains, positively and negatively charged grains are treated as distinct but interacting fluids. We use the ``central sink'' method to study the structure and evolution of the isothermal disk surrounding the forming protostar. The effects of the mass and magnetic flux accumulating in the central, opaque protostar on the isothermal disk are accounted for.  We include important physics ignored by previous calculations: (1) the decoupling of the ions from the magnetic field lines, which occurs at densities above $10^8 \\, {\\rm cm ^{-3}}$ (Desch \\& Mouschovias 2001) and which leaves the magnetic field frozen only in the much less massive and much more tenuous electron fluid; (2) the chemical and dynamical effects of the positively-charged grains, the abundance of which becomes significant in the same density regime. Phenomena that arise due to this new physics are discovered, such as the formation and dissipation of a series of shocks in a quasi-periodic fashion.  In \\S 2 we summarize the most important results of the contraction phase prior to the introduction of the central sink. The evolution at later times, in the presence of a central sink, is described in \\S 3. Specifically, we discuss: the evolution of the system until the first electron outflow occurs (\\S 3.1); the establishment of a quasi-periodic magnetic cycle (\\S 3.2) and the properties of one typical such cycle (\\S 3.3) ; the time evolution of the mass, magnetic flux, and mass-to-flux ratio in the central protostar (\\S 3.4); the properties of the shock in the neutral fluid (\\S 3.5); and the evolution of the period of the magnetic cycle (\\S 3.6). Finally, we investigate the stability of the supercritical core against a magnetic interchange instability (\\S 3.7). The most important conclusions are summarized in \\S 4. ",
        "conclusions": "We have studied the structure of the magnetic accretion disk surrounding a forming star, using the central-sink approximation to make the problem tractable numerically. The differences in the formulation of the problem compared to previous work (Ciolek \\& K\\\"{o}nigl 1998) have been discussed in detail in Paper I.  Multiple shocks develop and propagate in the neutral fluid, associated with a quasi-periodic magnetic cycle. Flux accumulation close to the central sink leads to the formation of a region of enhanced magnetic field strength, a {\\em ``magnetic wall''}, where magnetic forces exceed the gravitational forces. The magnetic wall drives outflows in the neutrals and associated shocks. In time, the magnetic wall disperses, and the neutrals are reaccelerated inward behind the shock (at smaller radii), dragging with them the electrons and, consequently, magnetic flux (since the field is frozen in the electrons). Flux can then reaccumulate, and the magnetic wall gets recreated, which causes the magnetic cycle to repeat.  The shocks in the neutrals propagate outward with supermagnetosonic speeds in the rest frame of the central sink. This is in contrast to the slow accretion shock expected by Li \\& McKee (1996) and found numerically by Ciolek \\& K\\\"{o}nigl (1998) because they assumed that the magnetic flux is frozen in the ion fluid. The infall velocities of the neutrals in the rest frame {\\em of the shock} vary between 2 and 6 times the magnetosonic speed.  The development and propagation of multiple outflows and shocks result in mass accretion onto the forming star that occurs in {\\em magnetically controlled bursts}. We refer to this phenomenon as {\\em spasmodic accretion}. During the outflow of the neutrals, the inner part of the disk is almost ``emptied'' of matter (its column density decreases significantly) and the mass accretion rate drops almost by two orders of magnitude below its maximum value. It is only when the neutrals get reaccelerated inward by the mass of the forming star that the inner part of the disk is refilled with matter and the mass accretion rate increases again. The time required for 1 ${\\rm M_\\odot}$ to accumulate in the central sink is $2.55 \\times 10^5 \\, {\\rm yr}$, characteristic of the relatively inefficient accretion process due to the magnetically-driven, repeated outflows in the neutrals.  The radial extent of the region which is affected by the magnetic cycle (in which the velocity of the neutrals exhibits the quasi-periodic oscillation characteristic of the magnetic cycle) includes almost the entire magnetically supercritical core. The magnetically supported envelope, as expected, is not affected by the magnetic flux redistribution in the core.  The period of the magnetic cycle decreases in time, and is well represented by the local flux-loss (ambipolar-diffusion) timescale. This is expected because the bottleneck in the chain of events comprising the magnetic cycle is the rebuilding of the magnetic wall behind the outward propagating shock in the neutrals. This reaccumulation of flux does indeed occur on the ambipolar-diffusion timescale, and it becomes faster as time progresses, since the gravitational field (dominated by the central point mass) monotonically increases with time, thereby increasing the maximum infall velocity of the neutrals (which in turn drag the electrons inward).  We found that the magnetically supercritical core is stable against interchange-type instabilities. This result is in contrast to the predictions by Li \\& McKee (1996) and Ciolek \\& K\\\"{onigl} (1998). The reason for this difference with the results of Ciolek \\& K\\\"{o}nigl is the fact that they assumed flux-freezing in the ions, which is not valid above $\\sim 10^8$ ${\\rm cm^{-3}}$ and which leads to much better coupling between the neutrals and the magnetic field in their calculation.  Ambipolar diffusion by itself during the accretion phase significantly contributes to the resolution of the magnetic flux problem of star formation by increasing the mass-to-flux ratio of the protostar by more than two orders of magnitude above its critical value. However, this is not sufficient to resolve the entire magnetic flux problem yet."
    },
    "0410/nucl-th0410066_arXiv.txt": {
        "abstract": "The roles of isospin asymmetry in nuclei and neutron stars are investigated using a range of potential and field-theoretical models of nucleonic matter. The parameters of these models are fixed by fitting the properties of homogeneous bulk matter and closed-shell nuclei.  We discuss and unravel the causes of correlations among the neutron skin thickness in heavy nuclei, the pressure of beta-equilibrated matter at a density of 0.1 fm$^{-3}$, the derivative of the nuclear symmetry energy at the same density and the radii of moderate mass neutron stars.  Constraints on the symmetry properties of nuclear matter from the binding energies of nuclei are examined.  The extent to which forthcoming neutron skin measurements will further delimit the symmetry properties is investigated.  The impact of symmetry energy constraints for the mass and moment of inertia contained within neutron star crusts and the threshold density for the nucleon direct Urca process, all of which are potentially measurable, is explored.  We also comment on the minimum neutron star radius, assuming that only nucleonic matter exists within the star. ",
        "introduction": "Strongly interacting matter in which more neutrons than protons exist is encountered in both heavy nuclei and neutron stars.  In stable nuclei the net asymmetry $I=(N-Z)/(N+Z)$ ranges up to about 0.24, but the neutron-to-proton asymmetry, $\\delta = (n_n-n_p)/n$ where $n_n$ and $n_p$ are the number densities of neutrons and protons and $n=n_n+n_p$, approaches unity in the nuclear surface. In the future, rare-isotope accelerator experiments will extend the range of $I$ to values well in excess of 0.2. In contrast, $I$ could be as large as 0.95 in the interiors of neutron stars. The physical properties of nuclei, such as their masses, neutron and proton density distributions (including their mean radii), collective excitations, fission properties, matter and momentum flows in high energy heavy-ion collisions, etc. all depend on the isospin structure of the strong interactions between nucleons (i.e., $nn$ and $pp$ interactions versus $np$ interactions).  The energetics associated with the $n-p$ asymmetry can be characterized by the so-called symmetry energy, $E_{sym}$, which is the leading coefficient of an expansion of the total energy with respect to asymmetry: $E(n,\\delta)\\approx E_0(n)+E_{sym}\\delta^2\\cdots$.  The energy $\\hat{\\mu} = \\mu_n - \\mu_p \\cong 4 E_{sym}\\delta$, where $\\mu_n$ and $\\mu_p$ are the neutron and proton chemical potentials, respectively, is crucial in determining reaction rates involving electrons and neutrinos, particle abundances, etc., in astrophysical contexts such as supernova dynamics, proto-neutron star evolution, the $r-$process, the long-term cooling of neutron stars, and the structure of cold-catalyzed neutron stars (i.e. their masses, radii and crustal extent), etc.  The pervasive role of the isospin dependence of strong interactions in nuclear processes in the laboratory and the cosmos is sketched in Fig. 1. In this work some of these connections will be discussed.   \\begin{figure}[htb] \\begin{center} \\includegraphics[scale=0.65,angle=0]{steiner_fig1.eps} \\end{center} \\caption{The multifaceted influence of the nuclear symmetry energy.} \\label{figinfluence} \\end{figure}  Recently, several empirical relationships have been discovered that underscore the role of isospin interactions in nuclei and neutron stars. These include correlations between:  \\begin{enumerate}  \\item {\\em The neutron star radius $R$ and the pressure $P$ of neutron-star matter}: Lattimer and Prakash \\cite{Lattimer00,Lattimer01} found that the quantity $RP^{-1/4}$ is approximately constant, for a given neutron star mass, for a wide variety of equations of state when the pressure $P$ of beta-equilibrated neutron-star matter is evaluated at a density in the range $n_0$ to $2n_0$, where $n_0$ denotes equilibrium nuclear matter density. Since the pressure of nearly pure neutron matter (a good approximation to neutron star matter) near $n_0$ is approximately given by $n^2\\partial E_{sym}/\\partial n$, the density dependence of the symmetry energy just above $n_0$ will be a critical factor in determining the neutron star radius.  \\item {\\em The neutron skin thickness in nuclei and the pressure of pure neutron matter at sub-nuclear density:} Typel and Brown \\cite{Brown00,Typel01} have noted that model calculations of the difference between neutron and proton rms radii $\\delta R ={\\langle r_n^2\\rangle}^{1/2} - {\\langle r_p^2\\rangle}^{1/2}$ are linearly correlated with the pressure of pure neutron matter at a density below $n_0$ characteristic of the mean density in the nuclear surface (e.g., 0.1 fm$^{-3})$.  The density dependence of the symmetry energy controls $\\delta R$ (we will call this the neutron skin thickness) in a heavy, neutron-rich nucleus.  Explicitly, $\\delta R$ is proportional to a specific average of $[1-E_{sym}(n_0)/E_{sym}(n)]$ in the nuclear surface, see Refs.~\\cite{Krivine84,Lattimer96} and Eq. (\\ref{eq:newt}) below.  \\item {\\em The neutron skin thickness in nuclei and the neutron star radius:} Horowitz and Piekarewicz \\cite{Horowitz01} have pointed out that models that yield smaller neutron skins in heavy nuclei tend to yield smaller neutron star radii. These authors, along with others \\cite{Horowitz01b,Michaels00}, have also pointed out the need for an accurate measurement of the neutron skin. \\end{enumerate}   Unlike proton distributions, neutron distributions in nuclei have remained uncertain to this date.  Recent studies of neutron densities from a global analysis of medium-energy proton scattering on $^{208}$Pb indicate that $0.07 < \\delta R < 0.16$ fm \\cite{Clark03}. Related information is also available from an analysis of antiprotonic atom data that gives $\\delta R = 0.15\\pm 0.02$ fm~\\cite{Trzcinska01}. In the latter work, nucleon density distributions are parameterized by Fermi functions and it is found that the half-density radii for neutrons and protons in heavy nuclei are the same, but the diffuseness parameter for the neutrons is larger than that for the protons. Skin thicknesses as large as 0.2 fm were obtained in earlier analyses; see the discussion of Karataglidis et al. \\cite{Karataglidis02}. Since these studies involve strongly interacting probes, even to this date the value of $\\delta R$ for a nucleus such as $^{208}$Pb is not accurately known. This situation should improve as it is expected that the neutron rms radius will be determined to about 1\\% accuracy by measuring the parity-violating electroweak asymmetry in the elastic scattering of polarized electrons from $^{208}$Pb \\cite{Horowitz01b}, an experiment planned at the Jefferson Laboratory \\cite{Michaels00}.  A common theme in the evaluation of the neutron skin thickness and the pressure of neutron star matter below and above the saturation density $n_0$ is the density dependence of the isospin asymmetric part of the nuclear interaction.  To highlight this dependence, we examine theoretical predictions of $\\delta R$ and the neutron star radius for an extensive set of potential and field-theoretical models.  For the former we employ Skyrme-like potential model interactions, including a parameterization~\\cite{Akmal98} of the microscopic calculations of Akmal and Pandharipande~\\cite{Akmal97}. For the latter field-theoretical models we use Walecka-type models~\\cite{Serot86} in which nucleons interact through the exchange of $\\sigma$, $\\omega$ and $\\rho$ mesons augmented with additional mesonic couplings (see, for example, Horowitz and Piekarewicz~\\cite{Horowitz01}) to better describe the properties of nuclei and explore the properties of neutron stars.  In order to provide baseline calculations, our investigations will be restricted to the case in which the high density phase of neutron stars contains nucleons (and enough electrons and muons to neutralize the matter) only.  Our goals in this work are to identify the causes of the correlations mentioned above through a semi-analytic analysis of detailed calculations and to uncover other possible correlations.  Toward these goals, we investigate an assortment of potential and field-theoretical models that cover a range of symmetry properties and incompressibilites, but which are constrained by the empirical properties of isospin symmetric and asymmetric nuclear matter, and the binding energies and radii of closed shell nuclei. The properties of nuclei are investigated through Hartree-Fock-Bogoliubov calculations for potential models and Hartree calculations for field theoretical models.  In both categories, our calculations include models that match the Akmal and Pandharipande~\\cite{Akmal97} equation of state (EOS) as results for nuclei with this model are not available in the literature.  Additional insights are provided through analytical and numerical analyses of isospin asymmetric semi-infinite matter in the potential and field theoretical approaches.  The results of these calculations are utilized in understanding several correlations in conjunction with inferences from nuclear masses.  The neutron skin-pressure and neutron star radius-pressure correlations are reexamined to provide improved relations.  Our calculations also provide an estimate of the smallest neutron star radii that ensue from these models.  The presentation is organized as follows.  In Sec. II, the Hamiltonian (Lagrangian) densities of the potential (field-theoretical) models used in this work are described. The variational approach to determine the bulk and surface properties of isospin asymmetric semi-infinite nucleonic matter in these models is developed in Sec. III. This section also contains a discussion of how these properties enter in the liquid droplet model description of nuclei. Our results and discussion of the various correlations, and their origins are presented in Sec. IV. Section V contains a brief description of other related correlations.  Discussion and conclusions are contained in Sec. VI. Appendix A lists the properties of the models used in this work. The coupling strengths of some newly-constructed models are given in Appendix B. ",
        "conclusions": "In this work, we have investigated the relationship between the symmetry energy of nucleonic matter and the neutron skin thicknesses of neutron-rich nuclei as well as the radii of neutron stars.  Precision measurements of neutron skins and neutron star radii are due to become available in the near future.  We have studied how these measurements can constrain the dependence of the symmetry energy on baryon density, $E_{sym}(n)$, in the vicinity of the nuclear saturation density $n_0$. This knowledge is crucial to understanding many astrophysical phenomena, including neutron star evolution, supernova explosions and nucleosynthesis, and binary mergers involving neutron stars. Simulations of these phenomena involve extrapolations of the symmetry energy to supranuclear densities.  A crucial aspect of our presentation is that care was taken to ensure that the nuclear force parameterizations, whether for potential or field-theoretical models, were constrained by fits to closed-shell nuclei.  In addition, all models were constrained to yield a maximum neutron star mass of at least 1.44M$_\\odot$, the larger of the accurately measured neutron star masses in the binary pulsar 1913+16. Without such constraints, much weaker correlations relating the neutron skin thickness to astrophysical quantities such as neutron star radii would have resulted.  Analytical representations of the relation between the symmetry energy and the neutron skin were developed for semi-infinite surfaces calculated in potential and field-theoretical models.  To lowest order, the neutron skin thickness $\\delta R ={\\langle r_n^2\\rangle}^{1/2} - {\\langle r_p^2\\rangle}^{1/2}$ is proportional to $\\delta_L S_s/S_v$ where $\\delta_L$ is the neutron excess in the nuclear center and $S_v [\\equiv E_{sym}(n_0)]$ and $S_s$ are related to the expansion parameters of the volume and surface symmetry energies, respectively, in the liquid drop or droplet models of nuclei.  This representation is validated by comparisons with results of finite nucleus calculations performed in the Hartree-Fock-Bogoliubov (for potential models) and Hartree (for field-theoretical models) approximations.  For the first time, results for nuclei are presented for the Akmal-Pandharipande-Ravenhall EOS parameterized according to both potential and field-theoretical models.  The semi-infinite surface representation predicts that the ratio $S_s/S_v$, or equivalently, the skin thickness $\\delta R$, is a particular average of the density-dependent symmetry energy in the surface.  In addition, phenomenological comparisons indicate that there is a relatively tight correlation between $\\delta R$ and the pressure of neutron star matter at a typical average surface density of $5n_0/8$.  These results are not independent, but together imply that the determination of $\\delta R$ offers a valuable constraint on the density dependence of the symmetry energy.  It was explicitly shown that the $\\delta R$ -- pressure correlation is a consequence of a more general correlation between $\\delta R$ and the derivative of the symmetry energy at the same density. In contrast to most, but not all, previous results, we have also demonstrated that field theoretical models can yield a neutron skin thickness in $^{208}$Pb of less than 0.2 fm.  Since a measurement of the neutron radius to about 1\\% accuracy implies an error of about 0.05 fm, the uncertainty in the neutron skin thickness will probably range from 10\\% to 30\\%, depending upon the magnitude of $\\delta R$.  On this basis, it is expected that the pressure at 5/8 $n_0$ could be determined to 20--50\\%.  An independent method of constraining $S_s/S_v$ is from a least squares fit of nuclear models to nuclear binding energies.  However, this fitting can only reliably establish a correlation between $S_s/S_v$ and $S_v$.  We have examined this in the liquid droplet approach using two plausible models.  The surface tension was parameterized either in terms of $\\mu_n$, implying a neutron skin, or in terms of $\\mu_n-\\mu_p$ which allowed for the Coulomb interaction of protons. While these models gave fits to the data of similar accuracy, the valley in $\\chi^2$ differed in the two cases. The behavior of the `$\\mu_n$' model was shown to be more consistent with the systematic correlation derived from parameter optimization of both potential and field-theoretical models of laboratory nuclei.   The density dependence of the symmetry energy has a number of astrophysical consequences.  One of the most important is its role in determining the composition of matter.  In beta equilibrium, the proton fraction is proportional to $E_{sym}(n)^3$ for small proton fractions.  The charge fraction of matter plays an essential role in establishing the threshold densities of hyperons and the quark-hadron phase transition.  If any of these threshold densities are exceeded in a neutron star, the possibility of enhanced neutrino emission by a direct Urca process exists.  The threshold for the direct Urca process involving nucleons alone is also determined by the charge fraction (for example, it is 1/9 in the absence of muons).  We have shown that the determination of $\\delta R$ to have a value greater than approximately 0.2 fm would imply that at least the nucleon direct Urca process exists in 1.4 M$_\\odot$ neutron stars.   An additional astrophysical application is the relation we established between $\\delta R$ and the radii $R$ of neutron stars.  Although Lattimer and Prakash established a phenomenological relation between $R$ for a given mass star and the pressure of neutron star matter above, but near, the nuclear saturation density, there exists additional uncertainty in relating this algebraically to the $\\delta R - P$ correlation.  However, as shown in this work, a useful relation directly relating $\\delta R$ and, for example, the radius of a 1.4M$_\\odot$ star, $R_{1.4}$, can be established.   The fact that the neutron skin thickness, which measures symmetry properties below nuclear density, and the astrophysical quantities involving the neutron star radius and the Urca threshold density, can be correlated is a consequence of a generic trend in the overall density dependence of the symmetry energy.  This trend was specifically demonstrated by displaying the linear relationship which exists between the density derivatives of the symmetry energy at $5n_0/8$ and $3n_0/2$.  This relationship is independent of the parameterization of the nuclear force.  Observation of neutron star radii may provide qualitative information on the stellar composition. Using the parametric freedom of the EOS's considered, while satisfying the aforementioned constraints, we concluded that the minimum radius achievable for a star composed of just nucleons and leptons is about 9 km. Observation of a significantly smaller radius would likely imply that some softening component (bose condensates, hyperons, quarks, etc.) is present in dense matter, although it could also indicate that the EOS's given by potential and field-theoretical models is qualitatively incorrect.  We have also explored the influence of the density dependence of the nuclear symmetry energy on the transition pressure marking the boundary between the core and the crust in a neutron star.  This transition pressure is closely approximated by the phase boundary separating uniform matter from matter in which nuclei exist.  It was found that the transition pressure was not correlated with the neutron skin thickness, although the transition density did show some rather weak correlation. More interesting is the fraction of the star's moment of inertia residing in the crust which is measurable and which depends on the transition pressure, mass and radius. Here there was some correlation with larger values of this fraction tending to imply larger values of the neutron skin thickness, but the correlation was not very robust. Probably the neutron skin and the transition pressure should be viewed as independent quantities whose measurement would provide separate constraints on the EOS."
    },
    "0410/astro-ph0410552_arXiv.txt": {
        "abstract": "We present a new measurement of atmospheric muons made during an ascent of the High Energy Antimatter Telescope balloon experiment.  The muon charge ratio $\\mu^{+}/\\mu^{-}$ as a function of atmospheric depth in the momentum interval 0.3--0.9 GeV/c is presented.  The differential $\\mu^{-}$ intensities in the 0.3--50 GeV/c range and for atmospheric depths between 4--960 g/cm$^{2}$ are also presented.  We compare these results with other measurements and model predictions.  We find that our charge ratio is $\\sim$1.1 for all atmospheric depths and is consistent, within errors, with other measurements and the model predictions.  We find that our measured $\\mu^{-}$ intensities are also consistent with other measurements, and with the model predictions, except at shallow atmospheric depths. ",
        "introduction": "Measurements and theoretical calculations of atmospheric neutrino fluxes have consistently disagreed, and this disagreement was interpreted in terms of neutrino oscillations. However, this interpretation requires an accurate understanding of neutrino fluxes in the atmosphere. Simulations of the absolute fluxes have been conducted (\\cite{barr:muons, wentz:muons, gaisser:muons}, and references therein) but suffer from systematic uncertainty in the normalization of the atmospheric neutrino spectrum.  These simulations also predict the spectrum of other atmospheric secondaries, specifically muons.  Indeed, the production of neutrinos in the atmosphere is closely coupled with that of muons, as they are produced together in pion and kaon decays, and as some muons themselves decay and contribute to the neutrino flux. Therefore, measurements of atmospheric muons by ascending high altitude balloon borne instruments can be used to reduce the neutrino model uncertainties. It should be pointed out that other experimentally accessible quantities can also be used to that end, such as the absolute primary cosmic ray flux impinging on the atmosphere, or other atmospheric secondaries such as gamma rays or secondary nucleons. The atmospheric muon flux is therefore an important element in any comprehensive strategy to improve our understanding of neutrino production in the atmosphere.  An earlier version of the High Energy Antimatter Telescope (HEAT-e$^\\pm$) instrument had been used to measure air-shower muons during atmospheric ascent~\\cite{coutu:muons}. The HEAT instrument is described elsewhere~\\cite{barwick:heat, beach:heat}. In its present configuration, HEAT-pbar is optimized to study antiprotons. It combines a superconducting magnet spectrometer using a drift-tube hodoscope (DTH), a time-of-flight system (TOF), and two stacks of multiwire proportional chambers (dE/dx). We report here a new measurement of the muon charge ratio $\\mu^{+}/\\mu^{-}$ as a function of atmospheric depth in the momentum interval 0.3--0.9 GeV/c, and differential $\\mu^{-}$ intensities in the 0.3--50 GeV/c range and for atmospheric depths between 4--960 g/cm$^{2}$ for a balloon flight from Fort Sumner, NM, USA on June 3, 2000. ",
        "conclusions": "We have measured with good statistics the muon charge ratio $\\mu^{+}/\\mu^{-}$ as a function of atmospheric depth in the momentum interval 0.3--0.9 GeV/c and the differential $\\mu^{-}$ intensities in the 0.3--50 GeV/c range and for atmospheric depths between 4--960 g/cm$^{2}$. We have found that our charge ratio is $\\sim$1.1 for all atmospheric depths and is consistent, within errors, with other measurements and the model predictions.  We have found that our measured $\\mu^{-}$ intensities are also largely consistent with other measurements, and with the model predictions, despite varying solar epochs and geomagnetic rigidity cutoffs. This reinforces the conclusion of \\cite{coutu:muons} that model calculations are adequate to the task of predicting neutrino fluxes. The discrepancy between our measurements and the model predictions at shallow atmospheric depths is of little import to our understanding of neutrino fluxes, as the bulk of the neutrinos come from decay processes much deeper in the atmosphere."
    },
    "0410/astro-ph0410287_arXiv.txt": {
        "abstract": "It has been suggested that many hot subdwarfs lurk in the pile of rejected UV-excess candidate stars from the Palomar-Green (PG) survey. This suggestion is not supported by available photometric data. ",
        "introduction": "During the PG survey for ultraviolet excess (UVX) objects, candidate UVX objects were those with (transformed) $U-B < -0.46$.  It was recognized that the large error in $U-B$, $\\sigma\\approx 0.38$, meant that color selection should be supplemented by spectroscopy for classification, since more accurate temperature information was likely available from the spectra than from $U-B$. Many {\\em candidate} UVX targets were indeed culled from the final PG catalog (Green et al. 1986: GSL86), because their classification spectra showed the Ca II K line in absorption. These K-line stars were thought not to be genuinely hot, but rather to be metal-poor subdwarf F or G stars that crept into the candidate list owing to a combination of low metal-line blanketing and photometric errors.  Hot subdwarf stars (especially sdB stars, but including some sdO stars) are understood to belong to the Extended Horizontal Branch (EHB), and as such are core-helium burning objects with very thin hydrogen envelopes. Recently, there has been renewed interest in scenarios of the origin of hot subdwarf stars that involve binary star processes (Roche-lobe overflow, common-envelope evolution) to strip the hydrogen-rich envelope away from the helium core, near the time of He-core ignition (Han et al.\\ 2002, 2003).  In comparing their population synthesis models with observations, Han et al.\\ (2003) drew attention to the K-line stars rejected from the PG catalog, as possibly representing a ``missing'' group of hot subdwarfs, hidden by their cooler and (somewhat) brighter binary companions.  Owing to dilution of the hot star's energy distribution by the companion, the $U-B$ color could be marginal for the PG color criterion, while the cool star would contribute a K line, much as a metal-poor subdwarf would show.  In this interpretation, therefore, these ``PG--rejects'' actually belong in the PG catalog, and moreover would constitute important evidence in favor of binary formation channels for sdB.  Han et al.\\ have put forward a hypothesis that can be tested.  The list of PG--rejects exists in a card file (with finding charts) kept by RFG.  Here we report our investigation to date of the PG--rejects. ",
        "conclusions": "A very few objects in our sample of PG--reject stars may plausibly be binary systems with a hot subdwarf star component.  Also, a few objects seem to have entered the PG--reject list by accident. The vast majority of the PG--reject stars, however, are sufficiently modeled as single stars consistent with their being the metal-poor sdF/sdG contaminants that GSL86 were guarding against. The color-color sequences of sdB + cool star binaries (with $M_V$ for main sequence companions!)  are well separated from the observed colors of the PG--reject stars.  There is at present no compelling evidence for large numbers of additional hot subdwarf stars hiding in binaries that were rejected from the PG catalog."
    },
    "0410/astro-ph0410414_arXiv.txt": {
        "abstract": "{Active galaxies are often associated with starbursts, which are possible sources of fuelling for the nuclear black holes. By means of optical ($3500-7500$ \\AA) spectroscopic and narrow-band photometric data (H$\\alpha$ + [NII]$\\lambda$6548,6583 and [OIII]$\\lambda$5007) it is possible to detect starburst features in active nuclei and to study the correlation with Seyfert type and the surrounding environment. Here we present evidences of young stellar populations (HI high order Balmer lines) and extended H$\\alpha$ emission (indicative of HII regions) in 6 nearby Seyfert galaxies ($z < 0.02$). These very preliminary results seem to suggest that  the nuclear starburst presence does not depend on the Seyfert type, favouring the hypothesis of the AGN Unified Model. By estimating the present SFR ($< 10^6$ years) and recent SFR ($< 10^9$ years) we observe that these objects have experienced most of star formation episodes in past epochs. This supports the hypothesis of a starburst-AGN connection, favouring an evolutionary scenario in which the AGN survives to the starburst. } \\authorrunning{A. Romano et al.} \\titlerunning{} ",
        "introduction": "One of the most important questions concerning the Active Galactic Nuclei (AGNs) phenomenon is the connection between starburst and nuclear activity. The role of starbursts in AGNs has been extensively discussed in the past and several theoretical scenarios have been proposed, from those in which the two phenomena are indirectly connected because both are triggered by and live on gas fuelling (e.g., mergers, interactions between galaxies, bars, etc.) to those in which the starbursts are directly bound with the AGNs as possible sources of fuelling for the nuclear black holes. The study of the extended circum- and extra-nuclear regions of ionized gas in active galaxies can shed more light on the interaction between the active nucleus and its environment, because the gas is a convenient tracer of massive star-formation (HII regions), ionization by non thermal processes (e.g., power-law continua), and/or shocks. Evidences of starburst features have been detected in numerous Seyfert 2 \\citep{cid}, but rarely in Seyfert 1 galaxies \\citep{rodriguez}. On the contrary, if the AGNs Unified Model \\citep{antonucci} is valid and a connection exists between AGN and starburst, one would expect to observe a similar distribution of circumnuclear star-forming regions among the different types of AGNs. In this poster we present preliminary results of an investigation of the nuclear and circumnuclear regions of 6 nearby Seyfert galaxies. Narrow-band images isolating the emission lines of H$\\alpha$+ [NII]$\\lambda$6548,6583 and [OIII]$\\lambda$5007 were used to search for spatially extended circumnuclear emission regions. These observations were followed up with long-slit, medium resolution spectrophotometry to probe further the ionization state of the gas. The studied galaxies involve all Seyfert types (three Seyfert 2, one Seyfert 1, one Seyfert 1.5 and one Seyfert 1.8) in order to check a possible correlation between the presence of starburst and the Seyfert type.  \\begin{figure} \\centering \\resizebox{10.0cm}{!}{\\includegraphics{fig1a_bw.eps}} \\vskip 20pt \\resizebox{10.0cm}{!}{\\includegraphics{fig1b_bw.eps}} \\caption{Upper panel: The $H\\alpha$ (left) and $[OIII]\\lambda 5007$ (right) emission-line images of ESO018-G09. Bottom panel: Section of the nuclear spectrum of ESO018-G09 (Seyfert 2), where we see the absorption lines of young stars.} \\label{Fig1} \\end{figure}   \\begin{figure} \\centering \\resizebox{10.0cm}{!}{\\includegraphics{fig2a_bw.eps}} \\vskip 20pt \\resizebox{10.0cm}{!}{\\includegraphics{fig2b_bw.eps}} \\caption{Upper panel: The $H\\alpha$ (left) and $[OIII]\\lambda 5007$ (right) emission-line images of ESO362-G18 (Seyfert 1.8). Bottom panel: The nuclear spectrum of ESO362-G18, where the absorption lines of young stars are detectable.} \\label{fit1} \\end{figure}    \\section {The observations}  The photometric observations were carried out in October 2001 and January 2002 at the Siding Spring Observatory (Australia) with 2.3m telescope. A 1kx1k CCD camera (pixel size = 24 $\\mu m$, scale = 0.6\\arcsec/px) was used in combination with H$\\alpha$+ [NII]$\\lambda$6548,6583 and [OIII]$\\lambda$5007 filters (FWHM $\\sim 75$ \\AA). Long-slit spectra of 4 out of 6 galaxies (ESO018-G09, ESO362-G18, ESO377-G24, ESO428-G14) were obtained at the same telescope with the Dual Beam Spectrograph (pixel size = 15 $\\mu m$, scale = 0.89\\arcsec/px), which allowed to obtain simultaneously the blue (3640-5550 \\AA) and the red (5540-7450 \\AA) parts of the optical wavelength range at medium resolution ($\\sim 2$ \\AA). The spectra of the other 2 galaxies were obtained in December 1990 and March 1991 at ESO-La Silla (Chile) at the 1.52m telescope with the B\\&C Spectrograph (pixel size = 15 $\\mu m$, scale = 0.82\\arcsec/px) at lower resolution. The covered spectral range was 4620-7470 \\AA\\ for NGC 3081 and 3730-6780 \\AA\\ for NGC 1365. The data were reduced and analyzed with the IRAF software. All the observations were bias subtracted, flat field corrected, cleaned from cosmic rays and sky subtracted. The emission-line images were calibrated in flux using a planetary nebula observed in the same runs with the same instrumental configuration. The two-dimensional spectra were calibrated in wavelength using a Ne-Ar lamp for SSO spectra and He-Ar lamp for ESO spectra, then calibrated in flux using a spectrophotometric standard star. For each galaxy we extracted spectra of the nucleus and several circum- and/or extra-nuclear regions, whose sizes were determined by means of both the H$\\alpha$ line profile along the slit and the extension of the emission in the H$\\alpha$ images. We identified the stellar absorption lines and bands in each spectrum. Then, the template of a S0 galaxy was subtracted to remove the underlying stellar population contribution. After this step, we identified the emission lines in the spectra of each object: their profiles were decomposed into multiple Gaussians using the IRAF task FITPROF. In particular, we measured the redshift, width and flux of each identified line. Then, the line fluxes were corrected for internal reddening using the observed Balmer decrement H$\\alpha$/H$\\beta$ chosen equal to 3.1 for the active nuclei and 2.85 for the starburst regions \\citep{osterbrock}, and applying the extinction curve by \\citet{cardelli}. ",
        "conclusions": ""
    },
    "0410/astro-ph0410622_arXiv.txt": {
        "abstract": "Using an artificial \\ion{H}{1} Ly$\\alpha$ spectrum we simulate the corresponding \\ion{He}{2} forest with fixed values of $\\eta$ and a Doppler parameter consisting of a thermal and a turbulent part. In addition metal lines with line strengths and line density as expected in the case of HS~1700+6416 are superimposed. FUSE-like noise is added. The analysis of the simulated spectra in terms of Doppler profiles recovers the input $\\eta = N($\\ion{He}{2}$)/N($\\ion{H}{1}$)$ with a scatter by a factor of 10. About 10\\,\\% have significantly lower $\\eta$-values for various reasons. The majority of the extremely high $\\eta$-values (up to 1000) are evidently caused by metal lines. We conclude that part of the scatter in $\\eta$ of previous analyses of the \\ion{He}{2} forest in HE~2347-4342 can be regarded as an artifact. However, we confirm that the correlation between small column densities (in voids) and high $\\eta$-values is not a methodical artifact and appears to be a true phenomenon. ",
        "introduction": "Using the column density ratio $\\eta = N($\\ion{He}{2}$)/N($\\ion{H}{1}$)$ the sources and fluctuations of the intergalactic ionizing continuum can be examined observing the \\ion{He}{2} Ly$\\alpha$ forest. Following theoretical predictions $\\eta$ is expected to be in the range of $\\sim 50 - 100$ \\citep{HM96, FGS98} assuming the diffuse background radiation of quasars. Recent analyses of the resolved \\ion{He}{2} Ly$\\alpha$ forest towards the QSO HE~2347-4342 reveal values of $\\eta$ ranging from 1 to $ >1000$ \\citep{krissetal01, shulletal04, zhengetal04}. Therefore, the authors suggest that the intergalactical UV background radiation is strongly variable on very small scales requiring the dominance of local sources. However, the scatter can only be explained partly by the different spectral indices of QSOs \\citep{telferetal02}, since the observed scales are much smaller than the typical distances between AGN. Another finding is that absorbers in \\ion{H}{1} voids show higher $\\eta$-values, a phenomenon that is not yet understood.  On the basis of simulated spectra we investigate, whether the typical quality of the present data is sufficient to recover a constant $\\eta$, and which effects are produced artificially by the analysis technique. Furthermore, we examine the impact of additional metal line absorption on the results. ",
        "conclusions": ""
    },
    "0410/astro-ph0410308_arXiv.txt": {
        "abstract": "\\vspace*{-0.4cm} We have completed a survey of 30 Virgo cluster galaxies in the \\ha\\ emission-line using Fabry-Perot interferometry. The goal of the survey is to obtain a high angular resolution sample of velocity fields of spirals and to study the environmental effects on their kinematics and dynamics. \\vspace*{-0.4cm} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410636_arXiv.txt": {
        "abstract": "Using a set of compilations of measurements for extragalactic radio sources we construct all-sky maps of the Faraday Rotation produced by the Galactic magnetic field. In order to generate the maps we treat the radio source positions as a kind of \"mask\" and construct combinations of spherical harmonic modes that are orthogonal on the masked sky. As long as relatively small multipoles are used the resulting maps are quite stable to changes in selection criteria for the sources, and show clearly the structure of the local Galactic magnetic field. We also suggest the use of these maps as templates for CMB foreground analysis, illustrating the idea with a cross-correlation analysis between the Wilkinson Microwave Anisotropy Probe (WMAP) data and our maps. We find a significant cross-correlation, indicating the presence of significant residual contamination. ",
        "introduction": "The origin of large-scale magnetic fields observed on galactic and cluster scales is unknown. The magnetic fields, with observed strengths of $\\sim\\!10^{-6}$G, could be the consequence of an amplification of a tiny seed ($\\la\\!10^{-20}$G) by a dynamo mechanism. Alternatively, the compression of a primordial seed ($\\sim\\!10^{-9}$G) by protogalactic collapse could lead to the fields we see today. Both scenarios require an initial primordial field. Furthermore, the two mechanisms need to explain the high redshift magnetic fields observed in galaxies \\cite{kpz92} and damped Lyman-$\\alpha$ clouds \\cite{wlo92}. Magnetic fields play a crucial role in star formation as well as possibly playing an active role in galaxy formation as a whole \\cite{w78,w02}. Thus, one of the most significant tasks in cosmology is unravelling the mystery behind magnetic fields; from the primordial field to our own Galactic field.  The cosmic microwave background (CMB) provides us with the most distant and extensive probe of the early universe. A primordial magnetic field will leave an imprint in this radiation. Various methods have been developed that seek these signatures. Barrow, Ferreira and Silk (1997) use the anisotropic expansion caused by the presence of a homogeneous primordial field to place limits on its size from large-angle CMB measurements. Others have sought correlation between different scales in the temperature anisotropies \\cite{cmkr04,ncov04} or computed the effects the field has on the polarisation-temperature cross-correlation \\cite{sf97,l04}. The existence of a magnetic field at last-scattering also leads to a possible measurable Faraday rotation of the polarised CMB light \\cite{kl96}.  At the other end of the scale, investigation of our own Galaxy's magnetic field has led to the development of a number of different techniques; Han (2004) provides a concise review of the subject. Techniques include observing Zeeman splitting, polarised starlight, synchrotron radiation, Faraday rotated light and polarised dust emission. However, even with all these methods, there are still outstanding problems. This has led to a lack of consensus on key issues: from the number of spiral arms; how the arms are connected; to the direction the field takes along the arms \\cite{v97,h04}.  To fully understand magnetic fields, we need a coherent picture throughout different epochs. Theories and models can then be tested against this observational picture.  A rotation measure (RM) map of the full sky has the potential to fulfil such a goal. RM values probe the integral of the magnetic field from the radiation source to the observer. Obviously, the information encoded in such a map depends on the location of the radiation that is rotated. Faraday rotated polarised CMB radiation will offer both a picture of the primordial field (at the surface of last scattering) and that of our Galaxy. The recent detection of CMB polarisation by DASI \\cite{klpc02} and confirmation of this via the Wilkinson Microwave Anisotropy Probe (WMAP; Kogut et al. 2003), have opened up a new avenue in CMB research. Future results from WMAP and Planck satellites will offer polarised data covering the full-sky over a range of frequencies (seven in the case of Planck). By forming a RM map with the data and looking at differing scales, we should be able to untangle the information in the data; on the largest scales the local magnetic field can be studied, whereas the primordial field can be studied on the smaller scales.  On the other hand, one of the main tools for probing local magnetic fields (such as our Galaxy's) involves utilising RM values from extragalactic sources. Catalogues containing RM values of extragalactic sources have been used to map the Galactic field (eg. Frick et al. 2001). Combining this data with rotation measures from pulsars that are located within our Galaxy, it is conceivable that a 3-dimensional image of the Galactic magnetic field can be built. However, current RM catalogues are both sparsely populated and unevenly sampled. Thus, astute methods are required to produce a RM map with the data at hand.  So, what do we intend to do? We attempt to map the RM values as a function of angular position giving ${\\cal R}(\\Omega)$, where we use ${\\cal R}$ to denote the Faraday rotation measure and $\\Omega$ for the angular position. Catalogues containing RM values of extragalactic sources are used to construct the function. The observed spatial distribution of the RM values can be expanded over a set of orthogonal basis functions. For analysis of data distributed on the sky, expansion over spherical harmonics seems natural \\begin{equation} {\\cal R}(\\Omega)=\\sum _{l=1}^{\\infty }\\sum _ {m=-l}^{m=+l}a_{l,m}Y_{l,m}(\\Omega ), \\end{equation} where the $a_{l,m}$ are the spherical harmonic coefficients and the $Y_{l,m}$ are the spherical harmonics. The properties of the spherical harmonics are well understood and the calculation of the $a_{l,m}$ will allow us to utilise routines within the HEALPix\\footnote{http://www.eso.org/science/healpix/} package \\cite{healpix} for visualisation purposes and further analysis. However, a non-uniform distribution of data points compromises spherical harmonic analysis due to the loss of orthogonality \\cite{g94}. It is more fruitful to analyse a system using orthogonal functions: the statistical properties of the coefficients are simplified. If nonorthogonal functions are used, the properties of the system and the basis are confused. Therefore, we would like to construct an orthonormal basis with functions closely related to those of the spherical harmonics. The spherical harmonic coefficients can then be obtained from the resultant coefficients of the orthonormal basis.  Spherical harmonic analysis of extragalactic sources has been previously performed by Seymour (1966,1984). However, the analysis was carried out using a different form of orthogonalisation and only on a set of 65 sources.  The RM map resulting from our method will be a useful tool for probing Galactic magnetic structure. The map will also be a valuable point of reference when investigating mechanisms that involve the Galactic magnetic field. For example, CMB foregrounds (synchrotron, dust, free-free emission) are correlated with rotation measures \\cite{dc04}. There is evidence of another foreground component \\cite{dtdg04}, labelled foreground X, that is spatially correlated with $100\\mu$m dust emission. Spinning dust grains \\cite{dl98} are the most popular candidates for causing this anomalous emission. A RM map can provide insight into the role of these grains as they may align with the local magnetic field. Foreground removal will be particularly challenging in CMB polarisation studies as foregrounds are more dominant than in the temperature anisotropies. Also, single frequency polarisation measurements will not be able to remove the effects of Faraday rotation through the Galactic magnetic field. Thus, the extent to which the results have been effected by the $E$-mode signal rotating into the $B$-mode signal (and vice versa) is unknown.  The layout of the paper is as follows. In the next section we describe the three rotation measure catalogues used in our analysis. In describing the data we clarify the meaning of extragalactic rotation measures. In Section \\ref{sec:basis} we illustrate a method to generate orthonormal basis functions for each catalogue. From the coefficients of the new basis, the spherical harmonic coefficients are calculated. In Section \\ref{sec:results} we present the resulting RM maps and discuss the observed features. In Section \\ref{sec:correlations} we give a brief application of the maps. Correlations are sought between the RM maps and cleaned CMB-only maps. The conclusions are presented and discussed in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we have presented a new method to generate all-sky RM maps from uneven and sparsely populated data samples. The method calculates a set of functions orthonormal to the data set. With these basis functions, the spherical harmonic coefficients are calculated and converted into sky maps using the HEALPix package. The method was applied to three catalogues; S81, B88 and F01 catalogues. Maps from each catalogue showed evidence of the magnetic field in our local Orion spur, the North-South asymmetry attributed to radio Loop I and a maxima/minima pair close to the Galactic centre that possibly corresponds to the magnetic field of two inner Galactic arms. A RM map constructed from the S81 catalogue also had a prominent maxima at $l^{\\mathrm{II}}\\!\\sim\\!50^{\\mathrm{o}}$ in the northern hemisphere.  In Section \\ref{sec:correlations}, we showed the benefits a RM map has to CMB foreground analysis. Phase correlations were sought between RM maps and those of CMB-only maps derived from the WMAP data. For both the WMAP team's ILC map and the Wiener-filtered map of TOH, phases corresponding to $l$=11 were found to be highly correlated. Naselsky, Doroshkevich \\& Verkhodanov (2003) found the same scale to display phase correlations when carrying out a similar analysis on the ILC map and foreground maps. Their detection of correlations at the same scale as our analysis, reaffirms that the RM catalogues are valuable independent tracers of CMB foregrounds \\cite{dc04}.  Modelling foregrounds will play a crucial role in CMB polarisation studies. Foreground contamination is expected to be more severe than in the temperature measurements \\cite{k99}. Consequently, superior templates for the individual foreground components are required. RM maps will help trace these components. Besides this, extrapolation of low frequency measurement of synchrotron polarisation to CMB frequencies has been shown to be complicated by Faraday rotation \\cite{dtok03}. Again, this underlines the importance of developing templates of the Faraday rotation of the Galactic sky. Efforts to map the RM sky will be greatly enhanced by increased source catalogues for both extragalactic sources and pulsars within our Galaxy. This may enable the formation of a 3-dimensional image of the Galactic magnetic field. Furthermore, attempts to map the RM sky will be enhanced by upcoming satellite CMB polarisation experiments which present unprecedented sky-coverage and resolution."
    },
    "0410/astro-ph0410293_arXiv.txt": {
        "abstract": "We explore the X-ray properties of a subset of the optically selected SDSS cluster sample of Goto et al. (2002), by analysing seven public XMM-{\\em Newton} pointings, with exposure times ranging from $\\sim$ 4 to 46 ksec. There are in total 17 SDSS clusters out of which only eight are detected at X-ray wavelengths with $\\rm f_{0.5-2keV}\\magcir 1.2 \\times 10^{-14}$ ergs cm$^{-2}$ s$^{-1}$. For the remaining 9 SDSS clusters we estimate their 3$\\sigma$ luminosity upper limits (corresponding to $L_x\\mincir 5\\times 10^{42}$ ergs/sec in the 0.5-2 keV band). This relatively low luminosity suggests that if real structures, these galaxy aggregations correspond to poor groups of galaxies. Using the SDSS photometric catalogue we also derive the cluster optical $r$-band luminosities. The resulting scaling relations ($L_{\\rm opt}-L_x$, $L_{\\rm opt}-T_x$) are consistent with those of other recent studies. ",
        "introduction": "The intense recent interest on galaxy cluster studies is based on the tight cosmological constraints that they can provide, not only from their large-scale distribution, baryon fraction, clustering and dynamics, but lately also from their internal dynamics, its evolution and non-linear physical processes, that give rise in a variety of phenomena related to their formation processes (eg. Peebles et al. 1989;  Borgani \\& Guzzo 2001; Rosati, Borgani \\& Norman 2002). However, in order for such studies to provide the tight constraints on cosmological parameters, that we have recently seen from the amazing CMB experiments (cf. {\\sc Boomerang, Arceops, Wmap} etc) it is essential to understand and eliminate all possible systematic biases that enter in the construction and identification procedures of cluster samples (eg. Nichol 2002).  To this end a large effort has recently been put forward to construct large, bias-free, samples of clusters of galaxies spanning a wide redshift range (eg. Postman et al. 1996; Olsen et al. 1999; Gladders \\& Yee 2000; Goto et al. 2002; Bahcall et al. 2003). A striking result emerging from such studies is that the identified clusters strongly depend not only on the detection algorithm but also on the wavelength used. Optical selection methods, based on galaxy overdensities in two-dimensions are biased due to projection effects, although more sophisticated methods, using color information and/or structural cluster features, provide a more uniform selection of clusters (cf. Postman et al 1996). However, even such algorithms may suffer from projection effects and biases towards more evolved galaxy populations. Alternatively, X-ray selected cluster samples are less prone to biases due to the centrally peaked X-ray emitting cluster core which can be seen even at large ($z\\sim 1$) distances (eg. Stocke et al. 1991; Castander et al. 1995; Ebeling et al. 1996a, 1996b; Scharf et al. 1997; Ebeling et al. 2000; B\\\"{o}hringer et al. 2001; Gioia et al. 2001; Rosati, Borgani \\& Norman 2002). Nevertheless even such identification procedure could be biased towards the most virialized (and hence more X-ray luminous) clusters. A novel cluster detection approach has recently been proposed using simultaneous multiwavelength data with the use of Virtual Observatory data (Schuecker et al. 2004). A generic problem, however, arises from our ignorance of the cluster formation and evolution details which further complicates the characterization of clusters found at high redshifts and their use as cosmological probes.  Attempts to understand the systematic biases that enter in cluster selection procedures have been recently presented in the literature (cf. Donahue et al 2002; Basilakos et al 2004 and references therein). This letter further contributes to this issue by investigating the X-ray properties of a subsample of the Goto et al (2002) clusters detected in the {\\sc Sloan Digital Sky Survey} (SDSS) with the Cut and Enhance method (CE), which is based on galaxy colors\\footnote{The CE catalogue is available from T.Goto, but it can also be found in http://www.astro.noa.gr/xray/data/catalogues.html}. We correlate the positions of these CE cluster candidates with the extended X-ray source position, in order to investigate which of the clusters show significant X-ray emission, and thus have a larger probability of being true aggregations of galaxies. Throughout this paper we adopt the concordance cosmological model ($\\Omega_{\\rm m}=1-\\Omega_{\\Lambda}=0.3$), although all relevant quantities are partametrized according to $H_{\\circ}=100 \\;h$ km s$^{-1}$Mpc$^{-1}$, in order to facilitate comparisons with other studies. ",
        "conclusions": "In order to investigate the X-ray emission of the Goto et al. (2002) SDSS clusters, we have analysed seven public XMM-{\\em Newton} pointings with exposure times ranging from 4 to 46 ksecs. Out of the 17 such SDSS clusters, found in the regions covered by these XMM-{\\em Newton} pointings, only eight ($\\sim 47\\%$) are found to have extended X-ray emission with limiting flux $f_x \\sim 1.2 \\times 10^{-14}$ ergs cm$^{-2}$ s$^{-1}$. Note that two of these eight SDSS clusters, which are also the most X-ray luminous ones, were the prime targets of two of the XMM-{\\em Newton} observations analysed. The remaining 9 clusters are probably either very poor and/or dynamically young clusters with weak X-ray emission or the results of projection effects. The derived $L_{\\rm opt}-L_x$ and $L_{\\rm opt}-T_x$ relations, using the clusters with detected X-ray emission, are consistent with those of other recent determinations.  The present work provides a first glimpse of the observations that the future {\\sc Duo} X-ray mission will perform. It will survey a wide angle (6000 sq. deg) in the 0.3-10 keV energy band with the aim of using $\\sim$ 10000 clusters to map the geometry of the Universe and determine with high precision its Dark Energy component. {\\sc Duo} will target areas covered by the SDSS in order to facilitate the cluster redshift determination. Our work helps to determine the  X-ray flux depth required to identify a significant fraction of the optically selected cluster population."
    },
    "0410/astro-ph0410546_arXiv.txt": {
        "abstract": "We present a new general mechanism to generate curvature perturbations after the end of the slow-roll phase of inflation. Our model is based on the simple assumption that the potential driving inflation is characterized by an underlying global symmetry which is slightly broken. ",
        "introduction": "}  One of the most successful predictions of the inflationary theory, the current paradigm for understanding the evolution of the early universe \\cite{guth81}, is the redshifting of quantum fluctuations of the field driving inflation -- the \\textit{inflaton} -- beyond the Hubble radius, leading to an imprint on the background scalar (density) and tensor (gravitational waves) metric perturbations \\cite{lrreview,muk81,hawking82,starobinsky82,guth82,bardeen83, Hu:2004xd,Mukhanov:1990me,Durrer:2004fx,Langlois:2004de,Brandenberger:2003vk} that subsequently seeds structure formation.  For simplicity, most inflation models assume that there is only one scalar field involved in the dynamics of inflation. This is also the case when the mechanism of converting the energy driving inflation into radiation is considered. In this work we point out a qualitatively new effect that might arise if one relaxes the assumption of a single dynamical field. In a multi-field scenario in which the inflationary potential is characterized by a broken symmetry, the quantum fluctuations generated during the inflationary stage represent fluctuations in the initial conditions for the dynamics of the inflaton in the subsequent stage, thus implying that the background dynamics after the slow-roll phase has ended will differ in different regions of the universe. Since the background fields are coupled to the other fields into which they decay, the fluctuations generated during the slow-roll phase will affect the subsequent decay process.  The present work, assuming that the inflaton decay into other fields through the non-perturbative process of \\textit{preheating} \\cite{Kofman:1997yn,Felder:1998vq}, is then aimed to understand whether isocurvature inflaton fluctuations, generated during the slow-roll stage, can lead to perturbations of the background metric through variations of the preheating efficiency. While the generation of curvature perturbations during the stages following the slow-roll phase has already been considered in some works \\cite{Dvali:2003em,Dvali:2003ar,Matarrese:2003tk,Lyth:2001nq,Liddle:1999hq, Wands:2000dp,Malik:2002jb,Gupta:2003jc}, the present work is the first one to show that in a multi-field scenario a global broken symmetry of the potential is sufficient to yield curvature perturbations. Curvature perturbations produced through this mechanism can even represent the main source of perturbations to the background metric if the inflationary potential is such that the mass required to produce quantum fluctuations along the field trajectory is large, so that the latter result exponentially suppressed.  The structure of the present work is the following. In Sec.\\ \\ref{sect:GI} we obtain a general formula for the curvature perturbations generated from an inhomogeneous preheating efficiency related to the quantum fluctuations produced during inflation. Sec.\\ \\ref{sect:Application} presents an application of the general result obtained in Sec.\\ \\ref{sect:GI} to the case of a broken $U(1)$ symmetry. The conclusions are contained in Sec.\\ \\ref{sect:Discussion}. ",
        "conclusions": "}  The analysis of this paper suggests that the production of curvature perturbations due to a broken symmetry of the inflationary potential and the resulting inhomogeneous efficiency of the preheating stage may be a rather common phenomenon. It is, in fact, appropriate to stress that while the magnitude of the curvature perturbations produced through this mechanism will depend on the details chosen for the specific model (potentials, coupling constants, interaction Lagrangians, preheating mechanism, and so forth), the mere fact that the inflationary potential has a minimum characterized by a broken symmetry is sufficient to guarantee the generation of curvature perturbations during the preheating phase. This is because perturbations in a direction orthogonal to the field trajectory yield evolutions of the background that are not just time translations of one other (as would be the case with perturbations along the direction of the trajectory) but that might differ substantially. Such different background evolutions then necessarily lead to different preheating efficiencies, thus resulting in perturbations in the comoving particle number and energy densities.  As it has been shown in Sec.\\ \\ref{sect:Application}, choosing a specific preheating model allows one to quantify the magnitude of the curvature perturbations produced by this mechanism and to assess whether these may or may not represent a dominant component with respect to the adiabatic perturbation produced during the slow-roll phase by fluctuations along the radial direction. In the present context, the choice of the \\textit{instant preheating} model of Felder {\\it et al.}\\ \\cite{Felder:1998vq} has been made because it seems plausible that the preheat field $\\chi$ may be coupled to some other fields. Moreover, the nature of such process allows one to obtain convenient analytic estimates of the comoving number density of particles produced that are not affected by the stochasticity, related to the build up of preheat particles, usually present in the standard preheating models \\cite{Kofman:1997yn}. Nonetheless, it seems important to stress that the main conclusions of this work do not depend on the choice of the preheating process, but only on the fact that $n_{\\chi}$ depends on the background history, which in turn depends on the initial conditions.  Finally, it seems important to note that if the preheat field is coupled to, and decays into, one single field, then the effect of a broken symmetry of the inflationary potential is to convert isocurvature perturbations into adiabatic perturbations \\cite{Gordon:2000hv}. In this sense, the above model resembles in spirit the curvaton model of Lyth and Wands \\cite{Lyth:2001nq}, but does not require the assumption of an external field."
    },
    "0410/astro-ph0410220_arXiv.txt": {
        "abstract": "We present {\\it Hubble Space Telescope\\/} blue spectra at intermediate spectral resolution for the nuclei of 23 nearby disk galaxies.  These objects were selected to have nebular emission in their nuclei, and span a range of emission-line classifications as well as Hubble types. In this paper we focus on the stellar population as revealed by the continuum spectral energy distribution measured within the central $0\\farcs13$ ($\\sim 8$ pc) of these galaxies. The data were modeled with linear combinations of single-age stellar population synthesis models. The large majority ($\\sim 80$\\%) of the surveyed nuclei have spectra whose features are consistent with a predominantly old ($\\gtrsim 5 \\times 10^9$ yr) stellar population. Approximately 25\\% of these nuclei show evidence of a component with age younger than 1 Gyr, with the incidence of these stars related to the nebular classification. Successful model fits imply an average reddening corresponding to $A_V \\approx 0.4$ mag and a stellar metallicity of (1--2.5)$Z_\\odot$. We discuss the implications of these results for the understanding of the star formation history in the environment of quiescent and active supermassive black holes. Our findings reinforce the picture wherein Seyfert nuclei and the majority of low-ionization nuclear emission-line regions (LINERs) are predominantly accretion-powered, and suggest that much of the central star formation in \\ion{H}{2} nuclei is actually circumnuclear. ",
        "introduction": "\\label{sec:Pops_intro}  The centers of galaxies hold special interest as the locations of unusual energetic phenomena. It is now commonly accepted that most galaxies harbor central supermassive black holes (SMBHs) that in the past may have shone as powerful active galactic nuclei (AGNs). The fueling of luminous AGNs requires substantial gas concentrations that might also be favorable sites for star formation, and the formation of a dense star cluster is in fact one means of spawning a SMBH \\citep[e.g.,][]{Ree84}. In the present Universe, many galactic nuclei exhibit emission-line activity at low levels that apparently derives from a variable mix of accretion power and star formation \\citep[e.g.,][]{Ho97}. Emission-line ratios provide a basis for classifying objects into standard categories that, to varying degrees of confidence, are identified with specific power sources. While classification as a Seyfert or \\ion{H}{2} nucleus is routinely associated with nonstellar accretion and O stars as dominant energy sources, respectively, many galaxies are classified as low-ionization nuclear emission-line regions \\citep[LINERs]{Heck80} or LINER/\\ion{H}{2} ``transition nuclei'' \\citep[e.g.,][]{Fil92,Ho93} which are more ambiguous in their interpretation.  Evidence increasingly suggests that LINERs are primarily accretion-powered \\citep[e.g.,][]{Ho03,Fil03}, but the situation is less clear for the transition sources.  The direct detection of young stellar populations in galactic nuclei is an important means of quantifying star formation in these environments and possible connections to emission-line activity. Nearby examples suggest that recent star formation is not unusual. The Milky Way, IC~342, and NGC~4449 all have central young stellar clusters \\citep{Kra95,Boe01,Boe99}, and our neighbors M33 and M31 also have blue nuclei that are most likely consistent with the presence of young stars \\citep{Lau98}. Moreover, imaging surveys with the {\\it Hubble Space Telescope\\/} ({\\it HST\\/}) have shown that a compact star cluster is often present in the center of spiral galaxies \\citep{Phi96,Caro98,Mat99}, in particular in the latest Hubble types \\citep{Boe02}. If proved young, such clusters may help account for the high frequency (41\\%) of \\ion{H}{2} nuclei in nearby galaxies \\citep{Ho97}, and may also have bearing on transition nuclei.  Most existing studies of stellar populations and nebular emission in galaxy centers have been carried out from the ground, and in probing ``nuclei'' are thus limited to studying regions that may actually be several hundreds of parsecs in scale. Considerable structure may exist on smaller scales, and the spatial distribution of star formation within such a region is of particular interest. A number of ultraviolet (UV) {\\it HST\\/} surveys have shown that star formation in galactic nuclei often takes place in circumnuclear rings \\citep[e.g.,][]{Mao96}. If all bulges harbor a SMBH, there are reasons to believe that it may be difficult to form stars in their very central regions. In particular, the tidal field induced by a SMBH may suffice to disrupt molecular clouds before they can gravitationally collapse. Observationally, however, a number of studies have shown that SMBHs and young stellar clusters can coexist on very small scales, at least in some Seyfert~2 and LINER nuclei \\citep{Heck97,Gonz98,Col02} and in the Milky Way \\citep[e.g.,][]{Gen03}.   Probing galaxy centers with the highest possible spatial resolution is thus of great value for understanding what powers ``nuclear'' activity, and how star formation proceeds in these zones. The recent {\\it HST\\/} surveys of \\citet{Pel99} and \\citet{Caro01,Caro02} represent excellent examples of how high-resolution images have already helped in making steps in this direction. The observed visual to near-infrared colors indicate that the vast majority of disk galaxy centers are very red, consistent with intermediate to very old stellar populations that are substantially obscured by interstellar material. Yet, with only broad-band colors it is difficult to disentangle the effects of an old stellar population from those of reddening by dust, super-solar metallicity, or an active nucleus.  For this reason we undertook a survey of nearby nuclei with the Space Telescope Imaging Spectrograph (STIS) on board {\\it HST\\/} (the SUNNS survey), acquiring moderate-resolution blue and red spectra for 23 galaxies to study both their nuclear stellar populations and emission-line kinematics, and to detect in particular the presence of very young stars or SMBHs. In this paper we present a study of the nuclear stellar populations of the SUNNS sample galaxies, based on the fitting of synthetic stellar population models to the observed blue spectra.  After we submitted this paper for publication, we learned of related work by \\citet{Gonz04} which reaches broadly similar conclusions.  The paper is organized as follows. In \\S 2 we show the nuclear spectra for our sample objects, in \\S 3 we describe our analysis method, and in \\S 4 we present our results. Our analysis includes models based on single-age stellar populations (\\S 4.1) or with extended star formation histories (\\S 4.2) and addresses the effect of super-solar metallicities (\\S 4.3). From the model fitting we derive conservative upper limits on the presence of very young stars and constraints on the need for absorption-line dilution by a featureless continuum (\\S 4.4).  We discuss these results in \\S 5 and draw our conclusions in \\S 6. ",
        "conclusions": "We have fitted stellar population templates to interpret the spectra for the inner $\\sim 10$ pc of a sample of galaxies with a variety of Hubble types and nuclear emission classifications. The primary result of our analysis is that the large majority of nuclei have spectral energy distributions consistent with simply old stellar populations. Only one object out of 23 is dominated by a young cluster of stars less than 1 Gyr old, which is also photometrically distinct. When using multi-age population fits to reflect more complex formation histories, we find a distribution of luminosity-weighted mean ages across the sample that is strongly peaked around 10 Gyr. This result persists if we account for dust reddening and super-solar metallicities as contributors to or possible alternative explanations for the quite red observed spectra. For most sample galaxies, only negligible amounts of stars can have formed in the central 10 pc during the second half of the Universe's life.  There is, however, a significant fraction of galactic nuclei that show modest amounts of young stars: evidence for stars younger than 1 Gyr is found in $\\sim25$\\% of the cases. Very young stars, $\\sim1$ Myr, are not required by the spectral template fits for any object, and their possible contribution to the flux at the blue end of our spectra ($\\sim 3100$~\\AA) is limited to less than 6\\% on average. With these upper limits, the production of ionizing photons by very young stars falls short of the requirements for powering the nebular emission in the central 10 pc of the Seyfert nuclei and most of the LINERs in our sample. Furthermore, none of these 23 galactic nuclei contain enough very young stars to explain the nebular luminosity observed over $\\sim 100$ pc scales, as typically observed from the ground. This shortfall of stellar UV photons holds true regardless of internal dust extinction. The presence of a young population ($\\lesssim$1 Gyr) appears, however, correlated with the classification of the emission-line ratios.  In Seyferts and most LINERs our findings are naturally explained, if ionization comes largely from AGN accretion, not stars. For \\ion{H}{2} nuclei, our results imply that much of the central star formation is actually circumnuclear --- it occurs in the range of a few 10 pc to a few 100 pc.  We note in closing that even for nearby galaxies at {\\it HST} resolution, the contamination from stars in the central portion of the surrounding bulge is an important limiting factor. We could not have detected a quantity of young stars as small as found at the center of the Milky Way."
    },
    "0410/astro-ph0410016_arXiv.txt": {
        "abstract": "The impact of $FUSE$ upon the fundamental parameters of OB stars and Wolf-Rayet stars is reviewed.  The stellar wind signatures available in the far-UV provide us with important additional diagnostics of effective temperature. Together with improved non-LTE stellar atmosphere models allowing for line blanketing and stellar winds, this has led to a downward revision in the spectral type-temperature calibration for O stars versus Vacca et al.  (1996) In addition, the Lyman continuum ionizing fluxes from O dwarfs are compared with previous calibrations of Panagia (1973) and Vacca et al.  We also discuss mass-loss rates in OB stars, such that agreement between recent theoretical predictions (Vink et al. 2000, 2001) and observations of O supergiants is possible, solely if winds are clumped in the far-UV and H$\\alpha$ line forming regions, as favoured by line profile comparisons for P\\,{\\sc v} 1118-28 (early to mid O) or  S\\,{\\sc iv} 1062-1073 (late O to early B) in $FUSE$ datasets. In contrast,  B supergiant wind strengths are predicted to be much higher than observations  indicates, especially if their winds are also clumped. Finally,  significant upward revisions in wind velocities of very late WN stars are indicated by N\\,{\\sc ii} 1085 resonance line observations, plus elemental abundances in OB and WR stars are briefly discussed. ",
        "introduction": "Massive stars ($\\geq 10 M_{\\odot}$)  play an important role in the ecology of their host galaxies, which belies their rarity. O stars are generally the dominant source of Lyman continuum photons in galaxies, whilst their fast, dense winds are the principal source of kinetic energy in young bursts of star formation. Wolf-Rayet stars, the chemically evolved descendants of OB stars, also contribute to chemical and kinetic enrichment of galaxies via their very dense stellar winds, plus Lyman continuum photons and ultimately as core-collapse Supernova, which further dynamically and chemically modify their environment.  Although the majority of our understanding about luminous, hot stars originates from ground-based optical spectroscopy, their energy distributions peak in the far- or extreme-UV. The first rocket UV observations revealed the characteristic P~Cygni signatures of mass-loss from massive stars (e.g. Morton 1967). Longward of Ly$\\alpha$, primarily $IUE$ has provided observations of a very large database of OB stars, within the Milky Way, plus Magellanic Cloud supergiants (Walborn et al. 1985), supplemented by $HST$ datasets of fainter Magellanic Cloud stars (Walborn et al. 1995a; Walborn et al. 2000; Evans et al. 2004b) whilst only a few early-type stars have been observed in the far-UV prior to $FUSE$, with $Copernicus$ (Snow \\& Morton 1976) and more recently HUT (Walborn et al. 1995b; Schulte-Ladbeck et al. 1995). Far-UV $FUSE$ atlases have been presented by Walborn et al. (2002) for Magellanic Cloud OB stars, by Pellerin et al. (2002) for Milky Way OB stars and by Willis et al. (2004) for Wolf-Rayet stars. ",
        "conclusions": ""
    },
    "0410/astro-ph0410366_arXiv.txt": {
        "abstract": "DG Leo is a spectroscopic triple system composed by 3 stars of late-A spectral type, one of which was suggested to be a $\\delta$ Scuti star. Seven nights of observations at high spectral and high time resolution at the Observatoire de Haute-Provence with the ELODIE spectrograph were used to obtain the component spectra by applying a Fourier transform spectral disentangling technique. Comparing these with synthetic spectra, the stellar fundamental parameters (effective temperature, surface gravity, projected rotation velocity and chemical composition) are derived. The inner binary consists of two Am components, at least one of which is not yet rotating synchronously at the orbital period though the orbit is a circular one. The distant third component is confirmed to be a $\\delta$ Scuti star with normal chemical composition. ",
        "introduction": "Different reasons can be invoked to illustrate the fact that main-sequence A-type stars occupy a very interesting region of the H-R diagram: \\begin{itemize} \\item[--] the transition from radiative to convective energy transport occurs at effective temperatures between 8500 and 6000 K (from A5 to F5), a range which encompasses the late A-type stars; \\item[--] the classical Cepheid instability strip, when extended downward, intersects the ZAMS \\citep[cf. Fig.~8 in][]{2001A&A...366..178R} at effective temperatures of 8800 K (blue edge) and 7500 K (red edge) (spectral types from A3 to F0); \\item[--] the presence of magnetism is clearly demonstrated from the mid-B to early-F spectral types; \\item[--] metal-line abundance anomalies are also frequently detected among the (non-magnetic) late A-type stars. \\end{itemize} The variety of (atmospheric) phenomena which may affect the stars located in that part of the H-R diagram in one or the other way is extremely rich: these include several possible pulsation driving mechanisms (acting in the $\\delta$ Scuti and SX Phe, $\\gamma$ Dor or roAp variable stars) and different processes of magnetism, diffusion, rotation and convection which are thought to boost or inhibit the presence of chemical peculiarities in the stellar atmospheres (Ap, Am, $\\rho$ Pup and $\\lambda$ Boo stars). Among the pulsators there exist: the $\\delta$ Scuti stars which are main-sequence or giant low-amplitude variable stars pulsating in radial and non-radial acoustic modes (p modes) with typical periods of 0.02--0.25 days; the SX Phe stars which are high-amplitude $\\delta$ Scuti stars of the old disk population \\citep{2000dsrs.conf....3B}; the $\\gamma$ Dor stars which are cooler and pulsate in non-radial gravity modes (g modes) with typical periods of 1--2 days \\citep{2000dsrs.conf....3B}; rapidly oscillating cooler magnetic Ap (roAp) stars that exhibit very high-overtone non-radial acoustic pulsations aligned with the magnetic axis, inclined relatively to the rotation axis (`oblique pulsator' model) and with typical variability periods of 5--16 min \\citep{2000dsrs.conf..287K}. Among the chemically peculiar (CP) stars of interest, we have the magnetic Ap stars that are generally variable owing to surface inhomogeneities coupled to rotation (model known as the `oblique rotator'); the non-magnetic metallic-lined stars such as the classical Am stars which have K-line and metal-line spectral types differing by at least 5 spectral subclasses, and the evolved Am stars of luminosity class IV and III called $\\rho$ Puppis \\citep[also $\\delta$ Delphini stars,][]{2000dsrs.conf..287K}; and the $\\lambda$ Bootis stars which show weak metal lines owing to underabundances of the Fe-peak elements \\citep{2000dsrs.conf..287K}. Some of these phenomena may or may not be mutually exclusive. For example, enhanced metallicity and pulsation can theoretically occur together in the cooler and/or evolved part of the instability strip while there is mutual exclusion for the hotter A-stars of the main--sequence \\citep{2000A&A...360..603T}. The role of multiplicity is an additional aspect that has to be considered, as the majority of the Am stars are binaries with orbital periods between 1 and 10 days \\citep{1996A&A...313..523B,1997A&A...326..655B}. Most of these binaries are tidally locked systems with synchronized orbital and rotational periods. Their projected equatorial velocities are generally below 100 \\kps~\\citep{2000dsrs.conf..287K}. Because of this tidal braking mechanism, the process of diffusion can act more efficiently. In the case of the (magnetic) Ap stars, the braking mechanism is magnetic. Their projected equatorial velocities are usually well below 100 \\kps~\\citep{2000dsrs.conf..287K}. Thus rotation is relevant for explaining the presence of chemical peculiarities. An unknown multiplicity status is further expected to be the cause of some continuum veiling in several $\\lambda$ Bootis stars leading to an underestimation of the abundance of the Fe-peak elements as was recently suggested in \\citet{2003A&A...412..447G} and \\citet{2003A&A...398..697F}. It also appears that about 50\\% of these exhibit pulsations of $\\delta$ Scuti type \\citep{2000dsrs.conf..410W}. Thus complex interactions between all of these processes exist which are not easily disentangled.  For all these reasons the analysis of the chemical composition of multiple systems having at least one pulsating component is an ideal tool to explore in an empirical way the interactions that may or may not exist between pulsation, diffusion, rotation and multiplicity. In the present paper, we are therefore dealing with the chemical analysis of DG Leo (HD~85040, HR~3889, HIP~48218, Kui~44), a known multiple system with a pulsating component that we describe in Sect.~\\ref{sec:dgleo}. Details about the spectroscopic observations we performed are given in Sect.\\ref{sec:observations} while the Fourier-transform technique that was adopted to disentangle the component spectra from the combined ones is described in Sect.~\\ref{sec:korel}. Orbital parameters of both spectroscopic and visual orbits are discussed in Sect.~\\ref{sec:aorb}. Fundamental parameters and chemical composition of all three components are derived in Sect.~\\ref{sec:fundpar} and Sect.~\\ref{sec:abundance analysis} respectively.   \\begin{figure} \\center \\begin{minipage}{8cm} \\center \\includegraphics[width=7cm,angle=0,clip=]{fcc.eps} \\caption{Two hours-averaged cross correlation function observed from JD2452651.6036 to JD2452651.6954 (i.e.: orbital phase~$\\sim$~0.98), when all three components are resolved. Each correlation peak belongs to one component of DG Leo (Aa, Ab or B).} \\label{fig:crossc} \\end{minipage} \\end{figure} ",
        "conclusions": "The Fourier Transform spectral disentangling technique developed by \\citet{1995A&AS..114..393H} was successfully applied to the study of the pulsating triple system DG Leo. It allowed to confirm the previously derived parameters of the spectroscopic orbit by \\citet{1977PASP...89..216F} and to increase the accuracy. Combining astrometric measurements with our new radial-velocity measurements, we were further able to compute for the first time the relative orbit of the visual system AB. However, we stress the fact that speckle and RV measurements at periastron would be invaluable. We would like therefore to be informed about such data, if existing, as these measurements should greatly increase the precision of the present solution and it could further enable the determination of the stellar masses of all 3 components in an independent way.  A detailed chemical analysis was performed on the disentangled spectra of the components {with} the Kurucz LTE model atmospheres combined with the {\\sc synspec} computer codes. All three components appear to have the same fundamental parameters (effective temperature and surface gravity) but DG Leo B was found to have a solar-like chemical composition, while the two components belonging to the close spectroscopic binary show an abundance pattern typical of Am stars. A different level of underabundance for scandium and more strongly for calcium is noted between the two close components (Aa and Ab). These differences probably reflect the sensitivity of the diffusion processes to issues such as rotation, mass loss and evolution. We further noticed that, while the spectroscopic orbit is circularized, the apparent rotation of both A components is {not yet} synchronized with the orbital motion. DG Leo is therefore a very interesting clue for understanding issues such as rotation synchronization and orbital circularization.  Our high--time--resolution spectroscopic data confirm the existence of line--profile variations (LPVs) caused by pulsation in component B with time scales very close to those recently detected by photometric studies \\citep{phot1}. The analysis and the mode identification of these LPVs will be combined with the multicolour photometric observations, with the goal to provide a better understanding of the non-trivial link existing between chemical composition, multiplicity and pulsation."
    },
    "0410/astro-ph0410685_arXiv.txt": {
        "abstract": "We review the current state and future prospect of ultra high energy cosmic ray physics and the relationship between cosmic rays and gamma-ray astrophysics. ",
        "introduction": "Cosmic rays have been known to have a {\\it cosmic} origin since Victor Hess took electroscopes in  balloons above 5000 m in 1912. A few decades later, Pierre Auger showed that cosmic ray primaries reach energies in excess of $10^{15}$  eV with the detection of extensive air-showers in 1938 \\cite{auger38}. Since then, cosmic rays have been observed from energies of $\\sim$  MeV to above $\\sim 10^{20}$ eV.  Figure 1 shows a compilation of direct and indirect (via air showers)   cosmic ray observations unified into a single spectrum. The spectrum is well fit by power-laws with spectral index $\\gamma \\simeq 2.7$ for energies below  $\\sim 10^{15}$  eV  and  $\\gamma \\simeq 3$ for energies above  $\\sim 10^{15}$  eV,  with a varying low energy cutoff due to solar magnetic fields. The origin of cosmic rays is still a mystery soon to become a century old (see, e.g., \\cite{bhatsigl,reviews} for recent reviews).  The leading theory for the origin of cosmic rays below $\\sim 10^{15}$ eV is Fermi acceleration in Galactic supernova remnant shocks, but direct evidence for this model is still lacking. The best direct test of this hypothesis is likely to come from gamma-ray astronomy. If supernovae shocks are sites of cosmic ray acceleration, the emission of gamma-rays from neutral pion decay around the accelerations site should be observed. As announced by the Hess collaboration, the detection of gamma-rays from RX J1713 \\cite{hess1713} and around the Galactic Center \\cite{hessGC} may be the long sought after smoking guns of Galactic cosmic ray acceleration.  Detailed spectral and spatial studies of the recently discovered sites and others yet to be discovered should resolve the mystery of lower energy cosmic rays within the next decades.  Cosmic rays at the highest energies give another opportunity to resolve the mysterious origin of cosmic rays. At energies above $\\sim 10^{19} eV$ magnetic bending of cosmic rays paths are likely to ease and  cosmic ray astronomy should become feasible. These highest energy messengers are likely to bring information of extra-galactic sources of cosmic rays that may or may nor correlate to the highest energy gamma-ray emitters. While cosmic background radiations limit the arrival of gamma-rays above $\\sim 10^{13} eV$ from distant galaxies a similar limitation may only occurs for cosmic rays  above $\\sim 10^{20} eV$.  Combined studies of gamma-rays and cosmic rays at the highest energies available to each type of messenger will reveal the workings of the extremely energetic universe.  \\begin{figure} \\includegraphics[height=.55\\textheight]{crspectrum.eps} \\caption{Spectrum of cosmic rays \\cite{bhatsigl}.} \\end{figure} ",
        "conclusions": "After decades of attempts to discover the origin of ultra-high energy cosmic rays, present results are still inconclusive. The results from past experiments show the need to understand and control systematic effects within each technique and to cross-calibrate the two techniques presently available for UHECR studies (ground arrays and fluorescence). In addition, the lack of sufficient statistics limits the discussion of an excess flux or a drop in flux around the GZK feature. Next generation experiments are gearing up to accumulate the necessary statistics while having a better handle on the systematics. In the following decade, we may see the growth of a new astronomy with ultra-high energy charged particles and finally resolve the almost century old puzzle of the origin of cosmic rays.    \\begin{theacknowledgments} This work was supported in part by the KICP under NSF PHY-0114422 , by  the NSF through grant AST-0071235, and the DOE grant DE-FG0291-ER40606. \\end{theacknowledgments}"
    },
    "0410/astro-ph0410199_arXiv.txt": {
        "abstract": "We have applied the method of Zeeman tomography to analyze the surface magnetic field structures of the white dwarfs PG\\,1015+014 and HE\\,1045$-$0908 from spin-phase resolved flux and circular polarization spectra obtained with FORS1 at the ESO VLT. We find for both objects field topologies that deviate significantly from centred dipoles. For HE\\,1045$-$0908, the frequency distribution of magnetic field strengths is sharply peaked at 16\\,MG for all rotational phases covered by our data but extends to field strengths at least five times this value. In the case of PG\\,1015+014 there are significant contributions to the frequency distribution in the range from 50 to 90\\,MG with the maximum near 70\\,MG. The detailed shape of the frequency distribution is strongly variable with respect to the rotational phase. ",
        "introduction": "The optical spectra of magnetic DA white dwarfs are dominated by broad Balmer absorption patterns which are characteristic for the magnetic field structure in the visible part of the photosphere. Our evolution-strategy based Zeeman tomography code allows the determination of the surface magnetic field geometry from a set of rotation-phase resolved flux and circular polarization spectra by fitting theoretical model spectra from a precomputed database on a grid of temperature $T$, field strength $B$, and the angle $\\psi$ between the field direction and the line of sight. The method's ability to reconstruct field geometries from synthetic data has been demonstrated by Euchner~et~al.\\ (2002). For the present study, the magnetic field geometry has been parametrized using a truncated multipole expansion up to order $l=3$. ",
        "conclusions": ""
    },
    "0410/astro-ph0410553.txt": {
        "abstract": "Interplanetary dust particles (IDPs) contain enigmatic sub-micron components called GEMS (Glass with Embedded Metal and Sulfides).  The compositions and structures of GEMS   indicate that they have been processed by exposure to ionizing radiation but details of the actual irradiation environment(s) have remained elusive.  Here we propose a mechanism and astrophysical site for GEMS formation that explains for the first time the following key properties of GEMS; they are stoichiometrically enriched in oxygen and systematically depleted in S, Mg, Ca and Fe (relative to solar abundances), most have normal (solar) oxygen isotopic compositions, they exhibit a strikingly narrow size distribution (0.1-0.5 $\\mu$m diameter), and some of them contain ``relict'' crystals within their silicate glass matrices.  We show that the compositions, size distribution, and survival of relict crystals are inconsistent with amorphization by particles accelerated by diffusive shock acceleration.  Instead, we propose that GEMS are formed from crystalline grains that condense in stellar outflows from massive stars in OB associations,  are accelerated in encounters with frequent supernova shocks inside the associated superbubble, and are implanted with atoms from the hot gas in the SB interior.  We thus reverse the usual roles of target and projectile.  Rather than being bombarded at rest by energetic ions, grains are accelerated and bombarded by a nearly monovelocity beam of atoms as viewed in their rest frame. Meyer, Drury and Ellison have proposed that galactic cosmic rays originate from ions sputtered from such accelerated dust grains. We suggest that GEMS are surviving members of a population of fast grains that constitute the long-sought source material for galactic cosmic rays.    Thus, representatives of the GCR source material may have been awaiting discovery in cosmic dust labs for the last thirty years. ",
        "introduction": "% %  %\tInterplanetary dust particles (IDPs) contain enigmatic sub-micron components called GEMS (Glass Embedded with Metals and Sulfides). GEMS consist of amorphous silicate glass grains with non-stoichiometric concentrations of oxygen, and depletions of Mg, S, Ca, and Fe, relative to Si, compared to solar system abundances.  They also contain nanometer-scale inclusions of kamacite (Fe-Ni metal) and pyrrhotite ($\\sim$FeS with up to 2\\% Ni). GEMS are found in a strikingly narrow size range (100-500nm); specifically, small GEMS ($<100$nm) are conspicuously absent from IDPs. GEMS are often pseudo-euhedral despite the fact that they are structurally amorphous (Fig. 1). This property suggests that GEMS were originally individual crystalline mineral grains that were amorphized by a large fluence of ionizing radiation before they were incorporated into the solar nebula. Some GEMS contain  internal relict grains that mimic the external shape and orientation of the euhedral structure of the entire grain.  In other words, these GEMS are pseudomorphs (Fig. 1).    \\begin{figure}[h!] \\begin{center} \\rotatebox{0}{\\resizebox{!}{8cm}{\\includegraphics*{f1.eps}}} \\caption{(a) Darkfield micrograph of GEMS with an embedded relict pyrrhotite (FeS) crystal.\u00ca (b) Same GEMS as in (a) with dotted lines highlighting the approximate outlines of the relict pyrrhotite crystal and the GEMS itself.\u00ca Comparison of the lines indicates that the GEMS is a pseudomorph of the relict pyrrhotite crystal.  (See also \\protect\\citet{bradley-dai04} and \\protect\\citet{bradley99}.)} \\label{gemsimage} \\end{center} \\end{figure}  Here we propose a mechanism for the formation of GEMS from shock-accelerated mineral grains in the hot, low-density ISM.  The organization of the paper is as follows. In section \\ref{bombardment} we review the phenomenology of atomic bombardment of submicron grains. In section \\ref{isolated-supernova} we consider and reject a model in which GEMS are produced by encounters with isolated supernova shocks. In section \\ref{shock-acceleration} we review the theory of dust acceleration by astrophysical shocks.   In section \\ref{superbubbles} we propose an astrophysical site --- the hot, low-density cavities blown by OB associations,  called superbubbles --- in which GEMS are synthesized. In section \\ref{connection} we discuss the connection between GEMS and galactic cosmic rays. Finally, in section \\ref{discussion} we conclude, and suggest new measurements that would test our hypothesis.    % % ",
        "conclusions": "\\label{discussion}  \\subsection{Restricted size range of GEMS}  In Fig. \\ref{survival}, we have shown in our model that only grains in the size range $100-500$ nm survive in the SB interior.    In this model, we neglected the possibility of complete penetration of grains by high-velocity atoms, which will further accelerate the destruction of small grains.   In grains that are large compared to the range of bombarding ions, implantation can dominate over sputtering since bombarding atoms stop inside the grain, and sputtering occurs only on entry.  Such grains can grow with time.  In contrast, grains that are small compared to the range of bombarding ions are quickly eroded because bombarding atoms do not stop inside the grain, and sputtering occurs both on entry and on exit, so that sputtering dominates over implantation. Further, since the ions leaving the grain are slower, they will generally be more effective at sputtering (Fig. \\ref{sputtering}), so the effective sputtering rate will be more than twice that of a large grain.  We have confirmed this effect using SRIM.  \\begin{figure}[h!] \\begin{center} %\\resizebox{!}{4cm}{\\includegraphics*{/export/users/domi/aerogels/calor-papers/Fig.s/tileofshotswtext.eps}} \\resizebox{!}{10cm}{\\includegraphics*{f10.eps}} \\caption{For grains that are large compared to the range of bombarding ions (top), implantation can dominate over sputtering since bombarding atoms stop inside the grain, and sputtering occurs only on entry.  Such grains can grow with time.  In contrast, grains that are small compared to the range of bombarding ions (bottom) can be quickly eroded because bombarding atoms do not stop inside the grain, and sputtering occurs both on entry and on exit, so that sputtering dominates over implantation.} \\label{small-sputtering} \\end{center} \\end{figure}      \\subsection{Are GEMS typical IS grains?}  Astronomical observations have shown that the vast majority of interstellar silicates are amorphous \\citep{kemper}.  The fact that GEMS are also amorphous together with the similarity between the 10$\\mu$m feature of IDPs and that of interstellar dust has led to the suggestion that GEMS are typical representatives of the interstellar dust \\citep{bradley94, bradley99}. The fact that some GEMS exhibit isotopic anomalies indicative of a circumstellar origin support this idea. \\citet{keller-messenger}  have recently pointed out, however,  that the bulk chemistry of GEMS is inconsistent with the typical elemental composition of interstellar dust as derived from gas phase depletions\\citep{gas-depletions}.   They propose that GEMS formed by multiple mechanisms in multiple environments (in the solar system, in  presolar environments, and in the interstellar medium, for example). Based on the bulk compositions of GEMS-rich IDPs, they conclude that most GEMS formed in the solar system.     However, the bulk compositions of GEMS and the IDPs that contain them are equally inconsistent with a solar system origin \\citep{schramm, bradley-ireland,bradley94}.  Moreover, the presence of relict { crystals} and the restricted size range of GEMS are not easily understood in terms of a solar system origin.   Our model is also consistent with the observation that while some GEMS have non-solar O isotopic compositions most have normal (solar) compositions (Keller and Messenger, 2004).  The model predicts that most of the O atoms in most GEMS were deposited from the gas phase leading inevitably to a highly averaged (solar) O isotopic composition.  Nevertheless, nothing in the model that we propose requires that GEMS be representative of average interstellar dust. The absence of average interstellar dust  --- amorphous silicates with bulk chemistry consistent with ISM gas phase depletions --- in IDPs is unexplained.   This paper implies that it is appropriate to revisit the topic of abundances of the major solid-forming elements, particularly Si, as derived from gas-phase depletions in the ISM.  We have proposed a mechanism by which GEMS may be formed from shock-accelerated fast dust in superbubbles, and which explains for the first time the basic properties of of GEMS,  their restricted size range, their compositions, and the presence of surviving internal relict { crystals} in some GEMS. Our model is tied intimately to  the acceleration of galactic cosmic rays;   indeed, parameters in our model are determined from direct observations of galactic cosmic rays in the solar system. The agreement between the predicted and observed GEMS rim thicknesses is particularly striking when we consider that this model includes only one unconstrained parameter --- the sputtering correction factor for SRIM. All other parameters are fixed either theoretically (e.g., the size of the SB core) or experimentally (e.g., the maximum cutoff rigidity is fixed by the observed GCR energy spectrum).    We have made a number of assumptions in our model, which we list here:  $\\bullet$  We made the {\\em ad hoc} assumption that sputtering yields are substantially smaller ($\\sim{1\\over6}$) than the values predicted for silicates by \\citet{tielens}. Although the uncertainties in sputtering yields are subject to uncertainties of a factor of two, we require a more substantial suppression of sputtering rates for GEMS.     It is expected that many interstellar grains acquire a thin carbonaceous mantle\\citep{tielens} that should suppress sputtering yields substantially.  Indeed, the non-thermal sputtering yield of graphite due to He bombardment for velocities $> 700$ km sec$^{-1}$ is less than 40\\% of the yield of silicates \\citet[Fig. 11]{tielens}.  If such a carbonaceous layer does indeed protect GEMS from sputtering, it must be thin if the external shape of relict { crystals} is to reflect that of the glassy rims.   If this is the case, it may lead to an instability in sputtering:  reaccelerated grains that slow to velocities near the peak of the sputtering yield may first lose their protective carbonaceous mantle, then are rapidly eroded because of the enhanced sputtering rate of the bare silicate surface. The preponderance of { relict crystalline FeS} ($\\sim$20\\% of GEMS) remains to be explained, as does the  comparative rarity of relict silicates.  We speculate that silicate minerals may sputter more readily than FeS.  $\\bullet$  We assumed that the magnetic field in the SB interior scales, at worst, with the ambient gas density, so is no smaller than 10 nG.   In fact, the measured GCR spectrum strongly suggests that  the magnetic field is orders of magnitude larger than this.  There is a long-standing problem in GCR astrophysics:  isolated SN shocks should not be capable of accelerating particles much beyond 10$^{14}$eV, and yet the GCR spectrum extends continuously, without steps, from 100 MeV through the knee to the ankle at $\\sim 3\\times10^{18}$ eV. The continuity and smoothness of the spectrum strongly suggests that all GCRs over this energy range must be accelerated by the same mechanism. \\citet{bykov-toptygin}  have recently suggested that GCRs below the knee at $10^{15}$ eV have been accelerated by individual SN, but GCRs above the knee, up to the ``ankle'' in the GCR spectrum at $\\sim 3\\times10^{18}$ eV, have been accelerated by multiple SN encounters.  Acceleration through the collaborative effects of multiple SN in an association is a compelling picture, but it requires a magnetic field of 30$\\mu$G in the SB interior --- more than four orders of magnitude larger than the minimum field strength required to confine fast grains to the metal-rich core.  $\\bullet$  We   assumed that the metal-rich SB core near the OB association mixes inefficiently with material evaporated from the cold SB shell, and that cloud poisoning is negligible.  The extent of mixing is not well-constrained observationally.  $\\bullet$  Finally, and crucially, we assumed that the $10^6$K gas in the metal-rich SB core has the same composition as the  IMF-averaged LC02 core-collapse SN ejecta, that is, that the amount of material depleted by grain condensation is negligible.  Dust condenses quickly in rapidly-cooling SN ejecta, as observed in SN 1987A.  However, we assume that even highly refractory elements from dust destroyed in the hot SB interior does not recondense, and that {\\em most}  dust in SN ejecta is destroyed.   This assumption is consistent with our picture, since small ($<100$ nm) grains, which contain most of the dust mass, are readily destroyed in the SB interior.   Our model makes specific predictions for future observations.  First, $^{22}$Ne/$^{20}$Ne should be much larger ($\\sim 5\\times$) than the solar value in GEMS, just as it is in GCRs. We would expect systematic compositional differences between grains originating as pyrrhotite as compared with the more rare grains that originate in other types, because of the presence of residual material from the original grain.  For example, GEMS containing pyrrhotite relict { crystals} should have larger bulk S than those containing forsterite or enstatite. The decay products of  $^{26}$Al and $^{60}$Fe may  be present, but may not be detectable, unless there is substantial inhomogeneity in Mg/Al or  Fe/Ni ratio in the SB interior that would allow the fossil radioactivities to be detected through positive correlations between, e.g., $^{26}$Mg/$^{24}$Mg  and $^{27}$Al/$^{24}$Mg. Similarly, although there is no evidence of $^{59}$Ni in the ``bulk''  GCRs, it is possible that relict $^{59}$Ni could be found in GEMS if there is substantial inhomogeneity in Ni/Co in the SB interior. Our model also predicts that GEMS   exhibit a $^{58}$Fe excess if the GCR $^{58}$Fe excess originates from core-collapse (type II/Ibc) SN ejecta in superbubbles. So far as we know, no study has yet been made of  the  distribution of amorphous rim thicknesses as a function of grain size.  Our model predicts that relict { crystals} will preferentially be found in large ($>200$nm) grains."
    },
    "0410/astro-ph0410276_arXiv.txt": {
        "abstract": "Dynamo activity caused by waves in a rotating magneto-plasma is investigated. In astrophysical environments such as accretion disks and at sufficiently small spatial scales, the Hall effect is likely to play an important role. It is shown that  a combination of the Coriolis force and Hall effect can produce a finite $\\alpha$-effect by generating net helicity  in the small scales. The shear/ion-cyclotron  normal mode of the Hall plasma is the dominant contributor to the dynamo action for short scale motions. ",
        "introduction": "In astrophysical objects, large scale magnetic fields are thought to be generated by helical turbulence (the so-called $\\alpha$-effect) and differential rotation (the $\\Omega$-effect) (see \\citet{Meneguzzi,Brandenburg}). However, note that a large scale magnetic field can in some cases be generated without fully helical turbulence or a net $\\alpha$-effect. Within the mean field approximation, there are physical effects which contribute to the mean electromotive force even when the $\\alpha$ coefficient is zero. A shear turbulent flow \\citep{Urpin}, the $\\omega\\times j$ term \\citep{Radler69, Geppert}, or magnetic instabilities in a stably stratified atmosphere \\citep{Spruit}, are some of the examples reported in the literature.  Helicity is naturally imparted to a rotating fluid by the Coriolis force \\citep{Moffat1,Moffat2,Moffat3,Moffat4}. However, at sufficiently small scales, where the Rossby number \\begin{equation} R_S = \\frac{U_0}{2 L_0 \\Omega} \\label{Rossby} \\end{equation} is larger than unity ($ U_0 $ and $ L_0 $ are characteristic velocities and lengths, and $ \\Omega $ is the rotation rate), the Coriolis force and the resulting induced helicity might become negligibly small. Nonetheless, it is worth noting that in helical turbulent flows, the kinetic helicity develops a direct cascade along with the energy \\citep{Eyink,Gom2004}. Therefore, some small scale turbulent flows can still be helical, even though the source of helicity remains at much larger spatial scales.  The Hall effect introduces a definite handedness or helicity on small scale fluid motions (\\citet{Wardle}, \\citet{Balbus}), since the mirror symmetry in the induction equation is broken (see Eqn~(\\ref{induc}) below). Therefore, one  should expect a net $\\alpha$-effect  in a Hall plasma. The Hall effect becomes relevant whenever the Hall length scale \\begin{equation} \\lambda = \\frac{c}{\\omega_{pi}}\\ \\frac{U_A}{U_0} \\; . \\label{LHall} \\end{equation} is larger than the dissipation scale, a category to which several objects of astrophysical interest belong \\citep{MGM1,MGM3}. Here, $ U_A $ is the characteristic Alfv\\'enic speed, $c$ is the speed of light, and $\\omega_{pi}$ is the ion plasma frequency.  In \\S 2 we write down the Hall MHD equations. The normal modes  sustained in this system are derived and listed in \\S 3. In \\S 4 we briefly summarize the role of normal mode fluctuations on MHD dynamos, and in \\S 5 this concept is extended to Hall MHD. In \\S 6 we show the effect of rotation on the $\\alpha$-effect. The main results of the present work are summarized in \\S 7. ",
        "conclusions": "We have shown that the Hall effect in conjunction with fluid rotation can generate helicity at small scales (i.e. produce small scale helical motions) leading to a net $\\alpha$-effect through the agency of the shear/ion-cyclotron normal mode of the plasma. This finding can be of considerable importance to the existence of large scale dynamo action in a variety of astrophysical objects. Our results are quite consistent with previous results obtained in the study of instabilities in accretion disks. \\citet{Wardle}, \\citet{Balbus}, and \\citet{Sano} showed that the magneto-rotational instability can be either enhanced or quenched by the Hall effect depending on the orientation of $ \\vOmega $ and $ \\vB_0 $, which is just a manifestation of the handedness introduced by the Hall effect. In a future work the detailed investigation of this mode of dynamo action will be carried out through  direct numerical simulations."
    },
    "0410/astro-ph0410089_arXiv.txt": {
        "abstract": "Building upon a previous analysis of P~{\\sc v} wind lines in LMC O stars, we analyze the P~{\\sc v} wind lines in a sample of Galactic O stars which have empirical mass loss rates determined from either their radio fluxes or H$\\alpha$\\/ profiles.  Since the wind analysis provides a measure of $\\dot{M} q$ where $q$ is the ionization fraction of the ion, we determine $q($P~{\\sc v}$)$ observationally.  In spite of model predictions that $q \\sim 1$ for mid-O stars, we find $q($P~{\\sc v}$) \\leq 0.15$ throughout the O stars.  We discuss the origin of this discrepancy. ",
        "introduction": "$}  Three approaches are normally used to determine stellar mass loss rates, $\\dot{M}$s.  All assume that the wind is homogeneous and spherically symmetric (SS) with a single, monotonic velocity law, and all {\\em should} agree.  The three approaches are:\\\\ \\noindent {\\it 1. Continuum excess from free-free emission.}  This samples the outer wind (the exact radius depends on wavelength), where it becomes optically thick to free-free emission.  It is only detectable for massive winds in nearby stars.  The radio wavelengths are considered ``cleanest'', because, in contrast to the IR/FIR, massive winds become optically thick at large radii ($\\geq 10 R_*$), where $v = v_\\infty$, a constant. This makes $\\rho_{wind}\\sim \\dot{M}/(r^2 v_\\infty)$, independent of $v(r)$. Furthermore, no photospheric correction is needed.  However, the radio flux can be non-thermal, so observations at multiple wavelengths are required to determine the spectral index of the emission.\\\\ \\noindent {\\it 2. H$\\alpha$\\/ emission.}  This samples the inner wind and is easily observed.  For massive winds, H$\\alpha$\\/ emission is related to $\\dot{M}$.  However, the exact form of the observed H$\\alpha$\\/ profile depends upon the $N=3$ departure coefficient for H in the wind and this, in turn, depends upon: the photospheric radiation field; the diffuse radiation field of the wind; and the wind velocity law in the acceleration region.  The shape of the ``photospheric'' H$\\alpha$\\/ profile is also required, and the observed $W_\\lambda (H\\alpha)$ can be strongly variable. Nevertheless, relatively sophisticated models for H$\\alpha$\\/ formation exist (e.g., Repolust et al.\\ 2004), and can provide reasonable agreement between available radio and H$\\alpha$\\/ $\\dot{M}$s when H$\\alpha$ emission is strong.  \\noindent {\\it 3. UV resonance lines.}  These sample the entire wind. Their shapes are determined by the radial optical depth of the wind, $\\tau_{rad} \\sim \\dot{M} q_i A_E$, where $A_E$ and $q_i$ are the abundance of element $E$ and its ionization fraction for stage $i$.  However observations of a dominant ion ($q_i \\sim1$), of known abundance are required to estimate $\\dot{M}$\\/ directly, and the wind lines of abundant, dominant ions are saturated in winds massive enough to be detected in the radio or to have reliable H$\\alpha$\\/ $\\dot{M}$s. ",
        "conclusions": "\\begin{itemize} \\vspace{-0.1in} \\item An erroneous P abundance is not the cause of the small LMC $q(P\\;V)$s. \\vspace{-0.1in} \\item Strong clumping/porosity must be the root of the problem. \\vspace{-0.1in} \\item  Large scale clumping can also strongly affect the radio and H$\\alpha$\\/ $\\dot{M}$s, so these should be re-evaluated. \\vspace{-0.1in} \\item \\underline{The good news:} the different measures of $\\dot{M}$\\/ are sensitive to different aspects of the wind flow, so bringing them all into agreement (together with the X-ray and O~{\\sc vi} wind line observations) will provide powerful constraints on models of how the winds are structured. \\end{itemize}   \\begin{figure}[!ht] \\plotfiddle{figure1_massa.eps}{2.4in}{0}{55}{55}{-170}{-210} \\caption{Mean P~{\\sc v} ionization fractions as a function of temperature for LMC O stars analyzed by Massa et al.\\ (2003).  Open, half filled and filled symbols denote stellar luminosities in the ranges $\\log L/L_{\\odot} > 6.0$, $6.0 \\geq \\log L/L_{\\odot} > 5.6$ and $5.6 \\geq \\log L/L_{\\odot}$, respectively.} \\end{figure}  \\begin{figure}[!ht] \\plotfiddle{figure2_massa.eps}{2.25in}{0}{55}{55}{-170}{-210} \\caption{Mean P~{\\sc v} ionization fractions versus temperature for Galactic O stars with radio or H$\\alpha$\\/ $\\dot{M}$s.  Filled symbols: radio $\\dot{M}$s, open circles: H$\\alpha$\\/ $\\dot{M}$s from Repolust et al.\\ (2004), open triangles: H$\\alpha$\\/ $\\dot{M}$s from Lamers \\& Leitherer (1993).} \\end{figure}"
    },
    "0410/astro-ph0410040_arXiv.txt": {
        "abstract": "We have developed a large CdTe pixel detector with dimensions of 23.7 $\\times$ 13.0 mm${}^{2}$ and a pixel size of 448 $\\times$ 448 $\\mu$m$^{2}$. The detector is based on recent technologies of an uniform CdTe single crystal, a two-dimensional ASIC, and stud bump-bonding to connect pixel electrodes on the CdTe surface to the ASIC. Good spectra are obtained from 1051 pixels out of total 1056 pixels. When we operate the detector at --50 $^{\\circ}$C, the energy resolution is 0.67 keV and 0.99 keV at 14 keV and 60 keV, respectively. Week-long stability of the detector is confirmed at operating temperatures of both --50$^{\\circ}$C and --20 $^{\\circ}$C. The detector also shows high uniformity: the peak positions for all pixels agree to within 0.82 \\%, and the average of the energy resolution is 1.04 keV at a temperature of --50 $^{\\circ}$C. When we normalized the peak area by the total counts detected by each pixel, a variation of 2.1 \\% is obtained. ",
        "introduction": "One of the primary fields of high energy astro-physics in the near future is the hard X-ray universe, where non-thermal processes such as particle acceleration and nucleo-systhesis become dominant. The combination of the two new technologies of hard X-ray focusing mirror optics \\cite{Ref:Yamashita} and hard X-ray imaging spectrometers\\cite{Ref:Takahashi_NeXT} at the focal plane will provide two orders of magnitude improvement in both detection sensitivity and imaging resolution \\cite{Ref:Kunieda,Ref:Fiona_ConX,Ref:NuSTAR}. For example, the NeXT mission\\cite{Ref:Kunieda,Ref:Proposal} proposed in Japan has mirrors with a focal length of 12 m and an angular resolution up to 15 arcsec\\cite{Ref:Ogasaka}. The detector is required to have an energy coverage from 5 keV to 80 keV, with an energy resolution better than 1.0 keV (FWHM) for 60 keV line $\\gamma$-rays\\cite{Ref:Takahashi_SPIE2004}. An aperture size of 20--30 mm in diameter with a spacial resolution of 200 -- 250 $\\mu$m is needed to take advantage of the performance of the mirror. The detector should have a good timing resolution in the range 10 -- 100 $\\mu$s on an event by event basis to reduce the intrinsic background in the detector by the anti-coincidence technique. Good uniformity in both the detection efficiency and the spectroscopic properties is necessary not only to obtain high quality image and spectra, but also to perform proper background subtraction.    Cadmium telluride (CdTe) and cadmium zinc telluride (CZT) are promising devices as the focal plane detector since they have a high detection efficiency comparable to NaI scintilators, a good energy resolution comparable to Ge detectors, and can be operated at room temperature. Thanks to significant progress on technologies of crystal growth, large area detectors based on CdTe and CZT are now available. Recently, Harrison et al. has developed a large area CZT pixel detector with a newly developed low-noise ASIC for the front end\\cite{Ref:Fiona_HEFT}. The detector shows a good energy resolution of $<$ 1 keV. However, the present high pressure Bridgeman method, which is often used to grow CZT crystal, only yields polycrystals and therefore the yield of obtaining large ($>$ 1 cm$^{2}$) and uniform portions of single CZT crystals is low. To control the mobility and carrier lifetime within the whole wafer seems to be difficult\\cite{Ref:GSato,Ref:MSuzuki}. This non-uniformity is the issue of the current CZT pixel detectors. On the other hand,  CdTe crystal grown by the Traveling Heater Method (THM-CdTe) can provide a single crystal as large as 40 mm $\\times$ 40 mm \\cite{Ref:Funaki}. Based on the THM-CdTe wafers,  we have been working on high performance CdTe detectors for both planar and pixel configuration\\cite{Ref:Takahashi_IEEE2002,Ref:KN-NIM2003,Ref:Mitani,Ref:Tanaka}.  In order to improve the current hard X-ray imaging detector by utilizing the recent CdTe technology, we have developed a pixel detector under the collaboration of ISAS and Caltech. In this collaboration, a two-dimensional large area analog ASIC and the read out system is prepared by Caltech. A large size CdTe crystal with pixel electrodes and the bump bonding technologies are prepared by ISAS. In this paper, uniformity of the pixel detector as well as basic performance of the CdTe pixel detector is presented. ",
        "conclusions": "We have developed a large CdTe pixel detector by utilizing recent technologies of an uniform CdTe single crystal, the two-dimensional ASIC, and the In/Au stud bump bonding. When the detector is operated at a temperature of --50 $^{\\circ}$C, the energy resolution (FWHM) for the 14 keV and 60 keV $\\gamma$-rays is 0.67 keV and 0.89 keV, respectively. The voltage scanning measurement shows that 10 \\% of the total event become multi-pixel events, which contributes to the tail structure in thespectra obtained. The detector shows high uniformity as well as the spectral and imaging performances. The location of the peak agrees to within 0.82 \\% after the analog gain is calibrated by pulser data, and the average energy resolution is 1.04 keV. The variation of the 60 keV peak counts is found to be 8.9 \\%, which is inconsistent with the result for the planar CdTe detector.  Although the reason for this variation is still under investigation, if we assume a ``normalized'' area, the resultant variation is 2.1 \\%. In the demonstration of imaging spectroscopy, the detector resolves a 0.5 mm wide line, which is the same as the positional resolution."
    },
    "0410/gr-qc0410007_arXiv.txt": {
        "abstract": "We place direct upper limits on the amplitude of gravitational waves from 28 isolated radio pulsars by a coherent multi-detector analysis of the data collected during the second science run of the LIGO interferometric detectors. These are the first {\\it direct} upper limits for 26 of the 28 pulsars. We use coordinated radio observations for the first time to build radio-guided phase templates for the expected gravitational wave signals. The unprecedented sensitivity of the detectors allow us to set strain upper limits as low as a few times $10^{-24}$. These strain limits translate into limits on the equatorial ellipticities of the pulsars, which are smaller than $10^{-5}$ for the four closest pulsars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410330_arXiv.txt": {
        "abstract": "We present 2D results of simulations of the magnetorotational core collapsed supernova. For the first time we obtain strong explosion for the core collapsed supernova. In 2D approximation we show that amplification of the toroidal magnetic field due to the differential rotation leads to the formation of MHD shockwave, which produces supernova explosion. The amounts of the ejected mass $0.1M_\\odot$ and energy  $\\sim 0.5\\div0.6 \\cdot 10^{51}$ergs can explain the energy output for supernova type II or type Ib/c explosions. The shape of the explosion is qualitatively depends on the initial configuration of the magnetic field, and may form strong ejection neat the equatorial plane, or produce mildly collimated jets. Our simulation show that during the evolution of the magnetic field the magnetorotational instability appears and leads to exponential growth of the magnetic field strength. ",
        "introduction": "The reliable explanation of the mechanism of core collapsed supernova is still open question for the modern astrophysics. We describe the results of 2D numerical simulation of magnetorotational mechanism for core collapsed supernova. The main idea of the magnetorotational mechanism for the core collapsed supernova was suggested by \\citet{bk1970}. After the core collapse central parts of the star rotate differentially. In the presence of initial poloidal magnetic field toroidal component of the magnetic field appears and amplifies with time. When the pressure of the magnetic field becomes comparable with the gas pressure the compression wave arise and moves along steeply decreasing density profile. In a short time this compression wave transforms to the MHD fast shock wave and produces supernova explosion. First 2D simulations of the magnetorotational mechanism were made by \\citet{lw1970} with unrealistically large initial magnetic field. 2D simulations of the similar problem were given in the papers of \\citet{ohnishi} and \\citet{symbalisty}. The magnetorotational processes with very large initial magnetic fields (typical for magnetars) were recently simulated numerically in the papers of \\citet{kotake} and \\citet{yamada}. The 1D simulations of the magnetorotational supernova were presented in the papers by \\citet{bk76}, \\citet{abkp}. In the last paper the simulations were performed for a wide range of the initial values of magnetic field. 2D simulations of the collapse of the magnetized rotating protostellar cloud are presented in the paper by \\citet{abkm}. \\begin{figure} \\centerline{\\includegraphics{moiseenko_fig1.eps},\\includegraphics{moiseenko_fig2.eps}} \\caption{Velocity field for $t=1.06$s after beginning the evolution of the toroidal magnetic field component for the {\\it dipole}-like initial magnetic field (left plot), velocity field for $t=0.2$s after beginning the evolution of the toroidal magnetic field component for the {\\it quadrupole}-like initial magnetic field (right plot).} \\label{veldipquad} \\end{figure} \\begin{figure} \\centerline{\\includegraphics{moiseenko_fig3.eps}\\includegraphics{moiseenko_fig4.eps}} \\caption{Time evolution for of the ejected mass (in relation to $M_\\odot$) and energy for the {\\it quadrupole}-like initial magnetic field, $\\alpha=10^{-6}$.} \\label{masenq} \\end{figure} \\begin{figure} \\centerline{\\includegraphics{moiseenko_fig5.eps}\\includegraphics{moiseenko_fig6.eps}} \\caption{Time evolution of the ejected mass (in relation to  $M_\\odot$) and energy for the {\\it dipole}-like initial magnetic field, $\\alpha=10^{-6}$.} \\label{masend} \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410106_arXiv.txt": {
        "abstract": "{ We present the results of a near-infrared adaptive optics survey with the aim to detect close companions to {\\it Hipparcos} members in the three subgroups of the nearby OB~association Sco~OB2: Upper Scorpius (US), Upper Centaurus Lupus (UCL) and Lower Centaurus Crux (LCC). We have targeted 199 A-type and late B-type stars in the $K_S$ band, and a subset also in the $J$ and $H$ band. We find 151 stellar components other than the target stars. A brightness criterion is used to separate these components into 77 background stars and 74 candidate physical companion stars. Out of these 74 candidate companions, 41 have not been reported before (14 in US; 13 in UCL; 14 in LCC). The angular separation between primaries and observed companion stars ranges from $0.22''$ to $12.4''$. At the mean distance of Sco~OB2 (130~pc) this corresponds to a projected separation of $28.6$~AU to $1612$~AU. Absolute magnitudes are derived for all primaries and observed companions using the parallax and interstellar extinction for each star individually. For each object we derive the mass from $K_S$, assuming an age of 5~Myr for the US subgroup, and 20~Myr for the UCL and LCC subgroups. Companion star masses range from $0.10~{\\rm M}_\\odot$ to $3.0~{\\rm M}_\\odot$. The mass ratio distribution follows $f(q) = q^{-\\Gamma}$ with $\\Gamma=0.33$, which excludes random pairing. No close ($\\rho \\leq 3.75''$) companion stars or background stars are found in the magnitude range $12~{\\rm mag}\\leq K_S \\leq 14~{\\rm mag}$. The lack of stars with these properties cannot be explained by low-number statistics, and may imply a lower limit on the companion mass of $\\sim 0.1~{\\rm M}_\\odot$. Close stellar components with $K_S>14~{\\rm mag}$ are observed. If these components are very low-mass companion stars, a gap in the companion mass distribution might be present. The small number of close low-mass companion stars could support the embryo-ejection formation scenario for brown dwarfs. Our findings are compared with and complementary to visual, spectroscopic, and astrometric data on binarity in Sco~OB2. We find an overall companion star fraction of 0.52 in this association. This is a lower limit since the data from the observations and from literature are hampered by observational biases and selection effects. This paper is the first step toward our goal to derive the primordial binary population in Sco~OB2. ",
        "introduction": "Duplicity and multiplicity properties of newly born stars are among the most important clues to understanding the process of star formation \\citep{blaauw1991}. Observations of star forming regions over the past two decades have revealed two important facts: (1) practically all (70--90\\%) stars form in clusters \\citep[e.g.,][]{LL2003} and, (2) within these clusters most stars are formed in binaries \\citep{mathieu1994}. Consequently, the star formation community has shifted its attention toward understanding the formation of multiple systems --- from binaries to star clusters --- by means of both observations and theory.  The observational progress in studies of very young embedded as well as exposed star clusters is extensively summarized in the review by \\cite{LL2003}. On the theoretical side the numerical simulations of cluster formation have become increasingly sophisticated, covering the very earliest phases of the development of massive dense cores in giant molecular clouds \\citep[e.g.,][]{klessen2000}, the subsequent clustered formation of stars and binaries \\citep[e.g.,][]{bate2003}, as well as the early evolution of the binary population during the phase of gas expulsion from the newly formed cluster \\citep{kroupa2001}. At the same time numerical simulations of older exposed clusters have become more realistic by incorporating detailed stellar and binary evolution effects in N-body simulations \\citep[e.g.,][]{spz2001}. This has led to the creation of a number of research networks that aim at synthesizing the modeling and observing efforts into a single framework which covers all the stages from the formation of a star cluster to its eventual dissolution into the Galactic field, an example of which is the MODEST collaboration \\citep{hut2003,sills2003}.  Such detailed models require stringent observational constraints in the form of a precise characterization of the stellar content of young clusters. Investigations of the stellar population in young clusters have mostly focused on single stars. However, as pointed out by \\cite{larson2001}, single stars only retain their mass from the time of formation whereas binaries retain three additional parameters, their mass ratio, angular momentum and eccentricity. Thus, the properties of the binary population can place much stronger constraints on the physical mechanisms underlying the star and cluster formation process.  Ideally, one would like an accurate description of the `primordial' binary population. This population was defined by \\cite{brown2001} as ``the population of binaries as established just after the formation process itself has finished, i.e., after the stars have stopped accreting gas from their surroundings, but before stellar or dynamical evolution have had a chance to alter the distribution of binary parameters''. This definition is not entirely satisfactory as stellar and dynamical evolution will take place already during the gas-accretion phase.  Here we revise this definition to: {\\em``the population of binaries as established just after the gas has been removed from the forming system, i.e., when the stars can no longer accrete gas from their surroundings''}. This refers to the same point in time, but the interpretation of the `primordial binary population' is somewhat different. The term now refers to the point in time beyond which the freshly formed binary population is affected by {\\em stellar/binary evolution and stellar dynamical effects only}. Interactions with a surrounding gaseous medium no longer take place.  From the point of view of theoretical/numerical models of star cluster formation and evolution the primordial binary population takes on the following meaning. It is the final population predicted by simulations of the formation of binaries and star clusters, and it is the initial population for simulations that follow the evolution of star clusters and take into account the details of stellar dynamics and star and binary evolution. The primordial binary population as defined in this paper can be identified with the initial binary population, defined by \\cite{kroupa1995b} as the binary population at the instant in time when the pre-main-sequence eigenevolution has ceased, and when dynamical evolution of the stellar cluster becomes effective. \\cite{kroupa1995a,kroupa1995c} infers the initial binary population by the so-called inverse dynamical population synthesis technique. This method involves the evolution of simulated stellar clusters forward in time for different initial binary populations, where the simulations are repeated until a satisfactory fit with the present day binary population is found.  Our aim is to obtain a detailed observational characterization of the primordial binary population as a function of stellar mass, binary parameters, and (star forming) environment. The most likely sites where this population can be found are very young (i.e., freshly exposed), low density stellar groupings containing a wide spectrum of stellar masses. The youth of such a stellar grouping implies that stellar evolution will have affected the binary parameters of only a handful of the most massive systems. The low stellar densities guarantee that little dynamical evolution has taken place after the gas has been removed from the forming system. These constraints naturally lead to the study of the local ensemble of OB associations. Star clusters are older and have a higher density than OB associations and are therefore less favorable. For example, in the Hyades and Pleiades, the binary population has significantly changed due to dynamical and stellar evolution \\citep{kroupa1995c}. Note that OB associations may start out as dense clusters \\citep{kroupa2001,kroupaboily2002}. However, they rapidly expel their gas and evolve to low-density systems. This halts any further dynamical evolution of the binary population.  \\cite{brown1999} define OB associations as ``young ($\\la$\\,50~Myr) stellar groupings of low density ($\\lesssim 0.1 {\\rm M}_\\odot {\\rm pc}^{-3}$) ---~such that they are likely to be unbound~--- containing a significant population of B stars.'' Their projected dimensions range from $\\sim$\\,10 to $\\sim$\\,100~pc and their mass spectra cover the mass range from O~stars all the way down to brown dwarfs. For reviews on OB associations we refer to \\cite{blaauw1991} and \\cite{brown1999}.  Thanks to the {\\it Hipparcos} Catalogue (ESA 1997) the stellar content of the nearby OB associations has been established with unprecedented accuracy to a completeness limit of $V\\,\\sim\\,10.5~{\\rm mag}$, or about 1~M$_\\odot$ for the stars in the nearest associations \\citep[][]{dezeeuw1999, hoogerwerf2000}. Beyond this limit the population of low-mass pre-main-sequence stars has been intensively studied in, e.g., the Sco~OB2 association \\citep[][]{preibisch2002,mamajek2002}.  The latter is also the closest and best studied of the OB associations in the solar vicinity and has been the most extensively surveyed for binaries. The association consists of three subgroups: Upper Scorpius (US, near the Ophiuchus star forming region, at a distance of 145\\,pc), Upper Centaurus Lupus (UCL, 140\\,pc) and Lower Centaurus Crux (LCC, 118\\,pc). The ages of the subgroups range from 5 to $\\sim20$\\,Myr and their stellar content has been established from OB stars down to brown dwarfs. \\citep[for details see][]{degeus1989, dezeeuw1999, debruijne1999, hoogerwerf2000, mamajek2002, preibisch2002}. Surveys targeting the binary population of Sco~OB2 include the radial-velocity study by \\cite{levato1987}, the speckle interferometry study by \\cite{koehler2000}, and the adaptive optics study by \\cite{shatsky2002}. Because of their brightness many of the B-star members of Sco OB2 have been included in numerous binary star surveys, in which a variety of techniques have been employed. The literature data on the binary population in Sco~OB2 is discussed by \\cite{brown2001} and reveals that between 30 and 40 per cent of the {\\it Hipparcos} members of Sco~OB2 are known to be binary or multiple systems. However, these data are incomplete and suffer from severe selection effects, which, if not properly understood, will prevent a meaningful interpretation of the multiplicity data for this association in terms of the primordial binary population. The first problem can be addressed by additional multiplicity surveys of Sco~OB2. In this paper we report on our adaptive optics survey of Sco~OB2 which was aimed at surveying all the {\\it Hipparcos} members of spectral type A and late B, using the ADONIS instrument on the 3.6m telescope at ESO, La Silla.  We begin by describing in Sect.~\\ref{sec:theaosurvey} our observations, the data reduction procedures and how stellar components other than the target stars were detected in our images. These components have to be separated into background stars and candidate physical companions. We describe how this was done in Sect.~\\ref{sec: backgroundstars}. The properties of the physical companions are described in Sect.~\\ref{sec: properties}. In Sect.~\\ref{sec: literaturedata} we discuss which of the physical companions are new by comparing our observations to data in the literature and we provide updated statistics of the binary population in Sco~OB2. We summarize this work in Sect.~\\ref{sec:conclusions} and outline the next steps of this study which are aimed at addressing in detail the problem of selection biases associated with multiplicity surveys and subsequently characterizing the primordial binary population. ",
        "conclusions": "\\label{sec:conclusions}      We carried out a near-infrared adaptive optics search for companions around 199 (mainly) A and late-B stars in the nearby OB association Sco~OB2. Our sample is a selection of the {\\it Hipparcos} membership list of Sco~OB2 in \\cite{dezeeuw1999}. We find a total of 151 stellar components around the target stars. We use a simple brightness criterion to separate candidate companion stars ($K_S < 12~{\\rm mag}$) and probable background stars ($K_S > 12~{\\rm mag}$). The validity of this criterion is verified in several ways (\\S~\\ref{sec: backgroundstars}). Of the detected components, 77 are likely background stars, and 74 are candidate companion stars.  The 74 candidate companions occupy the full range in $K_S$ magnitude, down to $K_S=12~{\\rm mag}$. At the distance of Sco~OB2, an M5 main sequence star has a $K_S$ magnitude of approximately $12~{\\rm mag}$. The angular separation between primary and companion ranges from $0.22''$ to $12.4''$. These values correspond to orbital periods of several decades to several thousands of years.  The $J$, $H$, and $K_S$ magnitudes for all observed objects are corrected for distance and interstellar extinction to find the absolute magnitudes $M_J$, $M_H$, and $M_{K_S}$. The subset of the components with multi-color observations are plotted in a color-magnitude diagram. All (except one) candidate companion stars are positioned close to the 5~Myr (for US) and 20~Myr (for UCL and LCC) isochrones. The background stars are positioned far away from the isochrones.  For all observed primaries and companion stars the mass is derived from the absolute magnitude $M_{K_S}$. The mass of the A and late-B primaries ranges between $1.4~{\\rm M}_\\odot$ and $7.7~{\\rm M}_\\odot$ and the mass of their secondaries between $0.1~{\\rm M}_\\odot$ and $3.0~{\\rm M}_\\odot$. Most observed companion stars are less massive than their primaries. Most of the systems with previously undocumented companion stars have a mass ratio smaller than 0.25. The mass ratio distribution for the observed objects peaks around $q=0.1$ and decreases for higher mass ratios. The minimum and maximum value for the mass ratio of the companion stars that we observed are $q=0.022 \\pm 0.008$ and $1.0 \\pm 0.34$. For systems with mass ratio $q \\approx 1$ we cannot say with absolute certainty which is the most massive and therefore the primary star. The mass ratio distribution for binaries with late-B and A type primaries follows the distribution $f(q)=q^{-\\Gamma}$, with $\\Gamma=0.33$. This is similar to the mass ratio distribution observed by \\cite{shatsky2002} for systems with B type primaries. Relatively few systems with low mass ratio are found, excluding random pairing between primaries and companion stars. The uncertainty in the age of the UCL and LCC subgroups does not effect the mass ratio distribution significantly.  A cross-check with visual, astrometric and spectroscopic binaries in literature shows that 41 of the candidate companion stars are new: 14 in US; 13 in UCL and 14 in LCC. The other 33 candidate companions are already documented in literature. We analyze the presently known data on binarity and multiplicity in Sco~OB2, including the 41 new companions that are found in our survey. We conclude that at least 41\\% of all {\\it Hipparcos} member stars of Sco~OB2 are either double or multiple, and we find a companion star fraction $F_{\\rm C} = 0.52$. These values are lower limits since the dataset is affected by observational and selection biases.     Our AO observations are close to completeness in the angular separation range $1'' \\lesssim \\rho \\lesssim 9''$. Next to the 199 target stars we find 50 companion stars in this angular separation range. This corresponds to an $F_{\\rm C}$ of 0.25 in this range of angular separation for A and late-B type stars in Sco~OB2. This value (not corrected for incompleteness) is slightly higher than the values for B type stars in Sco~OB2 \\citep[$F_{\\rm C}=0.20$ per decade of $\\rho$;][]{shatsky2002} and pre-main sequence stars in Sco~OB2 \\citep[$F_{\\rm C}=0.21$ per decade of $\\rho$;][]{koehler2000}.  At small angular separation ($\\rho \\leq 3.75''$) no components other than the target stars are detected for magnitudes $12~{\\rm mag}\\lesssim K_S\\lesssim 14~{\\rm mag}$. This `gap' cannot be explained by observational biases or low-number statistics. A fully populated mass spectrum for the companions (assuming random pairing of binary components from the same underlying IMF) is also incompatible with this gap. This implies a lower limit on the companion masses of $\\sim 0.1$~M$_\\odot$. This is consistent with our finding that the mass ratio distribution points to a deficit of low-mass companions compared to the random pairing case. On the other hand, if we assume that the sources fainter than $K_S\\approx14~{\\rm mag}$ are actually physical companions, a gap may be present in the companion mass distribution. We will carry out follow-up multi-color AO observations to further investigate this issue.  The gap described above might indicate a {\\it brown dwarf desert}, such as observed for solar-type stars in the solar neighbourhood \\citep{duquennoy1991}. The presence of this gap could support the embryo-ejection formation scenario for brown dwarfs \\citep[e.g.][]{reipurth2001,bate2003,kroupa2004}. This is further supported by the detection of 28 candidate free-floating brown dwarf members of Sco~OB2 by \\cite{martin2004}.  The observed $F_{\\rm M}$ decreases with decreasing primary mass. Part of this trend was ascribed by \\cite{brown2001} to observational and selection biases. We lend support to this conclusion with our new AO observations. In particular, the multiplicity fraction for A-type {\\it Hipparcos} members of US is doubled. Nevertheless our knowledge of the present day binary population in Sco~OB2 is still rather fragmentary. None of the multiplicity surveys of Sco~OB2 so far has been complete due to practical and time constraints, but also because of a lack of full knowledge of the membership of the association. In addition, each observational technique used in these surveys has its own biases. For example, visual binaries are only detected if the angular separation between the components is large enough with respect to the luminosity contrast, which means that with this technique one cannot find the very short-period (spectroscopic) binaries. Moreover the observational biases depend on the way in which a particular survey was carried out (including how background stars were weeded out).   We therefore intend to follow up this observational study by a careful investigation of the effect of selection biases on the interpretation of the results. This will be done through a detailed modeling of evolving OB associations using state-of-the-art N-body techniques coupled with a stellar and binary evolution code. The synthetic OB association will subsequently be `observed' by simulating in detail the various binary surveys that have been carried out. This modeling includes simulated photometry, adaptive optics imaging and {\\it Hipparcos} data, as well as synthetic radial velocity surveys. Examples of this type of approach can be found in \\citet{spz2001} (for photometric data) and \\citet{quist2000} (for {\\it Hipparcos} data), where the latter study very clearly demonstrates that an understanding of the observational selection biases much enhances the interpretation of binarity data. Finally, the time dependence of the modeling will also allow us to investigate to what extent stellar/binary evolution and stellar dynamical effects have altered the binary population over the lifetime of the association.  This combination of the data on the binary population in Sco OB2 and a comprehensive modeling of both the association and the observations should result in the most detailed description of the characteristics of the primordial binary population to date."
    },
    "0410/astro-ph0410383_arXiv.txt": {
        "abstract": "Starbursts are a significant component of the present-day universe, and offer unique laboratories for studying the processes that have regulated the formation and evolution of galaxies and the intergalactic medium. The combination of large aperture size, medium-to-high spectral resolution, and access to the feature-rich far-ultraviolet band make FUSE a uniquely valuable tool for studying starbursts. In this paper, I summarize several of FUSE's ``greatest hits'' for starbursts. FUSE observations of the strong interstellar absorption lines show that powerful starbursts drive bulk outflows of the neutral, warm, and coronal phases of the ISM with velocities of several hundred km/s. These are similar to the outflows seen in Lyman Break Galaxies at high redshift. The weakness of OVI emission associated with these flows implies that radiative cooling by coronal gas is not energetically significant. This increases the likelihood that the flows can eventually escape the galaxies and heat and enrich the intergalactic medium. FUSE observations show that local starburst galaxies are quite opaque below the Lyman edge, with typically no more than $\\sim$6\\% of the ionizing photons escaping (even in starbursts with strong galactic winds). This has potentially important implications for the role of star forming galaxies in the early reionization of the universe. FUSE observations of molecular hydrogen in starbursts show very low molecular gas fractions in the translucent ISM (even in starbursts where mm-wave data show that the ISM is primarily molecular). This implies that the molecular gas is all in dense clouds that are completely opaque in the far-UV. The high far-UV intensity in starbursts photodissociates molecular hydrogen in the diffuse ISM. Finally, FUSE observations of chemical abundances in the neutral ISM in dwarf starburst galaxies suggest that the metallicity in the HI (which dominates the baryonic mass budget) may be systematically lower than in the HII regions. This would have important implications for galactic chemical evolution. ",
        "introduction": "Starbursts are intense episodes of star formation that dominate their ``host'' galaxy. Classically, they been defined in terms of their short duration. That is, given the observed star formation rate, the time it would take to consume the presently available reservoir of interstellar gas and/or the time it would take to produce the present-day stellar mass is much less than the age of the universe. An alternative definition is that a starburst has a high intensity: the star formation rate per unit area is very large compared to normal galaxies. As shown by Kennicutt (1998), these two definitions are functionally equivalent: the gas consumption time is a systematically decreasing function of the star formation rate per unit area. Extreme starbursts have gas consumption times of only $\\sim10^8$ years and star formation rates per unit area thousands of times larger than the disk of the Milky Way.  Starbursts are an important source of light, metals, and high-mass star-formation in the local universe (e.g. Heckman 1998). Their cosmological relevance has been highlighted by their many similarities to star forming galaxies at high-redshift. In particular, local UV-bright starbursts appear to be good analogs to the Lyman Break Galaxies (Meurer et al. 1997; Shapley et al. 2003; Heckman et al. 2005) and so can be used as a ``training set'' to provide a thorough understanding of rest-frame UV spectral diagnostics that are critical for studying star-formation in the early universe. Starbursts can contain millions of OB stars, and hence they also offer a unique opportunity to test theories of the evolution of massive stars.  The far ultraviolet spectral window is a key one for understanding starbursts. This is where the intrinsic spectral energy distribution of a starburst stellar population peaks. These hot massive stars have photospheric and wind spectral features in the far-UV that provide key information about the starburst age, metallicity, and initial mass function (Robert et al. 2003; Pellerin, this conference). The far-UV region provides powerful (and in some cases, unique) diagnostics of the physical, chemical, and dynamical properties of the interstellar medium from cold molecular through hot coronal phases. Finally, FUSE is well suited to the problem: its LWRS aperture is an excellent match to the angular sizes of the brightest starbursts, and its spectral resolution allows the interstellar lines to be resolved in almost all cases.  In the following, I will describe what we have learned from FUSE about the interstellar medium in starbursts. ",
        "conclusions": "Starbursts are important components of the present day universe, and wonderful laboratories for studying galaxy evolution, massive stars, and the ISM. The far-UV band is rich with spectral features that provide unique diagnostics of the physical, chemical, and dynamical state of the ISM from its molecular to its coronal phases. The combination of relatively high spectral resolution and a large spectroscopic aperture make FUSE very well suited to the investigation of starbursts. In this review I have summarized some of the highlights of these investigations:  \\begin{itemize} \\item Powerful starbursts drive bulk outflows of the ISM into their galaxy halo at velocities of several hundred km/s. Similar outflows are seen in Lyman Break Galaxies at high-redshift. \\item Radiative cooling/quenching of these outflows by coronal gas is not dynamically significant. This supports the idea that they are the mechanism by which metals were ejected from low mass galaxies and the IGM was metal-enriched. \\item Present-day starbursts are quite opaque to their Lyman continuum radiation. This has interesting implications for the possible reionization of the universe by early starbursts. \\item The translucent ISM in starbursts has a very low molecular gas fraction (most likely due to a high ambient UV intensity in the ISM). \\item The neutral ISM in dwarf starbursts appears to be significantly less metal-enriched than the HII regions. If confirmed, this would have important implications for the chemical evolution of galaxies, since the neutral phase of the ISM dominates the baryonic mass-budget in dwarf galaxies like these. \\end{itemize}  The future prospects are very bright. The All-sky Imaging Survey of the GALEX mission (Martin  et al. 2005) can provide a large sample of starbursts in the local universe that are bright enough for FUSE to obtain very high quality spectra with moderate exposures times. This offers us the opportunity to attack the problems summarized above for a sample large enough to draw statistically robust conclusions across the broad range of fundamental starburst properties (mass, metallicity, star formation rate, etc)."
    },
    "0410/astro-ph0410456_arXiv.txt": {
        "abstract": "We present the deepest yet radio image of the Galactic jet source, SS\\,433, which reveals over two full precession cycles ($> 2 \\times 163$\\,days) of the jet axis.  Systematic and identifiable deviations from the traditional kinematical model for the jets are found: variations in jet speed, lasting for as long as tens of days, are necessary to match the detailed structure of each jet.  It is remarkable that these variations are equal and opposite, matching the two jets simultaneously.  This explains certain features of the correlated redshift residuals found in fits to the kinematic model of SS\\,433 reported in the literature.  Asymmetries in the image caused by light travel time enabled us to measure the jet speeds of particular points to be within a range from $0.24\\,c$ to $0.28\\,c$, consistent with, yet determined independently from, the speeds derived from the famous moving optical emission lines.  Taken together with the angular periodicity of the zigzag/corkscrew structure projected on the plane of the sky (produced by the precession of the jet axis), these measurements determine beyond all reasonable doubt the distance to SS\\,433 to be $5.5 \\pm 0.2$\\,kpc, significantly different from the distance most recently inferred using neutral hydrogen measurements together with the current rotation model for the Galaxy. ",
        "introduction": "\\label{sec:intro}  SS\\,433 is famous as the first known relativistic jet source in the Galaxy.  Red- and blue-shifted optical lines, indicating velocities of $0.26c$, were discovered by \\citet{Mar79a,Mar79b} and interpreted as being from gas accelerated in oppositely-directed jets \\citep{Fab79,Mil79}.  \\citet{Mar84} successfully fitted a kinematic precessing jet model, finding an intrinsic jet speed of $\\approx0.26c$, a precession period of $\\sim 163\\rm\\,days$, a cone opening angle of $\\sim20^\\circ$, and an inclination to the line-of-sight of the precession axis of $\\sim80^\\circ$.  Subsequent radio imaging was consistent with this model \\citep[e.g.,][]{Hje81,Ver87,Ver93}, with the images of \\cite{Hje81} resolving ambiguities in the parameters from optical studies.  While optical spectroscopic data observed over many precession cycles \\citep[e.g.][]{Mar84,Eik01} have confirmed the basic parameter values of what has come to be known as the kinematic model, a detailed analysis of any deviations of the data from this model is hindered by inherent degeneracies in the parameters to be fitted, since the measurement is that of line-of-sight Doppler shifts.  Detailed analyses from quasi-daily milli-arcsec scale monitoring are similarly afflicted by such degeneracies since, for example, a variation in observed proper motion can in principle arise from changes in jet velocity projected on the plane of the sky, or changes in the angle the jet axis makes with our line-of-sight, perhaps caused by a change in precession rate.  Moreover, in milli-arcsec scale monitoring only a limited fraction of the precession period of 162 days has been sampled to date.  However, the extended and spatially resolved jet output observed over several arcseconds, while not time-resolved in the conventional sense, is an historical record of the geometry of the jet ejection over two complete precession periods. ",
        "conclusions": "\\label{sec:conclusions}  We have presented the deepest yet radio image of SS\\,433 on arcsecond scales, from which we have derived its distance from us to be 5.5\\,kpc, independently of any assumptions from optical data.  We have identified deviations from the simple kinematic model in our image, which last over timescales of 10s of days.  These may be formally fitted in either of two extreme scenarios: by small oscillations purely in jet velocity about a mean of $0.26\\,c$ (which are perfectly symmetric in both jets) or by large oscillations purely in the rate of precession about a mean of 162 days (which are perfectly symmetric in both jets).  Additional information from redshift residual data \\citep{Eik01} strongly suggests that variations in velocity and in phase are both occurring.  We speculate that the phase (precession rate) variations might arise because of a varying effective moment of inertia of the nozzle (which in turn might arise because of variations in the mass transfer rate). We suggest that these same changes in mass distribution might cause variation in the inner radius of the accretion disc which may determine the speed of jet bolide ejection \\citep{Mei01}."
    },
    "0410/astro-ph0410726_arXiv.txt": {
        "abstract": " ",
        "introduction": "Red giant stars, both in the field and in globular clusters, present abundance anomalies that can not be explained by standard stellar evolution models. Some of these peculiarities, such as the decline of $^{12}{\\rm C}/^{13}{\\rm C}$, and that of Li and $^{12}{\\rm C}$ surface abundances for stars more luminous than the \\emph{bump}, clearly point towards the existence of extra-mixing processes at play inside the stars, the nature of which remains unclear. Rotation has often been invoked as a possible source for mixing inside Red Giant Branch (RGB) stars (\\cite{SM79},\\cite{CC95},\\cite{DT00}). In this framework, we present the first fully consistent computations of rotating low mass and low metallicity stars from the Zero Age Main Sequence (ZAMS) to the upper RGB. ",
        "conclusions": ""
    },
    "0410/astro-ph0410510_arXiv.txt": {
        "abstract": "\\noindent In the present paper we report on first results of a project in Brussels where we study the effects of stellar dynamics on the evolution of young dense stellar systems using the 3 decades expertise in massive star evolution and our population (number and spectral) synthesis code. We highlight an unconventionally formed object  scenario ({\\it UFO-scenario}) for Wolf Rayet binaries and study the effects of a luminous blue variable-type instability wind mass loss formalism on the formation of intermediate mass black holes. ",
        "introduction": "A population synthesis code calculates the temporal evolution of a population of single stars and of close binaries, in regions where star formation is continuous or in starbursts. Population number synthesis (PNS) predicts the number of stars of a certain type whereas population spectral synthesis (PSS) computes the effects on the integrated spectrum of a population of a certain class of stars. The latter is very useful in the case of young starbursts.  Notice that to make realistic PNS/PSS predictions of massive stars it is essential to use evolutionary tracks calculated with the most up to date wind rates, where we distinguish those during the core hydrogen burning (CHB) phase, the luminous blue variable (LBV) phase, the red supergiant (RSG) phase and the core helium burning (CHeB) phase when the star is classified as a Wolf-Rayet (WR) star. A description of our PNS and PSS code that follows the evolution of young starbursts can be found in Van Bever and Vanbeveren (2000, 2003), and references therein.  The effects of N-body stellar dynamics may be very important in dense stellar systems (Portegies Zwart et al., 2004 and references therein) and we therefore started recently with the implementation of this process in our codes. In the present paper we report first results (more details will be presented by Belkus et al., 2005). In section 2 we describe our model. In section 3 we further discuss the unconventionally formed object scenario ({\\it UFO-scenario}) introduced by Dany Vanbeveren, these proceedings and section 4 illustrates the effect of an LBV-type instability in the most massive stars on the formation of intermediate mass black holes (IMBHs). ",
        "conclusions": ""
    },
    "0410/astro-ph0410732_arXiv.txt": {
        "abstract": "{A sample of spectroscopic binaries and a sample of single planetary systems, both having main-sequence solar-type primary components, are selected in order to compare their eccentricities. The positions of the objects in the ($P.(1-e^2)^{3/2}$, $e$) plane is used to determine parts in the period--eccentricity diagram that are not affected by tidal circularization. The original eccentricities of binaries and planets are derived and compared. They seem to be weakly or not at all correlated with period in both samples, but two major differences are found :  \\noindent (1) The tidal circularization of planetary orbits is almost complete for periods shorter than 5 days, but it is not visible when $P.(1-e^2)^{3/2}$ is longer than this limit. This suggests that the circularization occurs rapidly after the end of the migration process and is probably simultaneous with the end of the formation of the planet. By contrast, we confirm that the circularization of the binary orbits is a process still progressing a long time after the formation of the systems.  \\noindent (2) Beyond the circularization limit, the eccentricities of the orbits of the planets are significantly smaller than those of binary orbits, and this discrepancy cannot be due to a selection effect. Moreover, the eccentricities of binaries with small mass ratios are quite similar to those of all binaries with $q<0.8$. This suggests that the low eccentricities of exoplanet orbits are not a consequence of low-mass secondaries in a universal process.  These remarks are in favor of the idea that binaries and exoplanets are two different classes of object from the point of view of their formation. ",
        "introduction": "It is well known that the orbits of the exoplanets with periods larger than 5 or 6 days have eccentricities significantly larger than those of the giant planets of the solar system. Several mechanisms were proposed to explain this feature, but, up to now, none is fully convincing. It was proposed that eccentric orbits could be a consequence of the dynamic evolution of systems initially involving several planets (Rasio \\& Ford \\cite{RaFo96}, Lin \\& Ida \\cite{LinI97}, Ford et al. \\cite{Ford01}, Papaloizou \\& Terquem \\cite{PapaTer01}, Rice et al. \\cite{Rice03}), but these models fail to produce the frequency of giant planets with semi-major axes smaller than about 1 AU. The giant planets close to their harboring stars are often assumed to be produced by migration within a disk (Ward \\cite{Ward97}, Masset \\& Papaloizou \\cite{MaPa03} and references therein), but this process hardly produces eccentric orbits (Papaloizou et al. \\cite{PaNeMa01}, Thommes \\& Lissauer \\cite{ThoLi03}), although Goldreich \\& Sari (\\cite{GolRei03}) and Woolfson (\\cite{Woolf03}) leave some room for hope. Therefore, it is tempting to consider that the exoplanets are generated by the same process as binary stars (Stepinski \\& Black \\cite{SteBla00}). This implies that giant exoplanets were not formed by gas accretion onto a heavy rocky core, as usually assumed, but by an alternative process. They could come from disk instabilities (Mayer et al. \\cite{Mayer02}, Boss \\cite{Boss02}, \\cite{Boss03}), but inward migration in a disk is then invoked again to explain the short-period orbits; alternatively, planets could be generated by fragmentation of a collapsing protostellar cloud, {\\it via} filament condensation and capture (Oxley \\& Woolfson \\cite{OxWoo04}), or even exactly as stellar components in binary systems (see discussion in Bodenheimer et al. \\cite{BoHuLi} and references therein). However, these models may be efficient in forming massive planets or brown dwarfs, but not planets around 1 Jupiter mass or less.  Note that the binary formation models are not very satisfactory either (see the review by Tohline \\cite{Toh02}). The large eccentricities of binaries are explained by fragmentation of collapsing cores and subsequent interactions between the forming stars (Bate et al. \\cite{BaBoBro02}, Goodwin et al. \\cite{GoWhiWa04}), but, as for exoplanets, the simulations do not provide the high frequency of close systems. Moreover, statistical investigations on main sequence binaries (Halbwachs et al. \\cite{Halb03}, Paper I hereafter) have shown that the close binaries (i.e. with semi-major axes less than a few AU) consist in two populations~: one with large eccentricities and mass ratios less than 0.8 (``non-twins'' hereafter), and one with moderate eccentricities and nearly identical components (``twins''). Additionally, the twins are more frequent among short-period binaries than among the others. At first, these properties were derived from binaries with F7--K primary components, but they are also valid for M-type dwarfs (Marchal et al. \\cite{MaDel03}).  In the present paper, the period--eccentricity diagram is used to compare the exoplanets with the binary stars~: our main purpose is to investigate if the properties of exoplanets may be considered as an extrapolation of the properties of binaries in the range of very low mass ratios. This would indicate whether the formation processes of these objects are similar. In the course of the paper, a few other points are also treated~: (1) the correlation between the eccentricity and the period or the angular momentum, (2) the relation between the eccentricity and the metallicity of planets, (3) the original distributions of eccentricities for planets and for binaries, considering the twins separately. Comparisons between planets and binaries in the period--eccentricity diagram were already presented by Mayor et al. (\\cite{Mayor01}), Mazeh \\& Zucker (\\cite{MaZu01}), and Udry (\\cite{Udry01}), who concluded that planets and binaries are very similar when periods longer than 50 days are considered. However, their samples contained around 30 or 40 planets, and a many others have been discovered since. The question needs to be re-considered.  The interpretation of the period--eccentricity diagram is rather complex, and our investigations are based on the method presented in Sect.~\\ref{method}. Sect.~\\ref{binaries} is dedicated to the binaries; we investigate if, additionally to the twins, other classes of mass ratio have specific distributions of eccentricities. A similar treatment is applied to exoplanets in Sect.~\\ref{planets}. Binaries and exoplanets are compared in Sect.~\\ref{comparison}, in which we pay attention to the difference in the selection effects of both samples. The consequences of our results are discussed in Sect.~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have found some relevant features in comparing the eccentricities of the SB to those of the exoplanets~:  \\begin{itemize} \\item The ($P_{\\rm sr}$ -- $e$) diagram, based on $P_{\\rm sr}$ defined in Eq.~(\\ref{Pcirc}), is a powerful tool for determining the limit between the circularized orbits and the others. The contrast between the two areas in the diagram suggests that the tidal effects are efficient only when the semilatus rectum of the orbit is less than a fixed limit.  \\item The transition from the circularization to the part not affected by tidal effects looks sharper for the exoplanets than for the SB. For the SB, it corresponds to  $P_{\\rm sr}$ between around 5 and 10 days, in agreement with Mathieu \\& Mazeh (\\cite{MaMa88}), Duquennoy et al. (\\cite{DMM92}), Mathieu et al. (\\cite{MaDuLaM92}) and Witte \\& Savonije (\\cite{WiSa02}), who consider that circularization is not restricted to the time of binary formation, but is still progressing during the whole lifetime of the main sequence components. For the exoplanets, the fast transition observed for planets with different ages is consistent with the idea that the tides were efficient only during the formation of the system. The inefficiency of tides for a formed planet is in agreement with the circularization time derived by Zahn (\\cite{Zahn77}), which is a function of $(1+q)/q$. If the planets were brought closer to their host stars by migration, that means that migration occurred when the formation of the planets or of the host stars was not completed. Tidal circularization was then dominated by the tidal bulge on the planet, which was hotter, and therefore larger than it is today.  \\item Before tidal effects had modified them, the eccentricities of planets or binaries were not strongly related to the periods, or to the angular momenta. It is even quite possible that they were not correlated at all with these parameters, since the absence of correlation is not clearly rejected by statistical tests, and also because our assumption that tides did not affect at all the orbits with $P_{\\rm sr} > P_{\\rm cutoff}$ is possibly a bit too simple. Since the planets in the sample are supposed to have migrated, this suggests that migration did not alter the eccentricities significantly; alternatively, it is possible that the eccentricities were modified, but almost independently of the final periods.  \\item The exoplanets have orbits with eccentricities significantly smaller than those of the SB with the same period and with mass ratios larger than 0.8 (the non-twin binaries). A similar feature has already been pointed out for $P<50$~days (Udry \\cite{Udry01}), but neglecting the difference between the distributions of periods of binaries and of planets. Moreover, it is now certain that the low eccentricities of planets are not an effect of the selection of the observed sample. Additionally, it seems that the distributions of the eccentricities are not related to the masses of the companions, neither among the non-twin binaries, nor among the planets. Therefore, this discrepancy is probably not an effect of the low masses of planets in a formation/evolution process common to planets and binaries~: this would imply a process depending on the secondary mass, but only around the transition between stellar and planetary companions. In fact, the SB most similar to the planets are the twins, perhaps because these systems were also interacting with a disk at the time of their formation, as Tokovinin (\\cite{Toko00}) suggested.  \\end{itemize}  Our most relevant conclusion is that the eccentricities of the exoplanet orbits are rather in favor of the hypothesis that exoplanets and binary stars are not the products of the same physical process. After the ``brown dwarf desert'' (Halbwachs et al. \\cite{Halb00}) separating the stellar components from the planets in the distribution of the secondary masses, it is another argument in that sense which was derived from statistical investigations."
    },
    "0410/astro-ph0410218_arXiv.txt": {
        "abstract": "It has long been known that X-ray and optical/UV variability in active X-ray binaries is sometimes correlated.  The expectation has been that this arises from rapid reprocessing of X-rays into longer wavelength photons, and hence that the lags between the bandpasses can be used to reconstruct an echo-map of the binary, revealing the geometry and spatial scale of reprocessing sites.  I will review how this can be done, what can be learned, and what progress has been made with real data.  In addition, an interesting twist to the story has emerged in the form of correlated variability that does not seem associated with reprocessing.  This extends to very short timescales, includes quasi-periodic behavior, and may be associated with optical synchrotron emission from a jet.  Finally, contrary to expectations we have recently discovered that similar correlations are seen even in an ostensibly `quiescent' object allowing similar analyses of reprocessing effects. ",
        "introduction": "X-ray binaries, and more specifically the Low Mass X-ray Binaries (LMXBs) that are the focus of this review, are variable objects.  This is a universal characteristic seen from the radio through optical to X-rays and $\\gamma$-rays.  The X-ray variability is usually dominated by instabilities in the accretion flow.  X-rays irradiate the outer accretion disk and companion star, resulting in reprocessed optical and UV radiation which is expected to be imprinted with the same variability as the X-ray signal is.  An important difference, however, is that the optical and X-rays originate from a volume of significant spatial extent, resulting in light travel time delays between the X-rays and the reprocessed emission.  It is then possible to infer information about the geometry and scale of the reprocessing region from the lags measured between X-ray and optical/UV variability; this technique is known as reverberation or echo-mapping, as the reprocessed light behaves as an echo.  In this work we will review the concepts underlying the echo-mapping technique and the successes achieved in applying it to LMXBs to date. Surprisingly the technique is applicable not only in the most X-ray luminous systems but in at least one {\\em quiescent} black hole as well.  In attempting to apply the technique it has also emerged that not all optical/UV correlations arise in reprocessing at all.  We will briefly consider an alternative kind of correlations that may originate in optical/UV synchrotron emission from a jet. ",
        "conclusions": "Correlated X-ray and optical/UV variability has now been widely, if not commonly, detected in X-ray binaries.  The detections span the lowest luminosities to the highests and suggest a more complex origin than the simple reprocessing originally envisioned.  Reprocessing does appear to be responsible for reprocessed type I X-ray bursts, and for correlated flickering in luminous X-ray binaries.  Typically these studies have shown that the disk appears to dominate the response, although in some cases the lags appear to lag for light travel times alone.  Where a reprocessed spectrum has been obtained it is consistent with predominantly optically thick thermal reprocessing.  The full potential of this technique has yet to be exploited, either through phase-resolved echo-tomography or through kinematically resolved emission line echo-mapping.  On the other hand, correlated variations in X-ray binaries in the low/hard state do not have the characteristics of reprocessed variability, and instead appear to originate in optical synchrotron emission, likely from a jet.  Understanding the variability properties in the context of a jet model remains a topic of ongoing research, but this may prove a very valuable diagnostic of the disk-jet connection.  At the lowest luminosities, there is evidence that emission line variability in V404~Cyg in quiescence is dominated by reprocessing within the disk.  The origin of continuum variability is less clear, and this might either be reprocessing, or synchrotron.  For these systems the low X-ray brightnesses are a challenge that Chandra and XMM are barely adequate for, and full exploitation of echo-mapping in quiescence will likely require a larger throughput mission such as Constellation-X."
    },
    "0410/astro-ph0410442_arXiv.txt": {
        "abstract": "{ Analyzing IUE ultraviolet spectra of $\\beta$ Cep pulsating stars we have noticed that multiperiodic variables have a larger mean metal abundance in the photosphere, [m/H], than monoperiodic ones. We apply statistical tests to verify this dichotomy. We obtain that, with a large probability, the multiperiodic $\\beta$ Cep stars have greater values of [m/H]. This result is consistent with the linear non-adiabatic theory of pulsation of early B-type stars. ",
        "introduction": "The $\\beta$ Cephei pulsators are a well-known and well studied group of early-type pulsating stars. Oscillations in these stars are strictly connected with the metal abundance $Z$, as they are driven by the classical $\\kappa$-mechanism operating in the layer of the metal opacity bump $(T\\approx 2\\cdot 10^5~{\\rm K})$ caused by the lines of Fe-group elements (Moskalik \\& Dziembowski 1992, Dziembowski \\& Pamyatnykh 1993, Kiriakidis et al. 1992 and Gautschy \\& Saio 1993). The size of the instability domain of $\\beta$ Cep stars is very sensitive to the heavy element abundance $Z$, and is smaller for lower values of $Z$ (Dziembowski \\& Pamyatnykh 1993, Pamyatnykh 1999). Therefore, the information about the metal abundance in $\\beta$ Cep stars is of special importance.  In previous works (Daszy\\'nska 2001, Daszy\\'nska et al. 2003) we derived the metallicity parameter in the atmosphere, [m/H], and the mean stellar  parameters for all $\\beta$ Cephei stars monitored by the {\\it International Ultraviolet Explorer}. We showed that the metallicity values obtained are independent of effective temperature and surface gravity. Additionally, they are not correlated with the stellar rotation and any of the pulsational parameters, like a dominant period or amplitudes of the light and the radial velocity variations.  In the scope of her PhD thesis, Daszy\\'nska (2001) suggested that, in general, the multiperiodic $\\beta$ Cep stars have higher values of [m/H] than the monoperiodic ones. We mentioned this result in Daszy\\'nska et al. (2003), and we confirmed it in Niemczura \\& Daszy\\'nska-Daszkiewicz 2004 (Paper I), where we applied a more accurate method of determination of stellar parameters. Using the results of Paper I, we verify here this hypothesis by means of statistical tests.  The paper is composed as follows. In Sect. 2 we remind derived values of [m/H]. Sect. 3 contains the test of statistical hypotheses about equality of two means. In Sect. 4 we discuss the dichotomy of $\\beta$ Cep variables with regard to the metallicity. Conclusions are given in Sect. 5. ",
        "conclusions": "Using statistical tests, we have shown that with a very high probability the multiperiodic $\\beta$ Cephei stars have a higher mean metallicity in their atmospheres than the monoperiodic ones. The statistical  analysis was based on our last determinations of the metallicity parameter [m/H] (Paper I). Getting a normal distribution of [m/H] for the mono- and multiperiodic stars, we applied the test about equality of variances of these two samples, represented by X and Y random variables. The test based on the $F$ statistic allowed clearly to accept the null hypothesis $H_0:~\\sigma^2_Y=\\sigma^2_Y$. Then we could test the hypothesis about the mean values of metallicity for the mono- and multiperiodic variables. We used the $T$ statistic to define the critical regions for a small size of random samples. The null hypothesis $H_0:~\\mu_X=\\mu_Y$ was rejected in favor of $H_1:~\\mu_X\\neq\\mu_Y$, as well as in favor of $H_2:~\\mu_X<\\mu_Y$.  The higher value of the mean metal abundance in the case of multiperiodic $\\beta$ Cep stars is understandable in the framework of the linear non-adiabatic theory, which predicts more unstable frequencies for higher metal abundance $Z$. Therefore, we tried to find a dependence between the range of observed frequencies and the metal abundance. We can not make a statement about any correlation between these two parameters, because the value of the weighted correlation coefficient $\\tilde\\rho$ amounts to -0.45. In fact, we should not expect more, as a number of detected modes is much lower than theory predicts. We have to remember also that the values of [m/H] give information mainly about photospheric metal abundances. Moreover, we do not take into account phenomena like diffusion, element mixing etc.  To improve this analysis, a few things must be done. Firstly, we need determinations of the [m/H] parameter for a larger sample of $\\beta$ Cep stars. Secondly, a high accuracy in the period search analysis is needed to estimate better the range of the excited modes in a given star. Here put our hope on next multi-site and multi technique campaigns, as well as on space missions. Finally, the analysis of the chemical composition of $\\beta$ Cep variables from high resolution spectra would be very helpful.  These results require further studies and must be treated with caution. Firstly, they were obtained for small samples of stars, and secondly, we do not know how many modes from the excited ones have been already identified. Our aim was only to show some properties which could be extracted from such determinations of the metal abundance parameter [m/H] for $\\beta$ Cep variables. Having in mind that from a theoretical point of view all $\\beta$ Cep stars should be multiperiodic ones, the metallicity dichotomy we obtained is rather a qualitative confirmation of the dependence of the pulsation instability in these stars on the heavy element abundance."
    },
    "0410/astro-ph0410674_arXiv.txt": {
        "abstract": "The \\emph{Hubble Space Telescope} is uniquely able to study planets that are observed to transit their parent stars.  The extremely stable platform afforded by an orbiting spacecraft, free from the contaminating effects of the Earth's atmosphere, enables \\emph{HST} to conduct ultra-high precision photometry and spectroscopy of known transiting extrasolar planet systems.  Among \\emph{HST}'s list of successful observations of the first such system, HD~209458, are (1) the first detection of the atmosphere of an extrasolar planet, (2) the determination that gas is escaping from the planet, and (3) a search for Earth-sized satellites and circumplanetary rings. Numerous wide-field, ground-based transit surveys are poised to uncover a gaggle of new worlds for which \\emph{HST} may undertake similar studies, such as the newly-discovered planet TrES-1.  With regard to the future of \\emph{Hubble}, it must be noted that it is the only observatory in existence capable of confirming transits of Earth-like planets that may be detected by NASA's \\emph{Kepler} mission.  \\emph{Kepler} could reveal Earth-like transits by the year 2010, but without a servicing mission it is very unlikely that \\emph{HST} would still be in operation. ",
        "introduction": "When both the photometric transits and radial velocity variations due to an extrasolar planet are observed, we are granted access to key quantities of the object that Doppler monitoring alone cannot provide.  In particular, precise measurements of the planetary mass and radius allow us to calculate the average density, and infer a composition.  Such estimates enable a meaningful evaluation of structural models of these objects, including whether or not they possess a core of rocky material.  These inferences, in turn, enable direct tests of competing scenarios of planet formation and evolution. Moreover, the transiting configuration permits numerous interesting follow-up studies, such as searches for planetary satellites and circumplanetary rings, and studies of the planetary atmosphere.  In this review presented at the Space Telescope Science Institute's 2004 May Symposium \\emph{``From Planets to Cosmology: Essential Science in Hubble's Final Years''}, I discuss the status of work in the field with a focus on the key contributions, both past and near-future, enabled by the \\emph{Hubble Space Telescope}. I begin by reviewing the properties of the transiting planets discovered to date, as well as the the numerous ground-based efforts to detect more of these objects.  I then consider the various follow-up studies of these gas-giant planets that are enabled by \\emph{HST}, as well as \\emph{HST}'s potentially unique role in following-up Earth-sized objects detected by NASA's \\emph{Kepler} mission.  I finish by discussing two \\emph{HST}-based searches for transiting gas-giant planets.  Throughout this contribution, I restrict my attention to studies of \\emph{transiting} planets; for an introduction to the broader range of \\emph{HST}-based studies of extrasolar planets, see Charbonneau (2004). ",
        "conclusions": ""
    },
    "0410/astro-ph0410397_arXiv.txt": {
        "abstract": "The central engine that drives gamma ray burst (GRB) explosions may derive from the ability of electrons/positrons and nucleons to tap into the momentum and energy from the large neutrino luminosity emitted by an accretion disk surrounding a black hole. This transfer of momentum and energy occurs due to neutrino absorption, scattering, and annihilation and the non-spherical geometry of the source both increases the annihilation efficiency and, close to the black hole, directs the momentum transfer towards the disk axis. We present annihilation efficiencies and the momentum/energy transfers for a number of accretion disk models and compute the critical densities of infalling material below which the transfer of neutrino momentum/energy will lead to an explosion. Models in which the neutrinos and antineutrinos become trapped within the disk have noticeably different momentum and energy deposition structure compared to thin disk models that may lead to significant differences in the explosion dynamics. ",
        "introduction": "Gamma ray bursts were first observed forty years ago, but only recently has significant progress been made in understanding their origin. Evidence is mounting that at least the `long duration' bursts are associated with a rare type of supernova event such as a `failed supernova' or `collapsar' \\citep{woo93,pac98,mac99,mac01,prog03}, with neutron-star mergers as a candidate for the shorter bursts \\citep{pac91,ruf99,ros02}. In either case, it is likely that an accretion disk surrounding a black hole forms that cannot cool efficiently by photon emission because of the high densities and temperatures. But with their much smaller cross sections the neutrinos, if any are produced, may escape and cool the disk.  The electron neutrinos and antineutrinos emitted from the disk are produced from electron capture by protons or positron capture by neutrons respectively. Both reactions modify the make-up of the disk with electron capture initially predominating as material moves inward toward the black hole. The frequency of these reactions (and their inverse) are strong functions of the thermal properties of the disk and grow considerably as material moves towards the black hole where the temperatures and densities are highest. In turn, the density and temperature are functions of the mass, $M_{BH}$, and spin, $a$, of the black hole, the viscosity, $\\alpha$, of the disk and the accretion rate, $\\dot{m}$ since these four parameters determine the potential and the rate at which material flows radially inward. In particular, as either $\\dot{m}$ or $a$ increase so does the density of the disk (particularly in the hottest region close to the black hole) leading to greater neutrino luminosity.  The neutrino scattering length within the disk is similarly affected by the thermal properties. If sufficiently high temperatures are reached the neutrinos and/or antineutrinos become trapped and, consequently, their spectrum is altered. Due to the change in nucleon composition as material flows inward, neutrinos and antineutrinos will have different optical depths and will not become trapped at the same radius. For the lower accretion rate disks, $\\dot{m} \\la 0.1 \\,\\rmn{M_{\\sun} \\,s^{-1}}$, the density of the disk is never sufficient to trap the neutrinos and so they immediately escape. But for the higher accretion rates the neutrinos \\emph{do} become trapped. Once trapped the neutrinos thermalize and the neutrinos emitted from these regions have temperatures of a few MeV, the exact value depends on the model, and are functions of the radius. In addition, neutrino trapping also has the effect of restricting neutrino emission to that portion of the disk beyond a `neutrinosurface' which, given the cylindrical symmetry, are roughly toroid in shape.  Descriptions of disks with low accretion rates are given in \\citet{pwf99} while models which do include neutrino trapping are given in \\citet{dim02}. Additional calculations of neutrino emission from accretion disks are given in \\citet{set04}.  After escaping the disk the neutrinos and antineutrinos do not propagate unimpeded but may interact with material above the disk transferring energy and momentum. The extent to which their energy and momentum is tapped is the primary interest of this paper since the imparted energy may\\footnote{Though some of the required energy may come by electromagnetic extraction of energy from the rotating black hole, e.g. through the mechanism of \\citet{Bla77}}  power the burst. The transfer of momentum and energy occurs by a number of processes that we shall discuss in \\S \\ref{sec:nuinn}. We take calculations of the neutrino flux emitted from every point on the disk from previous work in \\citet{sur04}. Our results are presented in \\S \\ref{sec:results} and to investigate the effects of each of the four parameters needed to describe the disk we compute the energy/momentum transferred in a number of models. The parameters for each are listed in Table \\ref{table:models}. ",
        "conclusions": "The accretion disk model of the central engine of GRBs is an attractive one whose ultimate success will rely on the efficient extraction of the energy and momentum of the neutrinos emitted from the disk. The non-spherical geometry of the source makes this easier to achieve than in the context of a spherical source such as in a Type II supernova. The neutrinos are essentially emitted from a ring: a configuration that allows neutrinos to approach one another at larger interaction angles, $\\gamma$, with the concomitant increase in center of mass energy. The bulk of neutrino annihilation occurs in an oblate spheroidal region, extending upwards to $z \\sim 40 \\;\\rmn{km}$ and outward to $r \\sim 90 \\;\\rmn{km}$ for our model \\modonedpn, with the highest deposited power in the vicinity of the inner edge of the disk. In comparison with other azimuths, the power deposited along the z axis is a minimum. The electron/positron pairs would locally thermalize so one should expect a large pressure gradient directed toward the disk axis. In addition neutrino absorption and scattering creates a region, close to the black hole, in which the momentum transfer is also directed towards the disk axis. Trapping of neutrinos and antineutrinos within the disk leads to significant differences in the energy and momentum deposition structure particularly in the region close to the black hole. This may lead to noticeable changes in the explosion dynamics.  These results suggest a scenario whereby the electron/positron pairs created due to neutrino annihilation would, at least in part, move initially `inward and upward' facing little resistance from the $e^{+}/e^{-}$ pairs created along the z axis and accelerated by neutrino scattering from behind. At the disk axis the converging $e^{+}/e^{-}$ gas would collide, further increasing the internal energy, and then the hot, lepton rich gas flowing rapidly along the z axis. The movement of the $e^{+}/e^{-}$ gas created by neutrino annihilation would have consequences for an explosion. The results in table \\ref{table:densities} used only the power deposited along the disk axis to calculate the density of material that could be ejected but the movement of the $e^{+}/e^{-}$ pairs from the `hot spots' around the inner disk would suggest that these densities may be too small. The collision at the disk axis of the $e^{+}/e^{-}$ gas from around the torus from would raise the internal energy and, furthermore, the increase in the number density would also increase the neutrino opacity leading to a greater transfer of their momentum. Larger critical densities would also imply the more prompt formation of a jet."
    },
    "0410/astro-ph0410180.txt": {
        "abstract": "{ We present results of a detailed chemical analysis performed on 23 main-sequence turnoff stars having $-3.4\\leq$[Fe/H]$\\leq-2.2$, a sample selected to be highly homogeneous in T$_{\\rm eff}$ and log($g$). We investigate the efficiency of mixing in the early Galaxy by means of the [Mg/Fe] ratio, and find that all values lie within a total range of 0.2 dex, with a standard deviation about the mean of 0.06 dex, consistent with measurement errors. This implies there is {\\it little or no intrinsic scatter} in the early ISM, as suggested also by the most recent results from high-quality VLT observations. These results are in contrast with inhomogeneous Galactic chemical evolution (iGCE) models adopting present supernova (SN) II yields, which predict a peak-to-peak scatter in [Mg/Fe] as high as 1 dex at very low metallicity, with a corresponding standard deviation of about 0.4 dex. We propose that cooling and mixing timescales should be investigated in iGCE models to account for the apparent disagreement with present observations. The contrast between the constancy and small dispersion of [Mg/Fe] reported here and the quite different behaviour of [Ba/Fe] indicates, according to this interpretation, that Mg and Ba are predominantly synthesised in different progenitor mass ranges. ",
        "introduction": "The chemical evolution of the Milky Way arises from the continuous exchange of material between stars and the interstellar medium (ISM). The surfaces of long-lived, low-mass stars retain the composition of the ISM from their formation, so the early phases of the ISM are recorded in metal-poor stars. Metal-poor stars exhibit enhancements of oxygen and other $\\alpha$-elements relative to iron, which are explained by the differing yields of supernovae (SN) of type II and type Ia. The O comes from SN II (Thielemann {\\it et al.} \\cite{Thielemann90})), whereas Fe is produced by both SN Ia and SN II (Nomoto {\\it et al.} \\cite{Nomoto84})). It is usually inferred that the early Galaxy must have been enriched only by SN II; only more recently have SN Ia started to pollute the ISM, reducing the $\\alpha$-enhancement at higher metallicities (Tinsley \\cite{Tinsley1979}). This implies that the chemical composition of metal-poor halo stars is determined exclusively by SN II events and subsequent ISM mixing in the early Galaxy.  Trends of the abundances of different elements with metallicity are not the only tracers of the enrichment of the Galaxy. The scatter in the relative abundances of low-metallicity stars indicate the level of mixing of the early ISM. Whether all of the observed scatter is an intrinsic property of the ISM, and thus due to incomplete mixing in the early Galaxy, or is possibly due to errors that exceed the formal estimates of uncertainties, is still unclear. Nucleosynthesis models invoke various sources to explain the observed abundances of heavy elements (e.g. Travaglio {\\it et al.} \\cite{Travaglio04}), rising the level of uncertainty in their nucleosynthesis production. However, for some lighter elements, such as the $\\alpha$-elements O and Mg, their expected modes of synthesis do not predict differences in the ejected yields for stars of similar mass. Any scatter observed in the abundances of these elements should thus reflect more directly the effects of enrichment and mixing events in the early Galaxy.  SN II progenitors evolve on timescales of a few to tens of Myrs. In the earliest phase of the Galaxy there may have been a period of star formation during which the zones enriched by SN II have not yet mixed prior to the formation of the next stars. In this case, the newly forming stars will have different chemical compositions. Audouze \\& Silk (\\cite{AS95}) proposed that, at very low metallicity, clouds in the ISM could have been polluted by a maximum of three SN II, consistent with the suggestion by Ryan {\\it et al.} (\\cite{RNB91}b) that single SN were sufficient for the enrichment of the gas that formed stars at [Fe/H] $\\sim - 3.5$. If the early ISM was dominated by local inhomogeneities, as predicted by inhomogeneous Galactic chemical evolution (iGCE) models (e.g. Ishimaru \\& Wanajo \\cite{IW99}; Tsujimoto \\& Shigeyama \\cite{Tsujimoto99}; Argast {\\it et al.} \\cite{Argast00}; Travaglio {\\it et al.} \\cite{Travaglio01}), and the [O/Fe] and [Mg/Fe] ratios depend on SN-progenitor mass or metallicity, then [O/Fe] and [Mg/Fe] should show a scatter in the abundances trend versus metallicity, the surface abundances of these stars reflecting the local enrichment of the ISM. Argast finds, for example, that although the ISM is well mixed at [Fe/H] $> -2.0$, there is essentially no mixing at [Fe/H] $< -3.0$. This gives rise to progressively more star-to-star scatter as [Fe/H] decreases below $-2.0$.  Although there have been a large number of observational studies, the usually required agglomeration of results from different authors introduces the possibility that the reported scatter of the $\\alpha$-element abundance may be due to a lack of internal consistency, as discussed by Norris {\\it et al.} (\\cite{Norris01}). Underscoring this possibility, Magain (\\cite{Magain87}) found that an homogeneous re-analysis of literature data comprising 21 stars over the range $-3.6 \\leq $[Fe/H]$ \\leq -0.9$ gave a mean Mg abundance [Mg/Fe]=0.45 and a standard deviation, as a measure of the scatter, of 0.1 dex (scatters of the order 0.15 dex were reported for [Ca/Fe] and [Si/Fe]). Ten of the stars studied by Magain fell in the interval $-3.4 \\leq $[Fe/H]$ \\leq -2.2$ which we study in the present paper. A similar result was found by Nissen {\\it et al.} (\\cite{Nissen94}) for higher metallicities: [Mg/Fe]=0.41 and standard deviation 0.07 dex, over the range $-2.0 \\leq $[Fe/H]$ \\leq -1.0$. Recently, Carretta {\\it et al.} (\\cite{Carretta02}) found that the scatter in their estimate of intermediate mass elemental abundances was compatible with their estimated observational uncertainties. The most recent study of Cayrel {\\it et al.} (\\cite{Cayrel03}), based on extremely high-quality VLT/UVES observations of giant stars from the HK survey, exhibit a low scatter in the abundances of most elements, down to metallicities as low as [Fe/H]=$-4.2$. These authors report a standard deviation about their mean [Mg/Fe] value of 0.13 dex, and as low as 0.05 dex for [Cr/Fe]. This indeed supports the hypothesis that there is no large scatter in the early ISM.  Other studies, e.g., Fuhrmann {\\it et al.} (\\cite{FAG95}) and Mashonkina {\\it et al.} (\\cite{MGTB03}) have reported higher scatters in the alpha-element ratios. Even though these studies suggest the existence of stars with a different enrichment history (see e.g., Shiegeyama \\& Tsujimoto (\\cite{ST03}) for a possible origin of metal-poor stars with low [$\\alpha$/Fe]), their average standard deviation is still of the order $\\sim 0.1$ dex, far below what is expected in the case of an early inhomogeneous ISM. However, among the eight extremely metal-poor stars ([Fe/H]$<-3.5$) considered in the review by Norris (\\cite{Norris04}), two objects show [Mg/Fe]$>1.0$. This indicates the existence of some real differences, but only for metallicity much lower than previously thought.  Clearly, studies of the mixing efficiency in the early Galaxy are fundamental for testing inhomogeneous GCE models. This can be accomplished by a detailed analysis of [Mg/Fe] in halo stars, and is the main purpose of this work. We concentrate on [Mg/Fe] because of the suitability of Mg measurements in stars at very low metallicity, and to avoid the still debated uncertainties on measurements of O. Studies of the scatter vs. metallicity for different stellar populations could also allow one to constrain the yields from SN II. Attempts to do this have so far been unsuccessful for individual stars (Chieffi \\& Limongi \\cite{CL02}) because of the difficulty of computing SN yields that match the composition of metal-poor stars. Fran\\c{c}ois {\\it et al.} (\\cite{Francois04}) analyse how SN II yields can be constrained by the observations on the basis of their homogeneous GCE model, suggesting that major revisions of SN II yields are required to match the elemental abundances observed at very low metallicities.  Iron yields are particularly difficult to predict, because Fe production depends on the uncertainties in the mass cuts adopted for SN II models as well as the degree of mixing prior to fallback (Nomoto {\\it et al.} \\cite{Nomoto97}), whereas Mg is produced in hydrostatic burning much farther out in the star, and hence is subject to fewer uncertainties. As a result, Mg could also be a better chemical chronometer compared to Fe; its study would aid our understanding of the age-metallicity relation, a key constraint on studies of chemical evolution and nucleosynthesis.  The primary motivation for this study is to produce an independent estimate of the level of elemental scatter in the early Galactic ISM, as unaffected as possible by uncertainties. Unlike the most recent study of Cayrel {\\it et al.} (2003), whose S/N ratio is higher ( $\\ge 200/1$, compared to our 100/1-150/1), and who thus obtain better precision on a single measurement, here we focus on the homogeneity of the sample, and demonstrate how our homogeneous chemical analysis leads to a better constraint on the observed scatter of our derived abundances. The aim is to avoid the introduction of star-to-star differences that may lead to increased and misleading errors. We thus concentrate on halo stars that are in the same evolutionary stage, have only small differences in stellar parameters, and make use of the same Mg absorption lines for all objects. It is of fundamental importance to understand whether the scatter we observe is intrinsic to the sample or due to deficiencies in the analysis. We present the results of the chemical analysis of a sample of main-sequence turnoff halo stars, focusing on [Fe/H] and [Mg/Fe]. We describe the definition of the sample and the observational data in \\S \\ref{sample}. Particular importance has been given to the use of techniques that produce the highest homogeneity of analysis. The methods adopted and the chemical analysis are presented in \\S \\ref{analysis}. Section \\ref{abundances} reports our results and the determination of their uncertainties. In \\S \\ref{discussion} we analyse the results and focus on their possible implications. ",
        "conclusions": "the ISM from which stars formed at [Fe/H] $>-3.5$ was on the whole already well-mixed, {\\it and} it was enriched by a sufficient number of SN II to show an IMF averaged pattern at [Fe/H]$<-3.0$, in spite of the expectations to the contrary from iGCE current models. Furthermore, the hypothesis by Cayrel {\\it et al.} (\\cite{Cayrel03}) of the existence of a plateau in the [el/Fe] ratios at very low metallicity is clearly supported by our [Mg/Fe] ratio.  We are left to puzzle over how this IMF average is achieved ove much of the halo on the short timescales ($\\leq$ 30 Myr) for the ejection of the full range of SN II progenitor masses (M $\\geq$ 10M$_{\\odot}$). This could imply that we are observing the primordial abundances from the first stars below [Fe/H]=$-3.0$, as Cayrel et al. suggested.  \\subsubsection{Mixing and cooling timescales}  Enrichment timescales are not well constrained. Prantzos (\\cite{Prantzos03}) shows that an early phase of infall and relaxing the instantaneous recycling approximation causes his halo outflow model to reproduce better the metallicity distribution of stars in the halo. Enrichment occurs much quicker in his model than in iGCE models, but the introduction of an early infall phase in his model slows the metallicity enrichment again, as a result of which the time needed to reach [Fe/H]=$-3.0$ increases from 35 Myr to 100 Myr. A similar delay in the enrichment could occur in iGCE models for a strong early infall phase. However, a slower enrichment of the Galaxy does not affect the predicted scatter in [Mg/Fe] vs. [Fe/H], but rather the age-metallicity relation.  De Avillez \\& Mac Low (\\cite{AM02}) studied the mixing timescale in a SNe-driven ISM, and show that inhomogeneities due to single SN II explosions take up to 350 Myr to be erased (for the present Galactic SN II rate). Even when mixing scales as small as a few kpc are considered, efficient mixing requires of order 120 Myr. Increasing the SN rate by a factor 10, they find that inhomogeneities take only some tens of Myr to disappear. If increasing the SN II rate reduces the time needed for efficient mixing, then for a short time ($\\sim$ 30Myr) it favours the explosion of higher mass progenitors (because of their shorter lifetimes) and a higher average [Mg/Fe] ratio. This effect could also be achieved with a top-heavy IMF (discussed above). This would avoid the ISM having regions with low [Mg/Fe] values, but would still require a progenitor upper mass limit of 20 M$_{\\odot}$ to avoid the production of [Mg/Fe] $>$ 0.4 except for at the very lowest metallicities (Norris (\\cite{Norris04})).  We asked above how an IMF-averaged [Mg/Fe] value could be achieved on the very short evolutionary timescales of stars at [Fe/H]$<-3.0$. This may be the wrong question to ask if {\\it cooling} timescales crucial to star formation, rather than {\\it mixing} timescales, are the critical factor. iGCE models lack reliable treatments of the cooling of hot SN II ejecta in a metal-poor environment. If the cooling time exceeds the evolutionary time of the lowest mass SN~II progenitors of 10 M$_{\\odot}$, i.e. 30 Myr, then a complete IMF-average of the [Mg/Fe] ratio will be achieved, not on the evolutionary time of the SN II, but on the cooling timescale of the ISM, and therefore on the formation timescale of the next stellar generation.  Oey (\\cite{Oey03}) investigates mixing and cooling processes in the ISM by means of her simple inhomogeneous model (SIM) of Galactic chemical evolution, which explicitly incorporates interstellar mixing and mass transport. As previously found by De Avillez \\& Mac Low (\\cite{AM02}), diffusion is inefficient compared to turbulent mixing. Even though ``turbulent processes are extremely difficult to constrain'', turbulent mixing is found to be extremely efficient in the case of a hot ionized medium (HIM). Thus, in a phase of HIM-dominated ISM, efficient mixing could indeed occur. Furthermore, a hot ISM has a low cooling efficiency, thus a two-phase ISM could experience a delay due to the cooling time between the SN II events and the incorporation of their ejecta in star forming regions. Once new stars could form from SN-enriched gas, they would form out of a well- mixed ISM. By a comparison with the observed metallicity distribution function and the lowest metallicity observed in stars ([Fe/H]=--4.0 at the time of Oey's analysis), Oey's points out that an inhomogeneous phase could have been extremely short lived and that the ISM would have then mixed efficiently. Furthermore, Recchi {\\it et al.} (\\cite{Recchi01}) find that, in the case of dwarf galaxies, a single sturbust leads to the development of galactic winds and to a quick mixing. However, dwarf galaxies evolve in a low gravitational potential and thus dynamical processes (such as galactic wind) may have larger influences on the mixing than in the case of the Galaxy.  % This would thus % explain the far better agreement of GCE models than iGCE models with our % results in the metallicity range considered. % Oey warns of the possibility that superwinds could be generated in % a HIM, possibly causing the loss of metals through Galactic outflows.  We believe further investigations of mixing and cooling timescales in relation to iGCE models are advisable, since there is no clear theoretical expectation that the problem of the missing inhomogeneities will be solved by adjustments of the SN II yields, the SN II progenitor mass range, or modification of the IMF."
    },
    "0410/astro-ph0410438_arXiv.txt": {
        "abstract": "{ A calibration of H$\\alpha$ as both a chromospheric diagnostic and an age indicator is presented, complementing the works previously done on this subject (\\cite{herbig}, \\cite{luca1}). The chromospheric diagnostic was built with a statistically significant sample, covering nine years of observations, and including 175 solar neighborhood stars. Regarding the age indicator, the presence of stars for which very accurate ages are determined, such as those belonging to clusters and kinematic groups, lends confidence to our analysis. We also investigate the possibility that stars of the same age might have gone through different tracks of chromospheric decay, identifying - within the same age range - effects of metallicity and mass. These parameters, however, as well as age, seem to be significant only for dwarf stars, losing their meaning when we analyze stars in the subgiant branch. This result suggests that, in these evolved stars, the emission mechanism cannot be magnetohydrodynamical in nature, in agreement with recent models (Fawzy et al. 2002c, and references therein). The Sun is found to be a typical star in its H$\\alpha$ chromospheric flux, for its age, mass and metallicity. As a byproduct of this work, we developed an automatic method to determine temperatures from the wings of H$\\alpha$, which means the suppression of the error inherent to the visual procedure used in the literature. ",
        "introduction": "Losing angular momentum through magnetized stellar winds, cool main sequence dwarfs have their rotation continuously braked, reducing the efficiency of their dynamos and, consequently, their degree of chromospheric activity.  Because of this, the chromospheric filling observed in high opacity spectral lines can be translated into a potential indicator of age, a quantity which is still one of the most uncertain parameters in stellar astrophysics.  Although this scenario has a remarkable simplicity, it is a subject which has remained largely unexplored in a quantitative way. Among the works which have been published, most tend to focus on the H and K lines of Ca\\,{\\small {II}} (e.g. \\cite{skumanich}, \\cite{linsky}), in which the chromospheric emission is more obvious, but also more affected by transient phenomena and phase modulation. The H$\\alpha$ line, although widely used to measure chromospheric activity in solar physics, has received less attention in relation to this particular problem, and the lack of a good calibration of H$\\alpha$ is a drawback for several reasons. First, the line is less sensitive to transient phenomena like flares, coronal mass ejections and localized magnetic explosions; extremely energetic phenomena that flood the X-ray and ultraviolet spectra with energy, but barely affect the visible. Second, it has the property of characterizing the mean chromospheric flux in a better way than Ca\\,{\\small {II}} H and K because, showing less chromospheric filling, phase modulations within an activity cycle are greatly reduced: the errors in computing the flux -- due, for instance, to normalization and determination of effective temperatures --, largely overcome the intrinsic modulation. Third, the modern solid state detectors have higher quantum efficiency in the red, behaving inversely to the old photographic plates. Also, the studied stars --- solar type ones ---  have their maximum flux in the visible region, favoring better accuracy in narrow band photometry centered on H$\\alpha$.  Nevertheless, even if H$\\alpha$ did not present any advantage, the need for another diagnostic is crucial. As said before, the majority of works which attacked the problem until the 80s analyzed only the calcium lines, considering them representative of the radiative losses in the chromosphere. The core of H$\\alpha$, however, is formed in different regions (Schoolman 1972), thus responding differently to changes in the physical conditions of the chromosphere. In this sense, calibrating H$\\alpha$ will help not just to better determine the energy budget, but also to better discern the structure of upper stellar atmospheres. With this goal, Herbig (1985, hereafter H85) and Pasquini \\& Pallavicini (1991, hereafter P$^2$91) presented the only works measuring the net chromospheric fluxes in H$\\alpha$, finding consistent results using two different photometric wavebands to calibrate the spectra into absolute flux. Of these two attempts, one (H85) went so far as to develop an age indicator; however, the small sample of 43 stars used, in which all but two are field stars, did not allow for a sufficiently accurate analysis.  \\begin{table*} \\begin{center} \\caption[]{Parameters of the three analyses of chromospheric activity on H$\\alpha$ to date} \\label{param} \\begin{tabular}{c c c c}\\hline  & Herbig    & Pasquini \\& Pallavicini   & This work         \\\\ & 1985      & 1991              & 2004          \\\\ \\hline Resolution in the center of H$\\alpha$ ({\\AA})   & 0.74          & 0.11              & 0.30          \\\\ N$^{\\d{o}}$ of stars        & 43        & 87            & 175           \\\\ Range of Temperature Classes    & F8-G3     & F8-K5         & F5-K0         \\\\ Luminosity Classes  & V         & IV \\& V           & IV \\& V       \\\\ Photometric Waveband    & Johnson's V & Willstrop's $\\lambda\\lambda$\\,6650-6600\\,{\\AA} & Willstrop's $\\lambda\\lambda$\\,6650-6600\\,{\\AA}\\\\ Age Indicator       & Yes       & No                & Yes           \\\\ \\hline  \\end{tabular} \\end{center} \\end{table*}  In this work, we intend to fill this gap, calibrating H$\\alpha$ as an absolute diagnostic of chromospheric activity, using a statistically significant sample of 175 solar neighborhood stars, and building an age indicator based on this diagnostic, calibrated with stars belonging to open clusters and stellar kinematic groups (SKG). We also use a large sample of field stars with accurate ages derived from theoretical isochrones. This large sample is unique in the literature, and allows for a better statistical analysis. For comparison, the parameters of the three analyses on the H$\\alpha$ chromospheric flux are shown in Table~\\ref{param}. This paper is divided as follows: in section 2 we describe the observations and reduction procedures. In the 3$^{\\rm rd}$ we build the chromospheric diagnostic with the purpose of, in section 4, calibrating the age indicator. In section 5 we investigate the influence of other parameters in the degree of chromospheric activity, outlining the results and proceeding to the conclusion in the 6$^{\\rm th}$ section. ",
        "conclusions": "With high quality data, we developed an accurate chromospheric diagnostic based on the radiative losses in H$\\alpha$ using the Barnes-Evans relation (\\cite{barnes}) and the photometric waveband $\\lambda\\lambda$\\,6550-6600\\,{\\AA} defined by Willstrop (1965), using the procedure established by P$^{\\rm 2}$91. The precision achieved, with an uncertainty of 0.45 (in units of 10$^{\\rm 5}$\\, erg\\,cm$^{\\rm -2}$\\,s$^{\\rm -1}$) is one of the best in the H$\\alpha$ literature. This is due to many reasons, including that the size of the sample used reduced the error inherent to the photospheric subtraction; also, the broader coverage of our spectra, compared to those used by P$^2$91, allowed for a reliable normalization; finally, the sample was observed exclusively with solid state detectors. We also found evidence that the chromospheric emission, as measured in H$\\alpha$, is insensitive to the activity cycle and rotational modulation, hence characterizing the mean chromospheric flux in a more adequate way than the Ca\\,{\\small {II}} K line.  As a byproduct of the analysis, we developed an elegant automated method to determine effective temperatures using the wings of H$\\alpha$, structured on the classical K-S test. Tests show that the two methods, visual and automatic, agree to within one standard deviation, which supports its efficiency. The automation of the procedure means, mainly, the suppression of the personal error in the visual procedure, as well as a great economy of time when analyzing several spectra.  With this diagnostic, we calibrated the age-activity relation using stars belonging to open clusters, as well as to SKGs, finding a well defined relation until $\\sim$2 Gyr. By inserting field stars with good age solutions, it is found that a significant numer of subgiants do not follow the relation, hence one should not expect their emission mechanism to be magnetohydrodynamical in nature. This result agrees with recent models that credit them as powered by acoustic wave heating, generating basal flux in the low chromosphere.  Regarding the dwarfs, we find that mass and metallicity differences seem to be needed to explain their emission, which is consistent with the dynamo model. It is also verified that the age-activity relation loses sensitivity after 2\\,Gyr, flattening beyond this value. The method is then complementary to the isochrones.  For future work, we intend to refine the calibration, inserting more clusters and kinematical groups, and to investigate the spread within the clusters as an effect of mass. Our present data do not allow for this analysis because the individual samples, with a maximum of 6 stars (Hyades and Ursae Majoris) are not statistically significant. In the immediate future, we shall compare the results presented here, with the ones of another work (in preparation), using the Ca\\,{\\small {II}} H \\& K lines. Such a study should contribute both to refining our knowledge of the detailed structure of the chromospheric activity vs. age relationship, and its stratification with stellar parameters, and to provide additional observational constraints to the study of upper stellar atmospheres."
    },
    "0410/astro-ph0410112_arXiv.txt": {
        "abstract": "{We present the first results from simulations of processes leading to planet formation in protoplanetary disks with different metallicities. For a given metallicity, we construct a two-dimensional grid of disk models with different initial masses and radii ($M_0$, $R_0$). For each disk, we follow the evolution of gas and solids from an early evolutionary stage, when all solids are in the form of small dust grains, to the stage when most solids have condensed into planetesimals. Then, based on the core accretion - gas capture scenario, we estimate the planet-bearing capability of the environment defined by the final planetesimal swarm and the still evolving gaseous component of the disk. We define the probability of planet-formation, $P_p$, as the normalized fractional area in the ($M_0$, $\\log R_0$) plane populated by disks that have formed planets inside 5~AU. With such a definition, and under the assumption that the population of planets discovered at $R$ $<$ 5~AU is not significantly contaminated by planets that have migrated from $R$ $>$ 5~AU, our results agree fairly well with the observed dependence between the probability that a star harbors a planet and the star's metal content. The agreement holds for the disk viscosity parameter $\\alpha$ ranging from $10^{-3}$ to $10^{-2}$, and it becomes much poorer when the redistribution of solids relative to the gas is not allowed for during the evolution of model disks.} ",
        "introduction": "Well over a hundred extrasolar planets have already been catalogued. It is very likely that radial velocity surveys performed in the solar neighbourhood have discovered a significant fraction of all Jupiter-like planets with periods shorter than $\\sim$10 yr (or, equivalently, with orbital radii smaller than $\\sim$5~AU), whose parent stars exhibit periodic radial velocity variations with amplitudes $\\geq 10$~m~s$^{-1}$ Such a collection of planets provides a good data set with which predictions of different theories of planet formation can be compared \\citep[e.g.][]{ida04}.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fe_bargraph_bw.eps}} \\caption{The probability of planet occurrence as a function of the star's iron abundance (on the standard logarithmic scale with zero representing the solar value). Reproduced with permission from \\cite{debra03}} \\label{f:fe_bar} \\end{figure}  Current observations strongly suggest that  planet-bearing stars tend to have higher metallicities than field stars (Fig. \\ref{f:fe_bar}). This correlation was reported for the first time by \\cite{santos00}, and later confirmed by \\cite{debra03}. The effect finds a natural explanation in the core accretion - gas capture (CAGC) model for giant planet formation. The model predicts a slow accretion of planetesimals onto a protoplanetary core until the mass of the core reaches a few $M_\\oplus$. Then, gas accretion becomes significant, and around the core an extended envelope is formed, whose mass increases faster than that of the core. Eventually, the envelope becomes more massive than the core, and a runaway accretion of gas ensues, which terminates when tidal effects become important or the protoplanetary disk is dissipated. The length of the core accretion phase decreases with increasing surface density of the planetesimal swarm \\citep{pollack96}, which, in turn, is an increasing function of the original metal content of the protoplanetary disk and its central star. Consequently, the giant planets are expected to form more easily in disks with higher metallicities.  However, dust evolves due to different physical processes than  the gaseous component of the disk \\citep{weiden93} As a result, a significant redistribution of solids takes place with respect to the gas, and, locally, the surface density of solids can be considerably enhanced compared to the initial one \\citep{weiden03,SV97}. Thus, statistical tests of the CAGC scenario must be based on evolving models of protoplanetary disks rather than static configurations which only serve as mass supply to planetary cores and envelopes. Such models must include the global evolution of solids. An extensive survey of such models was reported by \\cite{kac1}, and applied to the problem of extrasolar planets formation by \\cite{kac2}. In the latter paper it was found that effects of dust evolution allow for \\textit{in situ} formation of giant planets as close as 2 AU from the star. Here we obtain a much larger sample of disk models with different primordial metallicities, estimate how many of them could form planets, and compare the estimates to the observed correlation between stellar iron abundance and the probability of planet occurrence.  In \\S 2 we briefly explain our approach to the evolution of a protoplanetary disk and planet formation. The results are presented in \\S 3, while in \\S 4 we summarize them and critically assess the simplifications on which our method is based. ",
        "conclusions": "Based on a simple approach to the evolution of solids in protoplanetary disks we calculated rates of giant planet occurrence on orbits smaller than 5~AU. We find that in more metal-rich disks the planets form more easily, and we qualitatively reproduce the observed correlation between the probability that a giant planet would be found at a star and the metal content of that star. We are able to match this correlation quantitatively for disk models with the viscosity parameter $\\alpha$ ranging from $\\sim 10^{-3}$ to $\\sim10^{-2}$. We also show that in the low-metallicity regime the calculated occurrence rates agree with the observed ones only if the redistribution of solids is included in the models. In such models the surface density of solids locked in the final planetesimal swarm can be locally much higher than the initial surface density of the dust. This effect is particularly important if we want to match the observed value of limiting stellar metallicity $Z_{min}$ below which no planets have been detected. Again, only models which allow for the redistribution of solids are successful, while models which do not include this effect produce $Z_{min}$ which is significantly too high. This finding supports our approach to the evolution of protoplanetary disks, whose key ingredient is the radial drift of solids due to the gas drag.  The above conclusions are based on the assumption that our grid of models is a representative sample of the real distribution ${\\cal P}(M_0,R_0)$ of initial masses and radii of protoplanetary disks, which is not known. Such an assumption is acceptable provided that the real disks are uniformly (or almost uniformly) distributed on the $[M_0,\\lg R_0]$ plane. Adopting it we demand that very large disks be relatively rare, which seems to be an entirely reasonable requirement (note that ${\\cal P}(M_0,R_0)$ {\\it must} approach 0 for $R_0\\rightarrow\\infty$). We tested the sensitivity of our results to the way the initial parameters are sampled by repeating the calculations with initial models uniformly distributed on the $[M_0,\\lg j_0]$ plane. The results, normalized to the observational data the same way as before, are shown in Fig. \\ref{f:j_sc}. The agreement is poorer, but the same general trend is still well visible. Moreover, we see that $Z_{min}$ does not depend on the sampling procedure (this conclusion is restricted to samples which do not include extremely massive disks with $M_0>0.2~M_\\odot$; note however that such disks would most likely be gravitationally unstable).  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{j_in_sc_new.ps}} \\caption{Same as Fig. \\ref{f:rout_sc}, but based on disk models with initial parameters uniformly distributed on the $[M_0, \\lg j_0]$ plane.} \\label{f:j_sc} \\end{figure}  Our results are based on a highly simplified scenario of the evolution of solids in protoplanetary disks. The assumption concerning the size distribution of solid particles, according to which at each distance from the star all particles have the same diameter, seems to be particularly radical. However, detailed work by \\cite{morfill85},\\cite{mizuno88} and \\cite{weiden97} showed that the distribution of solids quickly converges to a stage in which most of the mass is concentrated in a narrow range of sizes at the momentary maximum size. Their conclusion has recently been corroborated by the behaviour of solids in the two-dimensional disk model by \\cite{weiden03}. Thus, despite its simplicity, our our approach to the size disitribution function can be regarded as a reasonable approximation.  Our next radical assumption is the 100\\% efficiency of coagulation In reality, one hardly expects that collisions between solid particles always lead to sticking without fragmentation. However, calculations reported by \\cite{kac4} show that if the main process responsible for particle growth is collisional coagulation, then its efficiency {\\it must} be high - otherwise the disk would rapidly lose most of its solid material, disabling planet formation or severely diminishing its possibility. Note also that recent observations by \\cite{chakrab04} indicate that asteroid-sized bodies can exist in disks as young as 0.1 Myr. Again, if they have grown due to coagulation, the sticking efficiency must have been very high.  Migration of planets due to gravitational interactions with the disk is not included in our approach. Obviously, we do not claim that the process of migration does not operate. On the contrary, we have to assume some form of  redistribution of planets in order to explain the origin of \"hot Jupiters\", whose orbits are much tighter than the tightest orbits allowed by our models ($\\sim 1$~AU). However, as long as we are interested only in global rates of planet occurrence, and not in the distribution of orbital parameters, our results remain valid provided that the number of planets which migrate to $r<5$~AU from larger orbits is negligible compared to the number of planets which are born there.  The fact that such a simple model can reproduce the available data has two possible interpretations. According to the pessimistic one we observe a mere coincidence which does not have any physical meaning. According to the optimistic one we identify the most important process governing the evolution of solids in a protoplanetary disk (radial drift, induced by the gas drag, and associated with collisional coagulation), and we provide support for the core accretion - envelope capture scenario of giant planet formation. More sophisticated disk models based on fewer assumptions, which we are presently working on, should help us decide which of the two is true.  Acknowledgments. KK and MR were supported by the grant No. 1 PO3D 026 26 from the Polish Ministry of Science. They gratefully acknowledge benefits from the activities of the RTN network \"PLANETS\" supported by the European Commission under the agreement No HPRN-CT-2002-00308. This work was supported in part by the NASA Origins of Solar Systems program through grant NAG 5-13285 to the University of California, Santa Cruz.  KK was also supported by the German Research Foundation (DFG) through the Emmy Noether grant WO 857/2-1. TS was supported by the LPI, which is operated by USRA under contract CAN-NCC5-679 with NASA. This is LPI contribution No. 1214."
    },
    "0410/astro-ph0410054_arXiv.txt": {
        "abstract": "We have conducted spectropolarimetry of 12 type II (obscured) quasar candidates selected from the spectroscopic database of the Sloan Digital Sky Survey based on their emission line properties. Polarization was detected in all objects, with nine being highly polarized ($>3\\%$) and with polarization reaching as high as 17\\% in two objects. Broad lines were detected in the polarized spectra of five objects. These observations prove beyond a reasonable doubt that the objects in our sample are indeed type II quasars, in that they harbor luminous UV-excess AGNs in their centers and that the direct view to the AGN is highly obscured. For three of the objects in this paper, we have obtained HST images in three bands. The HST observations, combined with the spectropolarimetry data, imply that scattering off material outside the obscuration plane is the dominant polarization mechanism. In all three objects the sizes of scattering regions are a few kpc. For one object, the extent of the scattering region, coupled with the characteristics of the polarized spectrum, argue strongly that dust scattering rather than electron scattering dominates the polarized light. Our observations are well-described by the basic orientation-based unification model of toroidal obscuration and off-plane scattering, implying that the model can be extended to include at least some high-luminosity AGNs. ",
        "introduction": "Unification models of active galactic nuclei (AGNs) aim to explain the differences in observed AGN properties as orientation effects due to the non-isotropic nature of obscuring material \\citep{anto93}. In these models, the observer's view to the central engine can be blocked along some lines of sight by very optically thick, toroidally concentrated dusty material, making UV, optical and soft X-ray emission from AGNs highly anisotropic. A strong UV continuum and broad permitted emission lines that originate very close to the central engine are typical of sight-lines without significant obscuration (type I AGNs). These features are not observed in unpolarized optical/UV light in sources with large amounts of obscuration along the line of sight (type II AGNs).  Unification models are well established for low-luminosity, nearby AGNs (Seyfert galaxies); whether they can be extended to the luminosity and redshift range of genuine quasars has long been controversial. Type II quasars are the luminous analogs of Seyfert 2 galaxies, but despite their high intrinsic luminosities they have been very hard to find, and only a handful of candidates have been described in the literature \\citep{klei88,mcca93,hine95,daws01,djor01,norm02,schm02,ster02,dell03,smit03}.  We have identified about 150 type II quasar candidates in the redshift range $0.3<z<0.8$ from the spectroscopic data of the Sloan Digital Sky Survey (\\citealt{z03}, hereafter Paper I). Selecting such a large number of these optically faint objects is possible because of the size of the spectroscopic survey (currently containing $3.7 \\times 10^5$ publicly available spectra; \\citealt{abaz04}), part of which is dedicated to the search for unusual objects. We are now conducting a multi-wavelength follow-up campaign of the most luminous objects in our sample to constrain physical models of the circumnuclear obscuration and its geometry, and to probe the unification model at high luminosities.  The first strong evidence for the unification model of AGNs came from spectropolarimetry of Seyfert galaxies pioneered by \\citet{anto85} for NGC1068. In many Seyfert 2 galaxies, while the direct emission from the AGN is obscured, some of the light from the central engine and the broad-line region is scattered by material outside the plane of the obscuring torus and reaches the observer. Because scattered radiation is polarized, observations of the polarized flux can thus offer an indirect view of the central parts of the obscured AGNs.  In this paper we present spectropolarimetry of twelve type II quasar candidates, probing the central parts of these AGNs otherwise hidden from the observer.  In Section \\ref{sec_data} we describe the observations and data reduction. In Section \\ref{sec_results} we summarize the results of the spectropolarimetry observations and present example spectra. We discuss different mechanisms of producing polarization in AGNs in Section \\ref{sec_discussion}, along with related observations obtained by the Hubble Space Telescope (HST). We summarize our results in Section \\ref{sec_conc}. An $h=0.7$, $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$ cosmology is assumed throughout. Optical emission and absorption lines are identified by their air wavelengths in \\AA\\ (e.g., [OIII]5007). Objects are identified by their J2000 coordinates in Table 1 (e.g., SDSS J084234.94+362503.1) and shortened to $hhmm+ddmm$ notation  in the text (SDSS~J0842+3625). ",
        "conclusions": "\\label{sec_conc}  Spectropolarimetry of AGNs offers the possibility of observing their central regions along multiple lines of sight. Some of the light from the AGN is scattered off the surrounding material and is therefore polarized. In type II AGNs the direct light from the central parts of the AGN is so heavily extincted that spectropolarimetry provides the only way to explore the central engine at optical wavelengths.  Out of the 12 type II quasar candidates in this study, nine show high polarization ($P>3\\%$). If a crude host galaxy correction is applied then all objects are highly polarized by this criterion. The optical continua of most objects are dominated by the highly polarized scattered light and the unpolarized host galaxy component; the direct light from the AGNs is a negligible contribution except for one objects which is also a red 2MASS AGN and is atypical of our sample. This conclusion is reinforced by the absence of near-IR excesses for most of our objects (Paper II).  Broad emission lines were detected in five objects in the polarized spectra. This is direct evidence that type II quasar candidates from the SDSS harbor central regions similar to those of type I AGNs, with the defining characteristics of a blue continuum and broad emission lines. The presence of the broad-line region is not ruled out for any of our objects due to the low S/N of the spectropolarimetry of these faint objects, with a typical sensitivity to broad lines of about 100\\AA\\ in equivalent width.  Three objects with spectropolarimetric observations have also been imaged in three bands as part of our HST program (SDSS~J1039+6430, SDSS~J1323$-$0159 and SDSS~J1413$-$0142, Figure \\ref{pic_hst} and Paper IV). For these objects, the combination of spectropolarimetry and HST imaging strongly favors scattering as the mechanism producing polarization, since in all three cases the plane of polarization is orthogonal to the orientation of the presumed scattering regions. The wavelength dependence of the scattering efficiency can in some cases be used to discriminate between dust and electron scattering, but it is difficult to determine observationally because of the highly uncertain host galaxy fraction. We use other arguments to show that the object SDSS~J1039+6430 is a strong candidate for dust scattering. In addition to having a very large scattering region ($\\sim$10 kpc) it shows narrow emission lines in the polarized spectrum, which rules out scattering by warm and hot electrons. In this object there is also evidence that the polarized spectrum is bluer than the intrinsic spectrum, inconsistent with wavelength-independent electron scattering. In the remaining two objects SDSS~J1323$-$0159 and SDSS~J1413$-$0142 the scattering regions as seen in the blue band of the HST images have extents over 1 kpc, but scattering by warm electrons ($\\sim 10^5$ K) is not excluded as the main polarization mechanism.  All of these observations can be well described by a simple unification model with obscuration in a plane and scattering by material outside this plane. The size of the scattering region is comparable to the size of the narrow-line region, explaining the low polarization of the narrow-line emission if the narrow-line gas is intermixed with the scattering material. Typical extinction values in the obscuring plane should be well in excess of $A_V=7$. We suggest that the typical scattering efficiency is around 1\\%.  One out of 12 objects in our program (SDSS~J1715+2807) shows a non-stellar red continuum, polarization fraction rising to the red and a red polarized spectrum. This object is also a red 2MASS AGN, and it is better understood as a heavily reddened quasar rather than a completely obscured type II AGN \\citep{smit03}. We suggest that polarization by dichroic extinction in the obscuring torus with $A_V\\simeq 4$ cannot be excluded as the primary polarization mechanism in this object.  Finally, SDSS~J0842+3625, the most luminous type II quasar in our sample, shows an extraordinary wavelength-independent polarization level of 16\\%. A broad H$\\gamma$ line is detected in the polarized spectrum, as well as narrow H$\\beta$ and [OIII]4959, 5007. We estimate that the intrinsic luminosity of this object in the near-UV (rest-frame 2560$-$3840\\AA) is at least 10$^{46}$ erg sec$^{-1}$."
    },
    "0410/astro-ph0410748_arXiv.txt": {
        "abstract": "We present a study of analytic models of starless cores whose line profiles have ``infall asymmetry,'' or blue-skewed shapes indicative of contracting motions. We compare the ability of two types of analytical radiative transfer models to reproduce the line profiles and infall speeds of centrally condensed starless cores whose infall speeds are spatially constant and range between 0 and 0.2~km~s$^{-1}$. The model line profiles of \\hcop~(\\jequals{1}{0}) and \\hcop~(\\jequals{3}{2}) are produced by a self-consistent Monte Carlo radiative transfer code. The analytic models assume that the excitation temperature in the front of the cloud is either constant (``two-layer'' model) or increases inward as a linear function of optical depth (``hill'' model). Each analytic model is matched to the line profile by rapid least-squares fitting.  The blue-asymmetric line profiles with two peaks, or with a blue shifted peak and a red shifted shoulder, can be well fit by one or both of the analytic models. For two-peak profiles is best matched by the ``{\\sc hill5}'' model (a five parameter version of the hill model), with an RMS error of 0.01~km~s$^{-1}$, while the ``{\\sc twolayer6}'' model underestimates the infall speed by a factor of $\\sim 2$. For red-shoulder profiles, the {\\sc hill5} and {\\sc twolayer6} fits reproduce infall speeds equally well, with an RMS error of 0.04~km~s$^{-1}$. The fits are most accurate when the line has a brightness temperature greater than 3~K. Our most accurate models tend to reproduce not only the line profile shape, but also match the excitation conditions along the line of sight. A better match to the line profile shape does not necessarily imply a better match to the infall speed and provide guidance on how to minimize the risk of obtaining a poor infall speed fit.  A peak signal to noise ratio of at least 30 in the molecular line observations is required for performing these analytic radiative transfer fits to the line profiles. Moderate amounts of depletion and beam smoothing do not adversely affect the accuracy of the infall speeds obtained from these models. ",
        "introduction": "There are four distinct stages in isolated star formation: (1) formation of a gravitationally bound core withing a molecular cloud, (2) the gravitationally driven collapse of that core, (3) the formation of a central star-like object, (4) dispersal of the remaining cloud material \\citep*{sal}. The initial two stages are predominantly studied through millimeter- and submillimeter-wave-length observations. The molecular line rotational transition emission visible in these wavelength regimes allow one to probe both the physical and chemical composition of the molecular cloud cores out of which stars form. Spectral analysis of the molecular line emission also allows one to investigate the kinematic motions within the molecular cloud, particularly the phase of gravitational collapse \\citep{meo} as well as the associated bipolar outflows which are ubiquitous around young stars \\citep{bac}.  In a dense cloud core, commonly used molecular line tracers of moderate optical depth, such as CS~(\\jequals{2}{1}), or \\hcop~(\\jequals{1}{0}) tend to become self-absorbed. As the cloud contracts under the influence of gravity, the central regions tend to become more dense than the outer regions. This sets up a gradient in the excitation temperature of the molecular gas within the cloud such that the excitation temperatures in the central regions of the cloud are greater than those in the outer regions, even if the kinetic temperature is constant within the molecular cloud. If the cloud is spherical and static the spectral line emission from this cloud in these tracers will be symmetric about the line of sight velocity and self-absorbed. However, if there is radial motion around the molecular cloud core, a velocity gradient develops around the center of the core, precisely where the excitation temperature is the greatest. Much of the emission from the rear of the cloud is not absorbed as it travels through the core, due to the high excitation temperature within the core, and continues unabsorbed through the front section of the cloud because the velocity gradient between the rear and front sections of the cloud results in a Doppler shift which is significant compared to the broadening by thermal motions. Emission from the front of the cloud however is absorbed by nearby molecular gas with lower excitation temperature and only a small Doppler shift. In the case of inward radial motion the rear emission is blue-shifted and the front emission is red-shifted, yielding the ``classic'' blue-asymmetric infall line profile. This line profile has become the predominant tool for investigating infall motions within star-forming molecular cloud cores \\citep{sl,zhou2,zhou,mmtwbg,gemm,lmt2}.  Although this blue-asymmetry results from infall, it is not unique to that phenomenon. The superposition of multiple clouds along the line of sight can also cause an asymmetric line profile. By observing an optically thin molecular line, which should remain symmetric around the line of sight velocity of the cloud core, it is possible to reduce the likelihood of mistaking separate clouds for an infalling cloud. Asymmetric line profiles can also arise in kinematically complex regions. Studies have shown that rotation and outflow can also produce asymmetric line profiles \\citep{slp,al,nmwb}. Even the asymmetric line emission of B335, once thought to be the prototypical example of spectrally detected inside-out collapse \\citep{zekw,zekw2,zhou}, is now also thought to also contain an outflow which gives rise to the line wings previously modeled as part of the collapse \\citep{wmmt}. In spite of the multiple molecular cloud configurations that can result in blue-asymmetric line profiles, they are still the best way to study infall motions available, provided we understand the context in which they occur and make clear the necessary caveats.  Understanding how blue-asymmetric line profiles arise from infall motion is easy, but creating a complete model matching the molecular emission from one or more tracers is a difficult and time-consuming task. Typically one constructs an infalling physical model and uses LVG, Monte Carlo \\citep{cegw} or microturbulent \\citep{zhou} radiative transfer to model the emission. By iteratively changing the model parameters one is able to match some or all of the observations reasonably well. This process can take several minutes per iteration, and becomes prohibitively slow as the number of physical parameters becomes large and as the number of molecular species increases. The derived cloud parameters are also highly dependent on the input physical model. Changing the physical model can lead to an entirely different set of solutions.  ``Starless cores'' are useful regions for observations of inward motion because they are numerous, nearby, relatively simple in their density structure, and their motions are not confused by outflows. They have been the subject of several studies \\citep{lm,lmt1,lmt2,bapabw,mm,tmcwc,tmcw,ts,plm}, some of which have used radiative transfer models to fit their line profiles. In this paper we present simple analytic radiative transfer models, based on the simplest available assumptions of possible excitation temperature trends within starless molecular cloud cores. These models can easily be fit to blue-asymmetric self-absorbed molecular line profiles, and produce an estimate of the infall velocity suggested by those line profiles. We also test the efficacy of the analytic radiative transfer models by comparing the infall velocities derived using those models with actual infall velocities in rigorously modeled condensing molecular clouds. Our main finding is that our analytic radiative transfer models can yield good estimates for the physical conditions of the starless cores with few assumptions and little computational investment. ",
        "conclusions": "We have shown that analytic radiative transfer models are effective at estimating the infall velocities of contracting clouds from the blue-asymmetric line profiles of species of moderate optical depth in the region of collapse. The models provide good velocity estimates with few assumptions. The greatest advantage of the analytic models is that they allow for a quick and reasonable estimate of infall velocity with a minimum of computational effort. In cases where nothing else is known about the source, they provide the only possible estimate of infall, and in regions where a more robust model can be constructed, the analytic models provide an excellent starting point for further modeling \\footnote{Implementations of the models used for this paper are available from the authors at \\url{http://cfa-www.harvard.edu/\\~{}cdevries/analytic\\_infall.html}}. The main conclusions of this work are as follows:  \\begin{enumerate} \\item Analytic radiative transfer models can reproduce the blue-asymmetric line profiles seen in infalling starless cores. In addition to the two-layer model first presented in \\citet{mmtww} we have found a second class of radiative transfer models (the hill models) and explored 5 variants of these two classes of analytic radiative transfer models. All the model variations were able to produce blue-asymmetric line profiles similar to those we observe in starless cores.  \\item Successful analytic radiative transfer models reproduce not only the line profile shape, but also match the initial excitation conditions along the line of sight up to or beyond an optical depth of 1. When this requirement is met, the modeled infall velocity successfully fits the simulated infall velocity. Those models which match the line profile shape, but do not match the line of sight excitation conditions cannot reveal the infall velocity and velocity dispersion in the cloud.  \\item The {\\sc hill5} analytic radiative transfer model is the preferred model when fitting a blue-asymmetric line profile. In simulations with constant velocity as at all cloud radii the infall velocities attained by fitting the {\\sc hill5} analytic radiative transfer model to the simulated line profiles match the infall speed with an RMS error 0.02~km~s$^{-1}$. The {\\sc hill5} model tends to perform especially well when there are two distinct peaks such that the intensity difference between the self-absorption trough and the redshifted peak is greater than 10\\% of the intensity of the blueshifted peak. The RMS error in the infall determination of {\\sc hill5} fits to simulated spectral lines with those characteristics from Simulations~A and D is 0.01~km~s$^{-1}$.  \\item The physical parameters of simulated clouds that produce blue-asymmetric line profiles with no distinct shoulder or redshifted peak are generally not well determined by any of the model variations we have studied in this paper. We recommend not using these models to estimate infall velocities from line profiles that do not have a separate redshifted peak or a distinct redshifted shoulder.  \\item A better match to the line profile shape does not necessarily imply a better match to the excitation profile or the physical parameters, such as the infall velocity, of the observed cloud. This can be seen most clearly in figure~\\ref{dipshoulder} where local minimum of the ``dip'' solution often provides a lower $\\chi^{2}$ fit to the line profile shape, but a poorer fit to the infall velocity than the ``shoulder'' solution. Following the instructions in \\S\\ref{method} when fitting a blue-asymmetric line profile will minimize the risk of obtaining an unusually poor estimate of the physical parameters from the analytic radiative transfer model fit to the line profile.  \\item All the analytic models lose sensitivity to the infall velocity if there is a significantly optically thick static envelope between the observer and the infall region. In these cases a more optically thin line can provide a better estimate to the infall velocity in the obscured region. Conversely an optically thin tracer may underestimate the infall velocity of an infalling envelope around a static cloud core. These methods are sensitive to the velocities and conditions in the regions in which the line is becoming optically thick.  \\item A peak signal to noise ratio greater than or equal to 30 is required for the systematic error on the infall velocity estimate to be greater than the random error. We recommend achieving a peak signal to noise ratio of 30 in order to minimize the random error of a line profile fit using our analytic radiative transfer models, however nothing is gained by achieving a higher signal to noise level. Achieving this signal to noise level can take a prohibitively long time for lines whose peak antenna temperature is less than or equal to 1~K.  \\item Depletion and beam smoothing can greatly affect the observed line profile in infalling starless cores. We conclude that, even in the presence of these effects, if the line profile is asymmetric and a good candidate to be fit using these analytic radiative transfer models (as described in \\S\\ref{method}), then generally the radiative transfer models do yield an accurate estimate of the infall velocity. If the cloud is depleted to such an extent or viewed with such a large beam that the line profile becomes very weak or symmetric, then these models will not accurately probe the infall region of the starless core. \\end{enumerate}"
    },
    "0410/astro-ph0410262_arXiv.txt": {
        "abstract": "The Faint Sky Variability Survey (FSVS) is aimed at finding photometric and/or astrometric variable objects between 16th and 24th mag on time-scales between tens of minutes and years with photometric precisions ranging from 3 millimag to 0.2 mag. An area of $\\sim$23 deg$^2$, located at mid and high Galactic latitudes, was covered using the Wide Field Camera (WFC) on the 2.5-m Isaac Newton Telescope (INT) on La Palma. Here we present some preliminary results on the variability of sources in the FSVS. ",
        "introduction": "The FSVS data set consists of 78 INT WFC fields. Each field is divided into four, corresponding to the 4 2kX4k CCDs that form the WFC. For each field, we took one set of B, I and V images on a given night. Several more images were taken in the V band on that night and on consecutive nights. Typically, fields were observed in V 10 -- 20 times within one week. Exposure times were 10 min with a dead time between observations of 2 min.  This observing pattern allows us to sample periodicity timescales from 2$\\times$(observing time + dead time) (i.e. 24 min) up to the maximum time span of the observations (i.e. a few days). The combination of the colour information and the variability allows us to distinguish between different types of variable systems. See Groot et al. (2003) for a full discussion on the FSVS data. ",
        "conclusions": ""
    },
    "0410/astro-ph0410239_arXiv.txt": {
        "abstract": "This paper describes the Third Data Release of the Sloan Digital Sky Survey (SDSS).  This release, containing data taken up through June 2003, includes imaging data in five bands over 5282 deg$^2$, photometric and astrometric catalogs of the 141 million objects detected in these imaging data, and spectra of 528,640 objects selected over 4188 deg$^2$.  The pipelines analyzing both images and spectroscopy are unchanged from those used in our Second Data Release. ",
        "introduction": "The Sloan Digital Sky Survey (SDSS; York \\etal\\ 2000) is carrying out an imaging and spectroscopic CCD survey of the sky at high Galactic latitudes, using a dedicated wide-field 2.5m telescope at Apache Point Observatory in South-East New Mexico.  The telescope saw first light in May 1998, and following an extensive period of commissioning, formal survey operations began in April 2000.  The resulting data have been distributed to the public via web interfaces accessible from the SDSS public web site\\footnote{\\tt http://www.sdss.org}, and have been described in a series of papers: \\begin{itemize} \\item The Early Data Release, including data taken during commissioning, Stoughton \\etal\\ (2002; hereafter EDR paper), consisting of 462 deg$^2$ of imaging data and spectra of 54,000 objects. \\item The First Data Release, Abazajian \\etal\\ (2003; hereafter DR1 paper), consisting of imaging data over 2099 square degrees and spectra of 186,240 objects. \\item The Second Data Release, Abazajian \\etal\\ (2004; hereafter DR2 paper), consisting of imaging data over 3324 square degrees and spectra of 367,360 objects. \\end{itemize}  These data allow investigations in all areas of optical astronomy, from asteroids to quasars.  Among the papers which have appeared since the Second Data Release are studies of the stellar masses of galaxies (Brinchmann \\etal\\ 2004; Kauffmann \\etal\\ 2004), measurements of the dark matter power spectrum from the Ly $\\alpha$ forest (McDonald \\etal\\ 2004) and corresponding constraints on cosmological parameters (Seljak \\etal\\ 2004); studies of quasars (Lacy \\etal\\ 2004) and galaxies (Hogg \\etal\\ 2004) by cross-correlating SDSS data with the Spitzer First-Look Survey; the discovery of new structures in the halo of M31 (Zucker \\etal\\ 2004ab); measurements of the detailed shape of the galaxy correlation function and relation to halo occupation models (Zehavi \\etal\\ 2004); and studies of populations of stars at the bottom of the HR diagram (West \\etal\\ 2004; Knapp \\etal\\ 2004).  The SDSS imaging data are taken on photometric moonless nights (with photometricity determined by an auxiliary telescope; Hogg \\etal\\ 2001) of good seeing with a wide-field imaging camera operating in drift-scan mode (Gunn \\etal\\ 1998).  Six parallel {\\em scanlines} on the sky, each $13^\\prime$ wide, are observed by each of the columns of CCDs.  Each of the five rows of the CCDs in the camera is fronted by a different filter, thus each scanline is observed in five filters, denoted (in order of observation) $r\\,i\\,u\\,z\\,g$ (Fukugita \\etal\\ 1996; Gunn \\etal\\ 1998; EDR paper).  The imaging data are processed by automated software pipelines that measure the properties of all detected objects (Lupton \\etal\\ 2001), and perform astrometric (Pier \\etal\\ 2003) and photometric calibration, the latter to a set of standard stars observed with the Photometric Telescope (Smith \\etal\\ 2003).  The resulting object catalogs are used to select targets for spectroscopy, including the ``Main'' sample of galaxies, magnitude-limited to $r_{Petrosian} = 17.77$ (Strauss \\etal\\ 2002), a sample of luminous red ellipticals to $r_{Petrosian} = 19.5$ (Eisenstein \\etal\\ 2001; DR2 paper), quasar candidates selected by their colors to $i = 19.1$ (for $z < 3$ candidates) and $i=20.2$ (for higher-redshift candidates) (Richards \\etal\\ 2002), as well as a host of additional targets, including optical counterparts to ROSAT X-ray sources (Anderson \\etal\\ 2003), unusual stars, and calibration observations (EDR paper).  All magnitude limits here are corrected for Galactic extinction following Schlegel, Finkbeiner, \\& Davis (1998). The list of spectroscopic targets is distributed among a series of spectroscopic {\\it tiles} of $3^\\circ$ diameter (Blanton \\etal\\ 2003) to maximize observing efficiency.  Each tile then forms the design for a spectroscopic plate: holes are drilled in aluminum plates corresponding to the position of each object for which spectra will be measured.  At the telescope, optical fibers feeding a pair of double spectrographs are plugged into each plate; the spectroscopic observations, carried out under conditions not pristine enough for imaging, are typically 45 minutes per plate.  The spectra are wavelength- and flux-calibrated, and run through an automatic pipeline to classify them and determine redshifts (EDR paper).  The previous data release papers describe the quality of the data; the basic attributes of the data are given in Table~\\ref{tab:characteristics}. ",
        "conclusions": ""
    },
    "0410/astro-ph0410525_arXiv.txt": {
        "abstract": "Live \\fe60 has recently been detected in a deep-ocean ferromanganese crust, isolated in layers dating from about 3 Myr ago. Since \\fe60 has a mean life of 2.2~Myr, a near-Earth supernova is the only likely source for such a signal, and we explore here the consequences of a supernova origin. We combine the \\fe60 data with several supernova nucleosynthesis models to calculate the supernova distance as a function of progenitor mass, finding an allowed range of $15-120$ pc. We also predict the signals expected for several other radioisotopes, which are independent of the supernova distance. Species likely to be present near or above background levels are \\be10, \\al26, \\mn53, \\hf182 and \\pu244. Of these, \\hf182 and \\pu244 are nearly background-free, presenting the best opportunities to provide strong confirmation of the supernova origin of the \\fe60 signal, and to demonstrate that at least some supernovae are the source for the {\\em r}-process. The accuracies of our predictions are hampered by large uncertainties in the predicted \\fe60 yields for supernovae of different masses, so the new crust data motivate a redoubled theoretical attack on this problem. ",
        "introduction": "The possibility of finding terrestrial signatures of near-earth supernova explosions has been intensively discussed over the last decade \\citep[see, e.g.,][]{efs,fe}. Recent evidence for a supernova remnant $12^\\circ$ in diameter, that may be associated with \\al26 $\\gamma$-rays, points to an explosion within the past few Myr and within 100--200 pc \\citep{antlia}. Moreover, the existence of the Local Bubble -- containing hot, rarefied gas in which the Sun resides \\citep{frisch} -- suggests that we live in the {\\em interior} of one or more supernova remnants \\citep[e.g.,][]{sc}. This enhances the prospects of identifying terrestrial evidence of supernova remnant engulfing the Solar System and the Earth.  Even if it were too far away to be hazardous for life on Earth \\citep[see, e.g.,][]{ruderman,ES}, a sufficiently close supernova explosion would deposit on the Earth supernova ejecta and related material produced by cosmic rays \\citep{efs}. This could have resulted in the possible uptake of supernova-produced elements in deep-ocean sediments or manganese-iron crusts. Following \\citet{efs}, \\citet{knie99} discovered the first evidence for live \\fe60 in a deep-ocean ferromanganese crust, at a level $\\sim 100$ times the expected background and consistent with an event in the last $\\sim 5$ Myr at a distance $\\sim 30$ pc \\citep{fe}. However, the structure and geometry of the crust made it difficult to reconstruct a reliable time history, and thus the nature of the \\fe60 signal remained uncertain.  Very recently, \\citet{knie} have studied a different ferromanganese crust, which is better suited for a high-precision study with good time resolution. They find a \\fe60 spike far above background in a crust layer which formed between 2.4 and 3.2 Myr ago. This detection confirms the earlier result, but now goes beyond detecting the presence of \\fe60 to revealing the details of its deposition history. The presence of a live radioisotopic signal in a single, isolated layer agrees with the straightforward expectations of the supernova hypothesis \\citep{efs,fe}. These data greatly strengthen the case that supernova ejecta were deposited on the Earth about 3 Myr ago.  This is a major new result, and one that finally gives a strong direct empirical basis for longstanding speculation and arguments based on indirect evidence regarding the possibility of near-Earth supernova explosions in the geological past.  In light of the \\citet{knie} result, the time is ripe to take the supernova hypothesis seriously and investigate its consequences quantitatively. Specifically, in this paper we show how ferromanganese crusts can be used as telescopes to probe directly the {\\em r}-process in supernovae. ",
        "conclusions": "\\citet{knie} measured \\fe60 and \\mn53 (mean life $\\tau_{53} = 5.3$ Myr) throughout the ferromanganese crust 237KD, using accelerator mass spectrometry techniques. Over the 28 layers studied, just 69 live \\fe60 atoms were detected. This signal was found uniquely in layers corresponding to a time interval~\\footnote{ The crust was analyzed with two different time resolutions but even the smaller timespan of 0.3~Myr in the finer resolution analysis is expected to be an artifact representing an upper limit on the physical deposition timescale, which is likely to be of order $\\sim 10$~kyr.} % \\beq t = 2.4-3.2 \\ {\\rm Myr} \\eeq before the present. The resulting isotopic fraction is \\beq \\frac{\\fe60}{{\\rm Fe}} = (1.9  \\pm 0.24) \\times 10^{-15}, \\eeq which is a highly significant result, despite its small magnitude, testifying to the sensitivity and precision of the experiment. The principal uncertainty in this result is due to instrumental background, essentially the contamination due to stable isobars.  Given the crust size and Fe composition, the isotopic fraction can be used to infer a surface density of \\fe60. To do this, one must also correct for radioactive decay and geometric effects, and the result is \\beq N_{60,{\\rm obs}} = 2.9 \\times 10^6  \\ {\\rm atoms/cm^2}. \\eeq In order to infer the incident fluence of \\fe60, this result must be corrected for geometry and an uptake factor. Assuming the terrestrial fallout is isotropic, a geometric factor 1/4 arises from the ratio of the Earth's full cross section ($\\pi R_\\oplus^2$) to its total surface area ($4\\pi R_\\oplus^2$). Also, the crust incorporates only a certain percentage of all deposited material incident upon it, and a so-called `uptake factor' $U < 1$ accounts for this effect \\citep[for further discussion, see][]{knie,knie99}, leading to \\beq \\label{eq:surf} N_i(t)=\\frac{U}{4}F_i\\exp(-t/\\tau). \\eeq The uptake factors are different for each sample and every element, and thus cannot be evaluated independently of the details of the sample. Consequently, in the following discussion we present results for $U = 1$, with the understanding that all the results presented must be reduced by the appropriate uptake factor, which may be inferred from studies of stable isotopes.  \\citet{knie} use Mn and Fe data to infer an iron uptake factor $U_{\\rm Fe} \\simeq 0.006$, from which they infer an incident fluence of \\beq F_{60,\\oplus} = 2.0 \\times 10^9  \\ {\\rm atoms/cm^2}. \\eeq This signal encodes key information about the supernova explosion, as we now see.     Following the prediction of \\citet{efs}, \\citet{knie99} found initial evidence for supernova \\fe60 in a ferromanganese crust, prompting \\citet{fe} to discuss other isotopic signatures of a nearby supernova. The exciting new result of \\citet{knie} confirms the earlier observation with higher significance and more precise dating, motivating this paper. Using the data of \\citet{knie} to estimate the local interstellar fluence of \\fe60 and combining this with the total mass of \\fe60 ejected by diverse supernovae \\citep[][]{rhhw,ww}, we have inferred a supernova distance range $D \\sim 10-100$ pc, consistent with \\citet{knie}.  The \\fe60 detection and nucleosynthesis models together allow us to make distance-independent predictions for other radioisotopes. For supernova-produced isotopes such as \\al26, \\mn53 and \\cl36, we find that yields of the first two of these isotopes are likely to be detectable above background. We also consider the {\\em r}-process elements \\pu244, \\hf182, and \\sm146, finding again that the first two are likely to exceed background. If the supernova responsible for the \\fe60 observations of \\citet{knie} was also a site for the {\\em r}-process, clear signals should be visible. This would finally provide direct proof that supernovae are sites for the {\\em r}-process.  The firm detection of \\fe60 also adds support to recent ideas regarding the origin of the Local Bubble. \\citet{sc} concluded that several events are needed to create the Local Bubble, the last occurring within the past 5 Myr in order to maintain the high interstellar temperature. \\citet{ma} argued that the need for multiple events points to correlated explosions within an OB association, and specifically suggested the Lower Centarus Crux (LCC) of the Sco-Cen association as the likely site. \\citet{bmc} have examined the possibility that this OB association could have been the site of the \\fe60 progenitor. From kinematic considerations they find that the closest approach of the LCC was about 100 pc, with considerable uncertainties. Taken at face value, this result is in the upper part of our range of distance estimates. This suggests that either the event was a distant one (and thus the \\fe60 yield was high), or that the LCC approach was closer than estimated, or that the LCC was not the event site.  The accuracies of our predictions for the fluences of these isotopes are limited by the wide variations in the \\fe60 yields found in calculations for different progenitor masses, and between different model calculations \\citep[][]{rhhw,ww}. Therefore, at the present stage our results give only approximate estimates of the expected orders of magnitude of the fluences of supernova isotopes. This underlines the great need for more theoretical studies of \\fe60 production in supernovae, able to reduce the calculational uncertainties. Such a reduction would be necessary before one could predict any biological effects that might be associated with the supernova responsible for the \\fe60 observations of \\citet{knie99} and \\citet{knie}.  Measurements using other ferromanganese crusts would be most welcome, and observations of other supernova radioisotopes have high priority. These would enable one to evade the uncertainties associated with the \\fe60 yield. Another helpful calibration point comes from the recent evidence \\citep{rhessi} for the $\\gamma$-ray line signal from the decay of interstellar \\fe60. This reflects the current Galactic \\fe60 inventory, and constrains \\fe60 nucleosynthesis averaged over all sources \\citep{prantzos}.  The important results of \\citet{knie99} and \\citet{knie} open up a new era of `deep-ocean crust astronomy'. After providing already the first observations of the terrestrial impact of a nearby supernova, their measurements now offer the prospect of verifying directly that the {\\em r}-process takes place, at least partly, in ejecta from supernovae."
    },
    "0410/astro-ph0410655_arXiv.txt": {
        "abstract": "We study the composition of the Galactic interstellar medium (ISM) toward the Galactic center region ($5^\\circ<|l|<20^\\circ$) by utilizing X-ray absorption features of three bright low-mass X-ray binaries (LMXBs), \\gxp, \\gxm, and \\gxz, observed with the \\Chandra\\ HETGS. We detect X-ray absorption fine structure (XAFS) of the Si K-edge, characterized by a narrow and a broad absorption feature at 1846 eV and $\\approx$1865 eV, respectively. Comparison with ground experimental data indicates that most of the ISM Si exists in the form of silicates, although a composition of ``pure'' forsterite is ruled out. The XAFS spectra of the sulfur K-edge indicate that a significant fraction of S exists in the gas phase. From each source, we derive the column densities of Mg, S, Si, and Fe from the K-edge depth and that of O (or H) from the absorption of the continuum. The elemental abundance ratios are found to be consistent between the three targets: the mean values of O/Si, Mg/Si, S/Si, and Fe/Si are determined to be 0.55$\\pm$0.17, 1.14$\\pm$0.13, 1.03$\\pm$0.12, and 0.97$\\pm$0.31 solar, respectively (90\\% error in the mean value). We discuss the origins of the overabundances of the heavy metals relative to O in the Galactic ISM by comparison with the abundance pattern of the intracluster medium in clusters of galaxies. Assuming that most of the Mg and Si atoms are depleted into silicates of either the proxine or olivine family, we estimate that the number ratio of Mg to Fe in olivine is $\\simgt$1.2 and that 17\\%--43\\% of the total O atoms in the ISM must be contained in silicate grains. ",
        "introduction": "The composition of the Galactic interstellar medium (ISM) contains key information for understanding the formation and evolution of the Galaxy. In particular, the elemental abundances of heavy metals are crucial since the ISM is in the last stage of chemical evolution of the Galactic disk, reflecting the history of metal enrichment by supernovae (SNe). In addition, it is of great importance to understand the chemical compositions of dust grains and the dust-to-gas ratio in the ISM. Dust grains determine the physical properties of the ISM and dynamics of star formation. For reviews of ISM abundances and dust grains in the Galaxy, see e.g., \\citet{sav96} and \\citet{dra03a}, respectively.  Many efforts have been made to determine reference (or ``cosmic'') abundances of the ISM (total of dust plus gas) by observations of stellar photospheres, such as the Sun, main sequence B-stars, and F/G stars, and \\ion{H}{2} regions, etc. However, all these measurements could have inherent problems. Usually, solar abundances, those measured from the photosphere of the Sun or from meteorites, are adopted as reference values \\citep{and89}, although the 4.5 Gyr old Sun may not truly represent the present ISM compositions. Since 1994 there have been discussions that the true ISM abundances are about 2/3 the solar values \\citep{sav96,sno96}, based on new measurements of averaged stellar abundances in B-stars \\citep{gie92} and metal-rich dwarf F/G stars \\citep{edv93} in the solar neighborhood. In B stars, it is suggested that element stratification due to diffusion is common: the measured abundances in the stellar surface do not always represent those of the ISM from which they form \\citep{hem03}. Although metal-rich F/G stars were expected to be young enough to reflect the recent ISM composition, \\citet{edv93} found that there is a large scatter in the age-metalicity relation of F/G stars, indicating that ``metal rich'' stars inevitably include old populations. In fact, \\citet{sof01} calculated averaged abundances of only the young ($\\tau <$2 Gyr) F and G disk stars from \\citet{edv93} and found that they become closer to the solar abundances. The abundances derived from collisionally excited lines in \\ion{H}{2} regions are subject to large systematic errors because of temperature fluctuations in the nebula \\citep[e.g.,][]{mat95,kin95}. It should be noted that there exists abundance gradients decreasing with Galactocentric distance (e.g., summarized in Table~1 of \\citealt*{chi01}), making the issue more complex.  Studies using UV absorption lines of interstellar gas by the Goddard High-Resolution Spectrograph onboard the {\\it Hubble Space Telescope} indicate that a large fraction of Mg, Si, and Fe are depleted into dust grains in the Galactic disk \\citep[see][]{sav96}. In fact, a large number of silicates are expected in the ISM in order to reproduce the extinction curve in the infrared band \\citep{dra84}. The measurement of abundance patterns in dust grains can tightly constrain the chemical composition of the dust. The dust abundances can be measured only indirectly, however, by subtracting the gas-phase abundances in the ISM from the total (gas+dust) cosmic abundances. This means that the results strongly depend on the cosmic abundances assumed. Dust models claim that the infrared 9 and 18 $\\mu$m absorption structure of silicates, due to the transitions of the bending and stretching modes of the S-O bond, is too strong to be explained by a sub-solar abundance of Si \\citep{sno96,mat96}. \\citet{mat98} found that the problem can be avoided by introducing ``fluffy'' dust grains, which is not supported by a recent X-ray observation of the dust-scattered halo around GX~13+1 \\citep*{smi02}, however. The missing Si (and C) problem will be significantly relaxed if the true ISM abundances are closer to solar \\citep{sof01}. Thus it is an urgent task to establish standard abundances for understanding the properties of the ISM.  High-resolution X-ray absorption spectroscopy of bright Galactic sources provides us with a powerful technique for constraining the properties of the ISM. Unlike UV absorption lines, one can measure the metal abundances in both the gas and the dust phase from the edge depth, leading to a direct determination of the ISM abundances based on simple physics. In addition, X-ray absorption fine structure (XAFS) can be used as a diagnostic of the chemical state of an element. Analysis of XAFS, including X-ray absorption near edge structure (XANES) and extended X-ray absorption fine structure, could constrain not only the dust-to-gas ratio but also the composition of dust material. Since the launch of \\chandra\\ in particular, which carries transmission grating spectrometers, extensive studies have been made using absorption structures of the oxygen and neon K-edge using bright X-ray binaries with hydrogen column densities of several $10^{21}$ cm$^{-2}$ (e.g., X~0614--091: \\citealt{pae01}; Cyg X-1: \\citealt{sch02}; Cyg X-2: \\citealt{tak02}; X Persei: \\citealt{cun04}; multiple targets: \\citealt*{jue04}), following the pioneering work with the {\\it Einstein} observatory by \\citet{sch86}. These sources are, however, not suitable for studies of heavier elements such as Mg, Si, S, and Fe, as they produce only shallow K-edge structures because of their low column densities.  In this paper, we investigate the ISM composition of Mg, Si, S, and Fe elements toward the Galactic center region by systematically analyzing the X-ray absorption features in the X-ray spectra of three persistent (neutron star) low-mass X-ray binaries (LMXBs), \\gxp, \\gxm, and \\gxz, which have absorption column densities of several $10^{22}$ cm$^{-2}$ \\citep[see e.g.,][]{chr97,asa00}. The excellent energy resolution of the \\chandra\\ HETGS \\citep{can00} gives us the best opportunities to study XAFS around the Mg to S K-edge. Combined with the absorption to the continuum, we also constrain the abundance ratios not only for these elements but also for O (and hence H, with a reasonable assumption for O/H).  As the three targets are located at $5^\\circ<|l|<20^\\circ$ within $|b|<1^{\\circ}$ and at likely distances of $\\simgt 7$ kpc from the Sun \\citep[see e.g., Table~4 of ][and references therein; summarized in our Table~1]{hom04}, we can study the averaged properties of the ISM in the Galactic disk at Galactocentric distances of 0.7--8.5 kpc. Recalling the good correlation between the location and column density of Galactic LMXBs, it is very likely that the majority of the absorption column toward these targets is attributable to the ISM, not to the circumstellar material intrinsic to the source. Furthermore, as we show below, we obtain similar abundance ratios consistently from the three targets, also supporting this idea. Thus, it is justified to assume that the contribution of circumstellar matter is negligible (this point is also discussed in \\S~4.2). The spectral model of the continuum from persistent LMXBs is well established to be a sum of a multi color disk (MCD) model plus a black body component \\citep{mit84}. Hence, there is little uncertainty in modeling of the interstellar absorption, which could otherwise be coupled to the unknown continuum shape. This is an advantage of using a persistent neutron-star LMXB as a background source. \\gxm\\ and \\gxz\\ are Z-sources, while \\gxp\\ is often called a hybrid source as it shares characteristics of both a Z and an atoll source \\citep{has89}. Table~1 summarizes the mean radio flux density, $K$ band magnitude, and column density toward the source estimated from our X-ray spectral fit (see below) for each target; because of the large extinction in the visible band no optical counterparts have been identified yet.  The paper is organized as follows. \\S~2 describes the observations and data reduction. In \\S~3 we present the analysis method and results on the XAFS and abundances. Implications of our results are discussed in \\S~4, and the conclusion is given in \\S~5. The solar and ``ISM'' abundances referred to in this paper correspond to the values by \\citet{and89} and by \\citet*{wil00}, respectively: the O, Mg, Si, S, and Fe abundances relative to H are $8.51\\times10^{-4}$, $3.80\\times10^{-5}$, $3.55\\times10^{-5}$, $1.62\\times10^{-5}$, and $4.68\\times10^{-5}$ in solar abundances, and $4.90\\times10^{-4}$, $2.51\\times10^{-5}$, $1.86\\times10^{-5}$, $1.23\\times10^{-5}$, and $2.69\\times10^{-5}$ in ``ISM'' abundances.\\footnote{These solar and ``ISM'' abundances are available with the ``abund angr'' and ``abund wilm'' command on XSPEC (version 11.2.0), respectively.} ",
        "conclusions": "We have systematically analyzed high resolution X-ray absorption spectra of three Galactic bright sources to determine the properties of the ISM toward the direction of the Galactic center within $(l, b) < (20^\\circ, \\pm 1^\\circ)$. The conclusions are summarized as follows.  1. The silicon K-edge XAFS is characterized by a narrow absorption feature at 1846 eV and a broad one around $\\approx$1865 eV. Comparison with ground experimental data indicates that most of the ISM Si exists as silicates. The energy of the narrow absorption line rules out a composition of ``pure'' forsterite.  2. The sulfur K-edge XAFS consists of an edge feature at 2474--2490 eV and a narrow absorption line at 2469.4$\\pm$0.7 eV, indicating that a significant fraction of S exists in the gas phase. A partial contribution from iron sulfides and/or solid sulfur is possible, however, which may be the dominant form of S trapped onto dust grains in molecular clouds.  3. The Mg K-edge is detected at 1307$\\pm$2 eV, possibly with complex, broad absorption features in the 1310--1325 eV range. These features are consistent with that expected from magnesium silicates, implying that Mg atoms in the ISM are mainly contained in silicates.  4. From the K-edge depth of each element and the continuum absorption, we have determined the O/Si, Mg/Si, S/Si, and Fe/Si abundance ratio to be 0.55$\\pm$0.17, 1.14$\\pm$0.13, 1.03$\\pm$0.12, and 0.97$\\pm$0.31 solar, respectively. Similar metal abundance patterns are observed in very metal rich stars with [Fe/H] $\\simeq$ 0.3--0.5 in the Galaxy and also from the ICM in the central part of clusters (or groups) of galaxies. Considering the fact that the Galactic bulge is an old system similar to cD galaxies, we infer that our result may be explained by a significant contribution from SNe~Ia that produce a high Si/Fe abundance ratio superposed on that from SNe~II responsible for the production of O and Mg.  5. Assuming that most of the Mg and Si atoms are depleted into silicates of either the proxine or olivine family, we estimate that the number ratio of Mg to Fe in olivine is $\\simgt$1.2, and that 17--43\\% of total O atoms in the ISM must be contained in silicate grains."
    },
    "0410/astro-ph0410463_arXiv.txt": {
        "abstract": "We present multicolour imaging for a sample of 33 dwarf and intermediate-luminosity galaxies in the field of the \\ngc\\ Group, complemented with mid-resolution spectroscopy for a subsample of 13 objects. With these data, a revised membership and morphological classification is made for the galaxies in the sample. We were able to confirm all but one of the ``definite members'' included in the spectroscopic subsample, which were originally classified based on morphological criteria; however, an important fraction of background galaxies is probably present among ``likely'' and ``possible'' members.  The presence of a nucleus could be detected in just five out of the nine galaxies originally classified as dE,N, thus confirming the intrisic difficulty of photographic-plate morphological classification for this kind of objects.  Our deep surface photometry provided clear evidences for disc structure in at least three galaxies previously catalogued as dE or dS0. Their transition-type properties are also evident from the colour-magnitude diagram, where they lie near the late-type galaxies locus, suggesting an evolutionary connection between a parent disc-galaxy population and at least part of present-day dEs.  Half a dozen new dSph candidates were also found, most of them at small projected distances from \\ngc, the central galaxy of the Group.  The \\ngc\\ Group appears clearly defined in redshift space, with a mean heliocentric radial velocity, $\\langle v_r \\rangle = 2461 \\pm 84 \\kms$ ($z = 0.0082$), and a moderate dispersion, $\\sigma_{v_r} = 431 \\kms$.  Our kinematical data show no luminosity segregation for early-type galaxies: both dwarf and bright E/S0 systems show very similar velocity distributions ($\\sigma_{v_r} \\sim 290~\\kms$), in contrast to late-type galaxies that seem to display a broader distribution ($\\sigma_{v_r} \\sim 680~\\kms$). ",
        "introduction": "}  A systematic observational work to map the low surface brightness (LSB) galaxy distribution in selected zones of the sky has been carried out by different teams, leading to complete surveys and morphological catalogues of some loose groups of galaxies and nearby clusters (\\citealp*{BST85, KKB85, IWO86}; \\citealp{DPCDK88}; \\citealp*{F89, FS90, JD97, SH97}).  These catalogues serve as a basic reference to any further analysis relying on accurate photometry or spectroscopy of individual galaxies \\citep*[e.g.,][]{BM88, BH91, CFG94, HM94, SHP97, C99}.  However, accurate observations of LSB galaxies still remain a quite difficult task both for photometry and spectroscopy, especially for clusters and groups of galaxies beyond the Virgo and Fornax Clusters.  These two nearby clusters are the environments where, by far, most detailed studies of the dwarf and low-luminosity galaxy population have been carried out till present, including recent spectroscopic and photometric surveys using wide-field detectors \\citep[e.g.,][]{HKRIQ99,HIVKR99,DPJG00,KDSBH00, DBP02, SDS03}.  A complementary view can be gained through the study of smaller groups. These have the advantage that even a moderate-sized sample would be fairly representative of the whole group population. At the same time, depth effects would be minimized allowing a better analysis of distance-depending quantities.  One of these rather loose groups appears surrounding the elliptical galaxy \\ngc. The \\ngc\\ Group was catalogued by \\citet[hereafter FS90]{FS90}, who list 162 galaxies within $\\sim 45$ arcmin of \\ngc, about 80 percent of which are dwarfs ($B_{\\rmn T} \\ge 16$ mag).  Surface $BV$ photometry for a small sample of dwarf and intermediate luminosity (mostly) elliptical galaxies in the \\ngc\\ Group (of which 6 objects in common with the present work) has been presented in \\citet[hereafter C99]{C99}, while \\citet{KRP04} completed these results with a morphological study of a brighter sample of member galaxies, based on $BRI$ photometry. \\citet{CB01} discussed the properties of a few particular dwarfs, including a possible link between dEs and blue compact dwarfs (BCD).  A systematic H\\,\\textsc{i} radio survey of this field has recently been carried out by \\citet{MCK04} providing new hints for the low-surface brightness galaxy population of this loose group. In this paper we further extend the \\citet{CB01} analysis and present multicolour surface photometry for 33 dwarf and intermediate-luminosity galaxies (plus six likely unclassified new members) in the field of the \\ngc\\ Group, along with mid-resolution spectroscopy for a subsample of 13 objects.  Our main goal is to clearly discriminate between members and background objects, in order to study (in a future paper of this series) the stellar populations and structural properties of a sample with as little background contamination as possible. At the same time, an analysis of the kinematical properties of the Group can be made.  We will arrange our discussion by presenting first, in Sec.~2, the observational input. Surface photometry for each galaxy in our sample is obtained in Sec.~3, allowing us to carry out a fully consistent morphological (re)-classification for each object and assess --in some cases with the combined help of spectroscopy-- its fiducial membership to the \\ngc\\ Group. Nucleated and dwarf spheroidals are reviewed in some detail in this section, and we will also present fresh data for six new galaxies, possible Group members.  The main kinematical parameters for the Group are briefly summarized in Sec.~4, while the photometric and structural properties of the full galaxy sample are dealt with in Sec.~5. Here we discuss in particular the surface-brightness vs.\\ total magnitude relationship of the \\ngc\\ members and pick up some notable candidates to transition-type objects. Given their relevance in the long-standing debate about the evolutionary connection between different dwarf-galaxy types (i.e.\\ dE -- dI -- BCD), these special objects will be further characterized in Sec.~6 comparing with the global properties of the \\ngc\\ galaxy population. The main issues of our discussion and our conclusions will finally be summarized in Sec.~7.   \\begin{figure} \\includegraphics[width=\\hsize]{fields.eps} \\caption{DSS image of the \\ngc\\ Group, showing our observed fields. The frame is 90 arcmin on a side, North up, East to the left.} \\label{f_dss} \\end{figure} ",
        "conclusions": "In this paper we presented systematic multicolor photometry for an extended sample of 33 dwarf and intermediate-luminosity galaxies in the group of \\ngc\\ (including observations for the E galaxy \\ngc\\ itself, and the other luminous Sb member NGC~5054). For 13 of these objects, also mid-resolution spectroscopy was collected (this nearly doubles the galaxy sample covered in the literature for this group) in order to derive full kinematical information for each individual target and assess the dynamical status for the \\ngc\\ Group as a whole.  The Group appears clearly defined in redshift space, with a mean heliocentric radial velocity, $\\langle v_r \\rangle = 2461 \\pm 84 \\kms$ ($z = 0.0082$), and a moderate dispersion, $\\sigma_{v_r} = 431 \\kms$.  We also found marginal evidence of a possibly related  sub-structure (at least four galaxies spectroscopically confirmed) at $v_r \\sim 5000-6000 \\kms$, while three other galaxy aggregates seem to project on the background of the \\ngc\\ field, respectively at $z = 0.045, 0.09$ and 0.28.  Our kinematical data show no luminosity segregation among the early-type galaxy sub-samples: both the dwarf and bright E/S0 populations show nearly identical velocity distributions ($\\sigma_{vr} \\sim 294~\\kms$ for dwarfs and $\\sigma_{vr} \\sim 287~\\kms$ for bright ellipticals), while late-type galaxy distribution is on the contrary \\textbf{sensibly} broader, with $\\sigma_{vr} \\sim 680~\\kms$.  On the basis of the \\griz imagery, and thanks to excellent seeing conditions of our observations, we tried a revised morphological and membership classification for the galaxies in the sample. We were able to confirm all but one (i.e.\\ galaxy N109) of the ``definite members'' included in the spectroscopic subsample, which were originally classified based on morphological criteria; however, an important fraction of background galaxies is probably present among ``likely'' and ``possible'' members.  The presence of a nucleus, defined as a pointlike source clearly standing out from the inwards extrapolation of the galaxy profile, could be detected in just five out of the nine galaxies originally classified as dE,N, thus confirming the intrisic difficulty of nuclei identification on photographic plates.  Deep surface photometry down to $\\mu(g) \\sim 27 \\msq$ allowed us to detect six new dSph candidates \\citepalias[plus two other originally reported by][and confirmed here]{C99}, most of them at small projected distances from \\ngc, the central galaxy of the Group.  Clear evidence for disc structures in at least three galaxies previously catalogued as dE or dS0 (namely, N153, N156 and N42), was also obtained. The fact that this was inferred through different kinds of evidences, not always simultaneously present in all objects, suggests that these galaxies probably conform an heterogeneous set. It is thus necessary to extend this study, in order to shed some light on the evolutionary scenarios that led to the present-day dE population in groups.  The \\ngc\\ Group is particularly suited since it has about half the members as the more studied Fornax Cluster, however its central density and velocity dispersion are larger. It would be thus interesting to test whether proposed models to transform discs into spheroids, like galaxy harassment, could work in different environments.  In a second paper of this series (Buzzoni \\& Cellone 2004, in preparation) we will further extend our study of the \\ngc\\ Group trying a more specific analysis of the distinctive properties and the evolutionary status of stellar populations in LSB galaxies of this Group."
    },
    "0410/astro-ph0410133_arXiv.txt": {
        "abstract": " ",
        "introduction": "We are certainly living very promising days for cosmology. In a few years we have witnessed an explosion of cosmological data recollection coming from telescope, balloon and satellite experiments looking at the large scale structure distribution of galaxies as well as cosmic microwave background radiation anisotropies \\cite {review1}. In particular projects like the Wilkinson Microwave Anisotropy Probe WMAP \\cite {wmap} and the Sloan Digital Sky Survey SDSS \\cite {sdss}, in spite of having yet to complete their full schedule of observations, have already provided a significant amount of data which has been used, together with some reasonable assumptions, to determine cosmological parameters in a variety of ways. In particular this should eventually make possible to discriminate among several competing models of inflation.  Here we study a supergravity model of inflation essentially depending on one free parameter which can be identified with the slope of the potential at the origin and an energy scale fixed by the amplitude of the anisotropies, of the order of magnitude of the supersymmetry breaking scale $10^{11} GeV$. The model naturally gives a $\\it{total}$ number of e-folds which is typically close to the required number for observable inflation. The model can also accommodate a spectral index bigger or smaller than one without extreme fine tuning. We show that it is possible to obtain reasonable numbers for cosmological parameters and, as an example, we reproduce values obtained recently by Tegmark $\\it{et. al.}$ \\cite {tegmark}, particularly the data provided in the sixth column of Table 4 of that paper, where cosmological parameters are determined from a six-parameter model using WMAP and SDSS data alone. We choose this sample set of parameters for simplicity.  The model described here actually defines a new type of models where the inflaton rolls at high energy from negative values and a positive curvature potential. There is no inflation yet, however at some point defined by the equation $\\eta\\equiv V''/V=1$ inflation starts. The potential curvature eventually changes to negative values and inflation ends when $\\eta=-1$, (see Fig.$1$). No spontaneous symmetry breaking type mechanism for inflation of the new type to occur is here required. As a consequence the $\\it{total}$ number of e-folds is bounded independently of the initial conditions for the inflaton and close to the required number for observable inflation, see Table $1$. The spectral index can take values bigger than one at the expense of increasing the slope of the potential and thus decreasing the total number of e-folds.  In section $2$ we give a complete description of the model and obtain some general expressions related to inflation in the slow-roll approximation. Section $3$ presents the main results in the form of a table. We give a brief description of the ways the essentially unique free parameter of the model is fixed and calculate all other relevant quantities. Finally we conclude in section 4 briefly discussing the main results. ",
        "conclusions": "\\label{con}  We have studied a very simple model of inflation which essentially depends on one free parameter denoted by $c$ which controls the shape of the potential for small $\\phi$ and thus the dynamics of inflation. The large-$|\\phi|$ behavior, being dictated by the exponential term of the potential, is insensitive to $c$. In this sense characteristics of inflation are independent of the initial conditions for the inflaton.  The parameter $c$ can be identified with the slope of the potential at the origin. For $c=0$ we have the flat case \\cite {w}, \\cite {ross-s} where the slope and the curvature both vanish at $\\phi=0$. In this case the $\\it{total}$ number of e-folds grows without limit as we approach the origin and the spectral index can only take values lower than one. For negative $c$, however, the slope of the potential during inflation never vanishes and the amount of inflation is bounded, i.e., the total number of e-folds is bounded. Finite negative $c$ values actually define a new type of models where the would be inflaton rolls at high energy from negative values and a positive curvature potential without inflating. At some point defined by the equation $\\eta=1$ inflation starts. The potential curvature eventually changes to negative values and inflation ends when $\\eta=-1$ (see Fig.$1$). Here the inflaton rolls all the way from high energy down to its global minimum. No pre-inflationary epoch where the inflaton is kept waiting at the origin for spontaneous symmetry breaking to occur is here required. The spectral index can take values even bigger than one at the expense of increasing the slope of the potential and thus decreasing the $\\it{total}$ number of e-folds. We illustrate this in the second and third rows of Table 1. We also show that it is possible to obtain reasonable numbers for cosmological parameters and, as an example, we reproduce values obtained recently by Tegmark $\\it{et. al.}$ \\cite {tegmark}, where cosmological parameters are determined from a six-parameter model using WMAP and SDSS data alone. We see that the results are in excellent agreement. Also, the reheat temperature is sufficiently low to avoid overproduction of unwanted states such as gravitinos. The $\\Lambda$ scale introduced in the superpotential can be related to supersymmetry breaking in the hidden sector being ${\\cal O}(10^{11} GeV)$ and the resulting scale of inflation $\\Delta\\equiv V(\\phi_*)^{1/4} \\approx (\\Lambda M)^{1/2}$ is close to the unification scale. In conclusion the good agreement with the data in such a simple model indicate that perhaps some of the underlying assumptions should be part of a more elaborated model coming from a more fundamental theory."
    },
    "0410/astro-ph0410419_arXiv.txt": {
        "abstract": "We use simulations of large-scale structure formation to study the build-up of magnetic fields (MFs) in the intergalactic medium.  Our basic assumption is that cosmological MFs grow in a magnetohydrodynamical (MHD) amplification process driven by structure formation out of a magnetic seed field present at high redshift. This approach is motivated by previous simulations of the MFs in galaxy clusters which, under the same hypothesis that we adopt here, succeeded in reproducing Faraday rotation measurements (RMs) in clusters of galaxies.  Our $\\Lambda$CDM initial conditions for the dark matter density fluctuations have been statistically constrained by the observed large-scale density field within a sphere of 110 Mpc around the Milky Way, based on the IRAS 1.2-Jy all-sky redshift survey.  As a result, the positions and masses of prominent galaxy clusters in our simulation coincide closely with their real counterparts in the Local Universe.  We find excellent agreement between RMs of our simulated galaxy clusters and observational data.  The improved numerical resolution of our simulations compared to previous work also allows us to study the MF in large-scale filaments, sheets and voids. By tracing the propagation of ultra high energy (UHE) protons in the simulated MF we construct full-sky maps of expected deflection angles of protons with arrival energies $E=10^{20}~\\eV$ and $ 4 \\times 10^{19}~\\eV$, respectively. Accounting only for the structures within 110 Mpc, we find that strong deflections are only produced if UHE protons cross galaxy clusters. The total area on the sky covered by these structures is however very small. Over still larger distances, multiple crossings of sheets and filaments may give rise to noticeable deflections over a significant fraction of the sky; the exact amount and angular distribution depends on the model adopted for the magnetic seed field. Based on our results we argue that over a large fraction of the sky the deflections are likely to remain smaller than the present experimental angular sensitivity.  Therefore, we conclude that forthcoming air shower experiments should be able to locate sources of UHE protons and shed more light on the nature of cosmological MFs. ",
        "introduction": "In spite of decades of experimental and theoretical research on the subject, the nature and origin of Ultra High Energy Cosmic Rays (UHECRs) are still unknown. The experimental study of Extended Air Showers (EAS) produced by the interaction of UHECRs with the atmosphere provided, so far, only weak constraints on the composition of the primary particles. One of the few established facts is that the primaries of EAS are not standard-model neutrinos, since in this case showers would develop too deep in the atmosphere.  Somewhat less certain are the constraints on the allowed fractions of heavy nuclei and photons.  The search of a possible asymmetry between vertical and inclined showers made by the Haverah Park collaboration \\citep{HavPar..photons} allowed the conclusion that at energies above $4 \\times 10^{19}~\\eV$ less than $50\\%$ of EAS primaries are photons. Concerning the fraction of nuclei, the analysis of the inclined events recorded by the Fly-Eye shower detector \\citep{FlyEye..comp} seems to favor a light composition above $10^{19}~\\eV$. Similar results concerning photonic and heavy nuclei fractions were obtained by AGASA \\citep{Shinozaki:2003vv} using muon density in air showers as a primary mass estimator. It is therefore plausible that a significant component of UHECRs is composed of protons or light nuclei.  Besides determining the composition and energy spectrum of UHECR primaries, one of the main goals of EAS experiments is to map the angular distribution of cosmic rays. The angular resolution of these experiments lies generally in the range $0.5-2$ degrees, and may thus in principle allow the identification of UHECR sources with astrophysical counterparts observed in some other channel (optical, radio, X-rays, gamma-rays).  While the arrival directions of neutral particles would directly point to the sources (gravitational lensing is negligible), protons and composite nuclei may be deflected by galactic and Extragalactic Magnetic Fields (EGMFs). Note that at extremely high energies, $E \\simgeq 10^{20}~\\eV$, UHE protons are not expected to be significantly deflected by the {\\em galactic} magnetic fields.  Even at lower energies, $E \\simgeq 4\\times 10^{19}~\\eV$, the turbulent component of the MF in the Milky Way does not produce sizable deflections, while the large-scale component gives rise to ordered deflections, such that a reconstruction of the original arrival directions may still be possible~\\citep{Tinyakov:2001ir, Tinyakov:2003nu}. However, the very attractive perspective to do charged-particle astronomy with UHECRs might be spoiled by the presence of strong EGMFs.  So far, direct evidence for the presence of EGMFs has been found only in galaxy clusters \\citep[see e.g.][a short review of EGMF observations will be also given in section \\ref{obs}]{Carilli.Taylor..2002}. The origin of these fields is still unknown. One possibility is that the Intra Cluster Medium (ICM) has been magnetized at low redshift by the ejecta of local sources. If this is the case, EGMFs should be largely confined within galaxy clusters and groups, so that significant deflections of charged UHECRs would have to be expected only if they cross these structures.  Another possibility is that the seeds of the MFs observed in galaxy clusters have originated at high redshift, before the clusters collapsed and formed gravitationally bound systems. We will discuss possible origins of the intra-cluster magnetic fields in more detail in Section \\ref{models}, and demonstrate in Section \\ref{results1} that this scenario -- which is the model studied in our work -- leads to cluster MFs which reproduce many observational facts. Note that in this case EGMFs should fill a large fraction of the volume of the universe, so that they might induce considerable deflections of UHECRs traveling over cosmological distances, provided the fields are strong enough. Clearly this is an issue which has to be investigated carefully.  Until recently, the possible effects of EGMFs on the propagation of UHECRs have been investigated by constructing toy models for the structure of the field. Typically, EGMFs were assumed to have a cellular structure, with a Kolmogorov power spectrum at scales smaller than the cell size, and a uniform correlation length and rms field strength. None of these assumptions, however, is well supported by observational or theoretical arguments.  Indeed, due to the large conductivity of the IGM, magnetic fields should follow the gas flow and become concentrated and amplified in the high density regions.  The assumption of a statistically homogeneous MF is hence not valid in a structured Universe on scales of interest for the problem of UHECR propagation.   In the last few years, several groups started to develop physically more realistic models based on numerical simulations \\citep{Ryu..1998, Sigl:2003ay, Dolag:2003ft, Dolag:2003ra}, combining the magneto-hydrodynamics (MHD) of the magnetized IGM with N-body simulations of the driving gravitational dynamics of the dark matter. Barring this common general approach, the few simulations performed so far differ however not only in their numerical methodology but also in several important aspects of the model assumptions employed, as well as in their final results.  The common problem faced by all of these attempts is that one wants to predict the observationally unknown strength and geometry of the magnetic field in low density environments (e.g. filaments and voids) based on a calibration of the underlying model against observations of magnetic fields in galaxy clusters.  This involves an extrapolation over several orders of magnitude in density into a regime where there is no data. Even having such uncertainties, it is nevertheless crucial to demonstrate that the model, in the first place, correctly reproduces available observational data on the cluster scale.  In a recent letter \\citep{Dolag:2003ra}, we presented first results of a novel simulation method for the magnetic structure of the Local Universe, where we achieved a number of improvements compared with previous works.  First, the initial conditions for the dark matter density fluctuations were constrained to reproduce, to good approximation, the positions and masses of the largest structures (galaxy clusters) observed around the Milky Way. This allows us to remove the ambiguity in the choice of observer position that exists in ordinary simulations, and to obtain the first simulated all-sky maps of expected UHE proton deflections in the magnetic large-scale structure around the Local Group. Since the simulation volume (a sphere with radius $\\sim 110\\, \\Mpc$ embedded in a box of $\\sim 340\\ \\Mpc$ on a side) is much larger than considered in previous work, even massive clusters are present in our simulation, allowing for a more accurate comparison with observations. Second, we used a Lagrangian simulation method with an adaptive resolution scheme, which provides for better spatial resolution where it is most needed, namely in the high density regions where the field becomes more tangled. This helps in simulating RMs in galaxy clusters accurately, and hence to establish a reliable baseline for comparing the MHD simulation results with observations. By tracing UHE proton propagation in the simulated magnetic structure obtained in this novel approach, we were then able to show that significant deflections are only expected when galaxy clusters are crossed. Since clusters cover only a very small fraction of the sky, UHECR astronomy should therefore be, in principle, possible.  In this paper, we present a more detailed analysis of the predicted structure of the magnetic field in galaxy clusters, filaments and voids, as well as UHE proton propagation in the simulated magnetic web.  Our simulations are very CPU demanding and so far we have made two runs with different initial magnetic seed field. The results of the first run were reported in \\citet{Dolag:2003ra}.  The analysis of the present paper is mainly based on results obtained in a second simulation, which has been improved in several respects. In this new run, the comoving intensity of the seed field \\footnote{We define it as a value which the MF seed would have at present if only the cosmological redshift is taken into account. With this definition $B_0$ does not depend upon initial time.} was chosen to be $B_0 = 2 \\times 10^{-12}$ G, a factor 5 smaller than in the previous simulation.  This seed field strength leads to a more reasonable match with the magnetic field strengths observed in galaxy clusters, whereas the first simulation was pushing toward the upper limit of field strengths allowed by the data. We changed the orientation of the uniform seed field in order to test a possible dependence of our results on this aspect of the initial conditions. We have also improved the MHD component of our code to account for the back-reaction of the magnetic field on the ionized gas, which in turn may affect the MF in the most dense regions. Our analysis takes into account energy losses of UHE protons on cosmic microwave background (CMB) photons, and we construct deflection maps of protons with an arrival energy of $4\\times 10^{19}~\\eV$ and $10^{20}~\\eV$.  This paper is structured as follows. In Section~2, we give an overview of observations of magnetic fields and the possible origin of magnetic fields in galaxy clusters, which provides the motivation for the model we explore.  Section~3 describes the performed simulations, the initial conditions used, and gives a brief description of the simulation code used. Section~4 discusses the predictions of our simulations for the strength and structure of magnetic fields in galaxy clusters and relates these results to observations of Faraday rotation and radio halos of clusters. Section~5 analyses the simulation results for the magnetic field in low density environments like filaments, sheets and voids. In Section~6, we discuss the implications of the simulated magnetic field for the propagation of UHE protons. Finally, we summarize our findings in Section~7. In that section we also compare our approach with that of Sigl et al. (2003,2004a) and discuss the possible reasons of our discrepant results. Throughout this work we assume a background cosmology with $H_0=0.7$, $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, and a baryonic fraction $\\Omega_b/\\Omega_m$ of 14\\%. ",
        "conclusions": "In this work we analyzed the results of a constrained cosmological MHD simulation of the Local Universe embedded within a periodic box of $\\simeq 343$ Mpc on a side. The high resolution region centered on the `Milky Way' covers a sphere with $\\simeq 115$ Mpc radius and is resolved by 50 million dark matter plus the same number of gas particles. The initial conditions where obtained using constraints from the IRAS 1.2-Jy galaxy survey complemented by a random realization of a $\\Lambda$CDM universe. As a result, the evolved simulation at $z=0$ reflects the large-scale structure observed in the real Universe quite well, making it possible to identify most of the prominent halos and structures found within the simulation with known galaxy clusters and superclusters of the Local Universe.  This setup removes the arbitrariness involved in choosing a suitable observer position in ordinary simulations when trying to estimate how much the observed UHECRs may be deflected by extra-galactic magnetic fields.   We extracted the most massive clusters within the high resolution region of the simulation, leading to a set of 16 clusters (volume limited sample) within a temperature range of $\\approx 3-8~{\\rm keV}$. Examining the properties of the magnetic field formed in these systems, we confirm findings from earlier work. Radial profiles of the magnetic field strength are similar to that of the gas density in the outer parts, but the central magnetic field value strongly scales with the cluster temperature. During cluster formation, the magnetic seed field is not only amplified by adiabatic compression but also by shear flows that drive magnetic induction, a process that is ultimately powered by anisotropic accretion and merging events.  The initial field geometry is wiped out completely by the violent cluster formation history, a result that is almost independent of the exact mechanism for generating the initial magnetic seed field, provided it is generated early enough, say before $z\\sim3$. This makes the strength of the comoving intensity of the seed field, $B_0$, essentially the only relevant free parameter of our model.   We constructed synthetic Faraday rotation measures from the simulated clusters and compared them to a variety of observational data. We demonstrated that a comoving seed field of $B_0 = 0.2 \\times 10^{-11}$ G reproduces the observed amplitude and the radial scaling of rotation measurements found in galaxy clusters very well.  The strong dependence of the cluster magnetic field on cluster temperature leads to correlations between other measurable physical quantities. The predictions of our simulations for these correlations can also be confronted with observations. We demonstrated a good agreement for the scaling of the X-ray surface brightness with the rotation measure. Using two different models for the relativistic electron population, we also demonstrated that the simulated clusters are able to account for the observed correlation between radio power and temperature.  In low density regions outside of clusters, such as filaments, sheets and voids, no significant non-linear structure formation takes place and the magnetic seed field evolves mainly due to adiabatic compression. While within these regions large scale flows can still reorient the magnetic field, it will remain correlated with the geometry of the initial seed field. The detailed structure of the field in these regions therefore depends strongly on the properties of the seed, and thereby ultimately on the mechanism which creates the seed field in the first place.   We calculated the expected deflections of ultra high energy cosmic rays by the magnetic field in the large-scale structure of the simulated Local Universe. Given that the simulation reproduces in a realistic way positions, masses and sky coverage of the prominent clusters and superclusters of the Local Universe, we have no freedom to choose the observer position. Instead, we expect that placing the observer at the position of the ``Milky Way'' will provide a good model for what one expects to see in the real Universe.   We have computed full sky maps of expected deflections for particles with an energy at the detector of $4\\times 10^{19}$eV and $1\\times 10^{20}$eV. These maps show that strong deflections occur only at crossing galaxy clusters. Filaments, or sheets seen in projection, can lead to some noticeable deflections, but UHECRs approaching from most areas of the sky will not suffer significant deflections by the extragalactic magnetic field.  For UHECRs with arrival energies of $1\\times 10^{20}$eV, we also demonstrated that, if energy losses are included, the deflection will be reduced, as the traveled distance will decrease (for UHECRs with the same injection energy) and parts of the distance will be traveled at higher energies. UHECRs with energies of $4\\times 10^{19}$eV could in principal originate from sources up to a distance of 500 Mpc, which cannot be directly probed by our simulation. We therefore extrapolated the distribution of deflection angles for these distances, still finding that most of the sky should be covered by areas with small deflections.  Using another earlier simulation, where we chose a five times larger initial seed field that pushes the magnetic field within galaxy clusters towards (or even slightly above) the upper limit of observationally allowed magnetic field strength, we verified that even in such a more extreme model a significant fraction of the sky is covered by deflections below or comparable to the resolution of current experiments. In this second simulation we also used a different initial direction of the magnetic seed field, allowing us to explicitly verify that the statistical distribution of deflection angles is insensitive to the direction of the initial seed field.  A major systematic uncertainty of our model is the structure of the initial magnetic seed field. However, we argue that our estimates for the typical deflection angles provide a robust upper limit.  This is because an initially homogeneous magnetic seed field is expected to lead to the maximum deflection for a given field strength.  Any tangled component in the seed field will reduce the deflection and therefore strengthen our final conclusion. Note also that our modeling neglects contributions to the intra-cluster magnetic field from local injection processes, related for example to galactic winds or AGNs, which probably occur also in the late phases of cluster formation. Such an additional contribution will increase the cluster magnetic field. If important and accounted for, this would then force us to lower the cosmological seed field in order to avoid exceeding the observational limits for the field in clusters. A side effect would then be a reduction of the field and of the deflection angles in low density regions, again strengthening our final conclusion.  In summary, we hence conclude that our model provides an upper limit for the expected deflections of UHE protons by the magnetic field embedded in the large-scale structure of the Local Universe.  Therefore our main conclusion is that the deflections of UHE protons with energies larger than $4\\times 10^{19}$ eV are not big enough to prevent the pointing of UHECR sources in a significant fraction of the sky. Charged particle astronomy should be possible. Note that for this purpose any deflection of the UHECRs within the primary sources or within the direct embedding environment -- like for sources embedded in galaxy clusters -- is irrelevant, as long as these structures are smaller than the angular resolution of the experiments, which is true even for galaxy clusters, barring the exception of a very small number of nearby clusters.   An opposite conclusion was recently reached by \\citet{2004astro.ph..1084S,Sigl:2003ay}, who on the basis of a different MHD simulation and differing physical assumptions found significant deflections of UHECRs, large enough to spoil charged particle astronomy.  Their large deflections are a consequence of the much larger EGMF found in their MHD simulation, not of their different approach adopted for modeling the propagation of cosmic ray primaries.  In fact, their field strength decreases more slowly than in our simulations when going from galaxy clusters to the field, so that strong MFs extend over a larger volume and give rise to larger protons deflections. This difference can be due to several factors. Those which we were able to identify are listed below: 1) spatial resolution of the numerical codes; 2) simulation volume; 3) observer position;   4) physical assumptions about seed field.  The last issue was studied recently in \\cite{Sigl:2004gi}. It was shown that the use of cosmological seed field, as opposed to Biermann battery process with assumed subsequent turbulent amplification (which is not really well defined) does bring the result closer to ours, but cannot account for all difference.  Regarding the third issue, we note that the average MFs strength and the mean proton deflection may be significantly increased even if the observer is placed in the void but nearby a highly magnetized regions, see e.g. Fig.~\\ref{fig:fig1}. Our simulations have clear advantage here. Using constrained simulations we remove the ambiguity in the choice of the observer position.  The first and second issues affect the result via different overall normalization of the original seed field. In both simulations the seed field is normalized to reproduce MF in the centers of clusters. If for some reason MF is undershoot in cluster centers in a given simulation, this will lead to an overestimation of MF in filaments and voids, and vise versa.  In the simulation of \\cite{2004astro.ph..1084S,Sigl:2003ay} the largest clusters have temperature of a few keV, corresponding to a mass significantly smaller than $10^{15}~M_\\odot$, while in our simulation we find clusters with temperature as large as $8~\\keV$ which is consistent with observations. Since the MF strength in the clusters is correlated with the cluster temperature \\citep{Dolag:2001}, the strength of ICMFs inferred from the RMs in real clusters cannot be used to normalize the MF in simulations in which the cluster population is not representative of the real one.  Numerical resolution affect our Lagrangian simulations and the Eulerian simulations of \\citet{2004astro.ph..1084S,Sigl:2003ay} in a similar way. Since we can resolve the internal structure of galaxy clusters much better than possible with fixed mesh codes, we are confident that our normalization to the observational data on clusters is reliable. (We note that the resolution of both codes is comparable in filaments.) Note also that the analysis of the RMs of polarized radio sources observed in these systems shows that the coherence length of ICMFs is typically $\\sim 10~\\kpc$, comparable to the resolution we reach in clusters, but ten times smaller than the resolution reached in \\citet{2004astro.ph..1084S,Sigl:2003ay}.  It should be stressed that in our simulations the adiabatic compression and inductive amplification give roughly equal contributions into the final MF in clusters, see Fig.~\\ref{fig:b_ampli}. This is a non-trivial result which has to be verified with future work.  These discrepancies highlight the significant uncertainties that still exist in predicting the structure of extragalactic magnetic fields. Clearly, more work will be required to fully resolve this issue, which is of significant relevance for the prospects of forthcoming experiments. Further progress in this will probably be driven both by refined numerical methods as well as by ever more demanding comparisons with observational data.  Our present success in matching the observed correlations between the X-ray surface brightness and RMs in individual galaxy clusters, as well as the combined RMs profiles and radio halo luminosity is already highly encouraging in this respect."
    },
    "0410/astro-ph0410243_arXiv.txt": {
        "abstract": "The HESS collaboration recently reported highly significant detection of TeV $\\gamma$-rays coincident with Sgr~A$^\\ast$ \\citep{HESS}.  In the context of other Galactic Center (GC) observations, this points to the following scenario: In the extreme advection-dominated accretion flow (ADAF) regime of the GC black hole (BH), synchrotron radio/sub-mm emission of $\\sim 100 \\,\\rm MeV$ electrons emanates from an inefficiently radiating turbulent magnetized corona within 20 $R_{\\rm S}$ (Schwarzschild radii) of the GCBH.  These electrons are accelerated through second-order Fermi processes by MHD turbulence, as suggested by Liu et al.\\ (2004). Closer to the innermost stable orbit of the ADAF, instabilities and shocks within the flow inject power-law electrons through first-order Fermi acceleration to make synchrotron X-ray flares observed with Chandra, XMM, and INTEGRAL. A subrelativistic MHD wind subtending a $\\sim 1\\,\\rm sr$ cone with power $\\gtrsim 10^{37}\\,\\rm erg\\,s^{-1}$ is driven by the ADAF from the vicinity of the GCBH. As in pulsar powered plerions, electrons are accelerated at the wind termination shock, at $\\gtrsim 10^{16.5}\\,$cm from the GCBH,  and Compton-scatter the ADAF and the far infra-red (FIR) dust radiation to TeV energies. The synchrotron radiation of these electrons forms the quiescent X-ray source resolved by Chandra.  The radio counterpart of this TeV/X-ray plerion, formed when the injected electrons cool on timescales of $\\gtrsim 10^4\\,\\rm yrs$, could explain the origin of nonthermal radio emission in the {\\it pc}-scale bar of the radio nebula Sgr A West. ",
        "introduction": "HESS measurements of TeV fluxes from Sgr A$^\\ast$ \\citep{HESS} confirm, with high significance, earlier reports by the Whipple \\citep{Whipple} and CANGAROO \\citep{CANGAROO} collaborations (albeit at higher flux levels than HESS) about TeV emission from the Galactic Center. This result adds important new information about the massive ($M_{GC}\\approx 4\\times 10^6 M_\\odot$; \\citet{ghe03}) BH at the center of the Galaxy.  When combined with the wealth of observational data in the radio to X-ray wavebands \\citep{mf01}, the TeV $\\gamma$-ray data point to a scenario involving an inner, inefficiently radiating magnetized corona, i.e.\\ the ADAF, and a subrelativistic outflow of particles and field from the ADAF. As in pulsar plerions like Crab Nebula (see \\citet{kc84}), the wind terminates at a shock, at $R \\gtrsim 3\\times 10^{16}\\,\\rm cm$ from the GCBH, to create a black-hole powered `plerion' where the TeV radiation is produced.  The VLBI size of the radio source at the location of \\sgr is $110 (\\pm 60)\\,\\mu$as at $\\nu = 215 \\,\\rm ~GHz$ and $190 (\\pm 30)\\,\\mu$as at $\\nu = 87 \\,\\rm GHz$ \\citep{vlbi}. Using VLBA, \\citet{bow04} resolve  \\sgr  at 42 GHz to have an  angular size of $240 (\\pm  10)\\,\\rm \\mu$, or $R_{rad} \\simeq 24 \\,R_{\\rm S}$ for $M_{GC} = 4\\times 10^{6} M_\\odot$ black hole and a GC distance $d= 8\\,\\rm kpc$, and find that the size increases with wavelength. The radio emission of \\sgr is rather stable, with total luminosity $L_{36} = L_{rad}/ 10^{36} \\,\\rm erg \\, s^{-1}\\simeq 1$ and spectral flux $S_\\nu \\propto \\nu^{0.3}$ that sharply falls off above $\\nu_0 \\sim 10^{12}\\,\\rm Hz$ \\citep{Z95,fal98}.  The near-infrared (NIR) flares show factors-of-2 flux variations on timescales of $\\sim 20\\,$min and reach powers of $\\approx 10^{35}\\,\\rm erg\\,s^{-1}$ \\citep{gen03}.  The spectra rise in $\\nu S_\\nu$ diagrams, and show powers $\\approx 10^{35}\\,\\rm erg\\,s^{-1}$, similar with the ({\\it nonsimultaneously}) detected X-ray flares \\citep{bag01,por03}. The 200\\,s X-ray variability timescale limits the engine size to $r =R/R_{\\rm S} \\lesssim 5$. The unabsorbed quiescent flux detected by Chandra observatory from \\sgr is $L_{X.qui}(2$--$10\\, \\rm keV) \\approx 2.4 \\times 10^{33}\\,\\rm erg\\, s^{-1}$, and has angular extent $\\theta_{\\rm FWHM}\\approx 1.4^{\\prime\\prime}$ \\citep{Chandra}, corresponding to source size $R_X\\approx 8.3\\times 10^{16}\\,\\rm cm$.  The TeV spectral fluxes in two HESS campaigns in June/July and July/August of 2003 are almost exactly the same, extending with photon index $\\approx \\! 2.1$-2.3 from $160\\,$GeV to 10 TeV, with total power $L_\\gamma  \\approx 10^{35}\\,\\rm erg\\, s^{-1}$. There is no sign of flux variation within the statistics available in each of these sets \\citep{HESS}.  Given the high statistical significance of the signal in both data sets, these results suggest a quiescent or quasi-stationary TeV source coincident within $\\leq \\! 1.5^\\prime$ of Sgr A$^\\ast$.  We propose a model for Sgr A$^\\ast$ where the quasi-stationary radio and flaring X-ray/NIR emissions are synchrotron radiation from the ADAF. The MHD wind from the ADAF powers a black-hole plerion where the quiescent X-ray and TeV emissions are produced by electrons accelerated at the wind shock.  The model explains the existing multi-wavelength data, and makes predictions that can be checked by various X-ray and gamma-ray detectors. ",
        "conclusions": "A model consisting of magnetized coronal ADAF (synchrotron) radio emission within $\\approx 20~R_{\\rm S}$ of the GCBH, and of a black-hole plerion powered by a wind from the ADAF resolves many puzzling observations of Sgr A$^\\ast$.  X-ray flares are synchrotron radiation of electrons accelerated through first-order Fermi process by shocks within a few \\rs of the GCBH.  Electrons accelerated to $\\gamma\\gtrsim 10^8$ at the wind termination shock at $\\gtrsim 3\\times 10^{16}\\,$cm creates the GCBH plerion.  Its synchrotron X-ray emission has been resolved as an $\\simeq 1.4^{\\prime\\prime}$ source with Chandra \\citep{Chandra}.  The multi-TeV electrons Compton scatter the radio photons from \\sgr and FIR photons from the dust ring of Sgr A West to produce the TeV emission.  TeV emission from the GCBH plerion is nearly stationary because the cooling time of TeV electrons is $\\sim 100$ yrs.  A ``jet-ADAF\" model for \\sgr has been proposed by \\citet{fm00} and \\citet{ymf02}. Our model, though also based on energy outflow from the ADAF, differs greatly.  In particular, the quiescent and flaring X-ray components are produced at different sites.  The origin of the X-ray and NIR flares are explained as synchrotron emission of electrons accelerated during episodes of instabilities very close to the BH, not as Compton radiation in the SSC scenario of the ``jet-ADAF\" model. Importantly, the TeV flux cannot be easily explained if the observed X-rays were due to the SSC mechanism in the vicinity of the GCBH, as suggested by \\citet{fm00}.   Propagation of GeV electrons from the plerion on timescales of $\\gtrsim 10^4 \\,\\rm yr$ with speeds $\\sim 100 \\,\\rm km\\, s^{-1}$could significantly contribute to the radio synchrotron flux of Sgr A West. In particular, it could explain the suggested nonthermal origin of radiation in the bar of Sgr A West \\citep{bar}, formed in the outflow direction.  We also predict that quasi-stationary Compton and bremsstrahlung fluxes from the pc-scale plerion, coincident with the central parts of Sgr A West, will be significantly detected and possibly resolved with GLAST at GeV energies. But the expected highly variable Compton counterpart of the synchrotron X-ray flares from the GCBH vicinity is too weak to be detectable with GLAST or HESS.  The synchrotron extension of the flare emission in the hard X-ray/soft $\\gamma$-ray domain explains the $\\simeq 40\\,$min episode of a profound increase of flux detected by INTEGRAL from the direction of \\sgr \\citep{bel04}.  The absence of apparent TeV flux variations may suggest a proton origin for the TeV radiation \\citep{an04}. Indeed, proton and ion acceleration through first- and second-processes in the ADAF and through first-order acceleration at the wind termination shock could make cosmic rays to produce TeV emission through nuclear $pp$-interactions with $n\\sim 10^{3}\\,\\rm cm^{-3}$ dense gas on $pc$ scales, and could form extended TeV emission possibly already detected with HESS \\citep{HESS}.  The hadronic origin of the TeV radiation in the ADAF itself in the BH vicinity requires, however, an extremely dense gas target or extremely large proton powers $\\gtrsim 10^{39}\\,\\rm erg\\, s^{-1}$.   Finally, we note that the unidentified EGRET source 3EG J1746-2851 towards the GC \\citep{EGRET1} is significantly displaced \\citep{dh02} from the direction of the GCBH, and is unlikely to be related to Sgr A$^\\ast$. It is probably emission from a young pulsar, though not with the ``mouse\" PSR J1747-2958 \\citep{mc03}. A young pulsar with $\\gamma$-ray properties like Vela but with apparent $\\gamma$-ray power $\\approx 10 \\times$ larger could have been missed in pulsar surveys due to the large dispersion measure towards the GC. We suggest a deeper, higher radio-frequency and X-ray search at the refined location of 3EG J1746-2851."
    },
    "0410/astro-ph0410596_arXiv.txt": {
        "abstract": "We present Doppler images of the young K5V - K7V rapid rotator LO Peg from seven nights of continuous spectroscopy obtained in 1998 from July 04 to July 10. The images reveal the presence of a strong polar cap with appendages extending to mid-latitudes, but no starspots are seen below 15\\degr. We briefly discuss the distribution of spots in light of recent flux transport simulations which are able to reproduce the observed latitude dependence.  With the full timeseries of spectra, of which 314 are useful, many phases are observed three times over the seven nights of observations.  Using starspots as tracers of a solar-like latitudinal differential rotation in our image reconstructions, we find that the equatorial regions complete one more rotation than the polar regions every \\mbox{181 $\\pm$ 35} d. LO Peg is the second coolest star for which such a measurement has been made using indirect imaging methods. The degree of latitudinal shear is less than that seen in G and early K dwarfs, suggesting a trend in which differential rotation decreases with stellar mass in (pre-)main sequence objects. ",
        "introduction": "Several methods which enable us to measure the surface differential rotation in late-type stars with convective envelopes have appeared in the literature. Variability of stellar activity diagnostics such as the Ca {\\sc ii} H \\& K emission cores has been used to determine stellar rotation periods as exemplified by the long term Mount Wilson Survey \\citep{wilson78,baliunas95}. In many stars, the measured rotation period is found to vary periodically from one epoch to another over timescales similar to the 11 year solar magnetic activity cycle. If we adopt the solar paradigm, in which changes in latitude of the active regions occur throughout the cycle, we can attribute the measured period changes to a latitude dependent rotation in stars. The scatter in measured periods gives some indication of the degree of differential rotation shear on the star \\citep*{donahue96}.  A similar method based on photometric modulation due to starspots, in a sample of stars, has been found to yield a trend in the degree of surface shear with rotation rate \\citep{hall91dynamo,henry95diffrot}. Long term variations of the photometric period been shown to be indicative of magnetic activity cycles \\citep{messina02}. More recently, \\citet{reiners03diffrot} have found that main-sequence stars of spectral types F and G  exhibit high degrees of differential rotation. The rotational shear manifests itself in the Fourier transforms of the rotationally broadened line profiles. This allows precise measurements of shear from line profiles alone for stars with intermediate rotation rates.  A further method of determining the latitude dependence of rotation is to use spot distributions from indirect imaging techniques. Here we do not have to make assumptions about the latitude extent of starspots in order to determine the magnitude of shear between the equator and pole. We are however restricted to observing rapid rotators if we are to obtain sufficient resolution elements across a rotationally broadened profile to make this technique viable. The number of resolution elements is determined by a number of stellar broadening mechanisms, but is ultimately limited by the instrumental resolution. By cross-correlating constant latitude strips in Doppler images derived at closely related epochs, the latitude dependent surface shear can be determined \\citep{donati97doppler}. Alternatively, a differential rotation law (usually assumed to be solar-like) can be incorporated into the imaging model, and the whole data set from several nights combined to derive an image (\\citealt{donati00rxj1508}; \\citealt*{petit02}).   \\begin{table*} \\begin{minipage}{135mm} \\caption{Journal of observations of LO Peg and standards on 1998 July 04 - 10} \\protect\\label{tab:obs_journal} \\vspace{5mm} \\begin{center} \\begin{tabular}{lccccl} \\hline Object          & UT Start    \t& UT End      \t\t& Exp time [s]   \t& No. of frames& Comments\t\t\\\\ \\hline  & \t\t&               \t\t&  1998 Jul 04 &              \t&\t\t\t\\\\ \\vspace{-3mm} \\\\ Gl 673\t\t& 22:27\t& \t\t\t& 100\t\t& 3\t\t& K5V template          \\\\ LO Peg\t& 01:32\t& 05:57\t\t& 200\t\t& 49\t\t&\t\t\t\\\\ \\hline & \t\t&              \t \t& 1998 Jul 05\t&              \t&\t\t\t\\\\ \\vspace{-3mm} \\\\ Moon\t\t& 21:57\t& 22:07\t\t& 200, 30, 60   & 3\t\t& Solar template        \t\\\\ LO Peg\t& 01:46\t& 05:56\t\t& 200\t\t& 47\t\t&\t\t\t\\\\ \\hline & \t\t&               \t\t&  1998 Jul 06 &              \t&\t\t\t\\\\ \\vspace{-3mm} \\\\ Gl 687\t\t& 21:38\t& 21:53\t        \t& 300, 200, 200& 3\t\t& M3V template\t\\\\ Gl 649\t\t& 22:03\t& 22:17\t\t& 200\t\t& 3\t\t& M1V template\t\\\\ LO Peg\t& 01:40\t& 05:54\t\t& 200\t\t& 47\t\t&\t\t\t\\\\ \\hline & \t\t&               \t\t& 1998 Jul 07  &              \t&\t\t\t\\\\ \\vspace{-3mm} \\\\ LO Peg\t& 01:46\t& 06:01\t\t& 200\t\t& 48\t\t&\t\t\t\\\\ \\hline & \t\t&               \t\t& 1998 Jul 08\t&              \t&\t\t\t\\\\ \\vspace{-3mm} \\\\ LO Peg\t& 01:34\t& 05:59\t\t& 200\t\t& 50\t\t&\t\t\t\\\\ \\hline & \t\t&               \t\t& 1998 Jul 09  &              \t&\t\t\t\\\\ \\vspace{-3mm} \\\\ LO Peg\t& 01:51\t& 06:00\t\t& 200\t\t& 47\t\t&\t\t\t\\\\ \\hline & \t\t&               \t\t& 1998 Jul 10  &              \t&\t\t\t\\\\ \\vspace{-3mm} \\\\ Gl 447\t\t& 20:49\t& 21:23\t\t& 600\t\t& 3\t\t& M4V template \t\\\\ LO Peg\t& 01:41\t& 06:00\t\t& 200\t\t& 47\t\t&\t\t\t\\\\ \\hline  \\end{tabular} \\end{center} \\end{minipage} \\end{table*}   LO Peg (BD+22 4409) is among the least massive of the young rapid rotators in the solar neighbourhood. It was detected by the {\\em ROSAT WFC EUV} all-sky survey as the source RE J2131+23, and by the {\\em Extreme Ultraviolet Explorer} survey as the source EUVE J2131+23.3 \\citep{malina94}. \\citet{jeffries93kinematics} identified LO Peg as a member of the Local Association on the basis of its Galactic space motions and a large EUV to bolometric flux ratio. LO Peg was first studied in detail by \\citet{jeffries94bd22d4409} who determined an axial rotation period of \\hbox{0.42375 d} from V-band photometric observations. It was also noted that a slightly shorter period of \\hbox{0.38417 d} gave a physically acceptable solution with low false alarm probability and folded lightcurve with little scatter. Photometry yielded a visual magnitude of \\mbox{$V = 9.19 \\pm 0.05$} and colours \\mbox{$(B-V) = 1.08 \\pm 0.02$}, \\mbox{$(U-B) = 0.82 \\pm 0.08$}, \\mbox{$(V-R)_{KC} = 0.59 \\pm 0.02$} and \\mbox{$(V-I)_{KC} = 1.23 \\pm 0.03$}. Study of the Lithium 6708 \\AA~line yielded an abundance of N(Li) = {$1.30 \\pm 0.25$}, while the iron abundance [Fe/H] = \\mbox{$0.0 \\pm 0.1$}. The equatorial rotation velocity was determined to be \\vsini~= \\mbox{$69 \\pm 1$} \\kms. \\citet{jeffries94bd22d4409} also obtained a model fit to the spectrum of LO Peg, and found a spectral type of K5 - K7 with an age in excess of 30 Myr.  \\begin{figure*} \\begin{center}  \\begin{tabular}{cc} \\includegraphics[width=5.30cm,angle=0]{me759_fig01.ps} & \\includegraphics[width=5.30cm,angle=0]{me759_fig02.ps} \\\\ \\end{tabular} \\smallskip \\begin{tabular}{cc} \\includegraphics[width=5.30cm,angle=0]{me759_fig03.ps} & \\includegraphics[width=5.30cm,angle=0]{me759_fig04.ps} \\\\ \\end{tabular} \\smallskip \\begin{tabular}{ccc} \\includegraphics[width=5.30cm,angle=0]{me759_fig05.ps} & \\includegraphics[width=5.30cm,angle=0]{me759_fig06.ps} & \\includegraphics[width=5.30cm,angle=0]{me759_fig07.ps} \\\\ \\end{tabular}  \\end{center}  \\caption[LO Peg time series spectra ]{Time series spectra for 1998 July 04\\,-\\,10, with phase plotted against velocity. The mean profile has been subtracted from each time series. White features correspond to starspot signatures. } \\label{lopeg_spectra} \\end{figure*}    \\citet{jeffries94bd22d4409} also determined the axial inclination of LO Peg to be 50\\degs. By studying the \\ha\\ line, which appears in emission, these authors showed that LO Peg does not exhibit prominence features, which appear as \\ha\\ absorption transients in other similar stars as the cool optically thick clouds cross the stellar disc \\citep{cameron89cloud}. The lack of observational evidence indicates that if prominences do exist, they must be present only above low latitudes where they will not cross the stellar disc in the observer's line of sight. More recently, \\citet{eibe99bd} found that variability of the \\ha\\ line indicates strong downflows of absorbing material. A previous imaging study by \\citet*{lister99lopeg} has yielded surface Doppler images showing a polar cap and a low latitude band of features on the surface. Observations at two epochs allowed these authors to distinguish between the two proposed axial rotation periods discussed above, indicating that the longer 0.42375~d period is preferred.  The K5V\\,-K7V spectral type make LO Peg an important object because no other single stars of this mass have been studied. In fact, no stars between K3V and  and M1V have been imaged to date. At this spectral type, the radiative core decreases in size as the star approaches the fully convective regime at spectral type early to mid-M. If a solar-like interface dynamo is at work, according to \\citet{schussler96buoy} and \\citet{granzer00} we may expect only intermediate to high { latitude eruption of magnetic flux}. In this paper, we present images derived from data obtained over seven nights and determine the differential rotation for LO Peg. ",
        "conclusions": "\\subsection{Starspot distributions} With a spectral type of K5 - K7, LO Peg is one of the latest spectral type dwarfs which have been Doppler imaged. Additionally, it is the only main sequence star in the \\hbox{mid-late K} spectral type range to have been studied in this way. The evolution of starspots on the one-week timescale of the current observations is clearly quite small. Over the longer five year time base, the starspot morphology has also changed very little in terms of latitudinal distribution. The polar cap appears to have remained both stable and strong at both epochs of observation. This is intriguing because the slightly earlier spectral type (K3V), but more rapidly rotating \\hbox{(P = 0.38 d)} Speedy Mic showed no evidence for a polar cap \\citep{barnes01speedy} when observed at a single epoch in 1998. The early M dwarf stars, EK Dra (RE1816+541) and HK Aqr have similar periods to LO Peg but do not possess polar caps, at more than one epoch in the latter case \\citep*{barnes04hkaqr}. There is however evidence that the young K0V star AB Dor \\citep*{kurster94abdor} did not possess a polar cap in 1989, in which case these large spots may disappear periodically. If such behaviour is (quasi-)cyclic, it may be related to dynamo activity, although there is as yet no convincing evidence for this.  By simulating a star ten times more active than the Sun, \\citet{schrijver01polar} found that the higher rate of flux emergence, at around 30 times solar, resulted in the development of polar caps. But magnetic flux, as traced by starspots on LO Peg, clearly does not emerge at low latitudes in the way we see on the Sun. With the flux emergence patterns seen on young rapid rotators in mind, \\citet{mackay04} have investigated the role of meridional flows for stars where flux emerges up to latitude 70\\degs. In order to reproduce the intermingled polarity pattern as seen on the young rapid rotator, AB Dor, strong meridional flows of around 100 \\ms\\ are required.  Unfortunately we do not have magnetic maps of LO Peg which would allow us to compare our results more closely, but the the flux distribution patterns do demand comparison with the simulation maps. \\citet{mackay04} have shown that much of the flux piles up at high latitudes, with less at intermediate latitudes. In the case where the rate of flux emergence shows cyclic behaviour (with an 11 year cycle), but no change in the latitude at which flux erupts over that cycle, fewer spots are seen at low latitudes than intermediate and high latitudes. While the amount of flux at low latitudes varies throughout the cycle, features within a few degrees of the equator are present at all times. However when a solar-like variation of flux emergence with latitude, varying between $\\pm$70\\degs (rather than the solar $\\pm$40\\degs) to $\\pm$ 5\\degs, is simulated, there are times when no flux is present in the equatorial regions.  A distribution resembling this situation is certainly seen on the present images of LO Peg where we see no spots below 15\\degs. If observation can be reconciled with this picture of activity, long term Doppler imaging projects are required. It is worth noting however that while \\citet{baliunas95} found period variability, indicating the presence of  active regions at different latitudes and differential rotation, the younger stars in the Mount Wilson Survey tended to show no preferred period for magnetic cycles.  A solar type interface dynamo predicts only spots at intermediate latitude in rapid rotators (\\citealt{schussler96buoy} and \\citealt{granzer00}). { While spots {\\em do} appear at intermediate latitudes in both 1993 and 1998, there is some structure at lower latitudes, while the images are dominated by polar or circumpolar spots. In many cases however, in addition to polar and circumpolar structure, spots my be seen in bands at low latitudes.} This is the case for the young post-T Tauri star, RX J1508.6-4423 \\citep{donati00rxj1508} and Alpha Per G dwarfs \\citep{barnes01aper} for example. The role of meridional flows have been touched upon above and are known, in the solar case at least, to possess maximum amplitude at equatorial latitudes (e.g. \\citealt*{komm93meridional}). Flow rates following the solar pattern, but with greater amplitude, as modelled by \\citet{mackay04}, will then result in a lower chance of seeing starspots at low latitudes. Another likely reason for no spots being observed at low latitudes is related to the low axial inclination of LO Peg. At 45\\degs, spots in the equatorial latitudes of the star will appear severely foreshortened. If the S/N of the data is not sufficiently high, and only small spots exist, our imaging program may not find sufficient evidence for their presence.  \\subsection{Differential rotation} The differential rotation rate which we have measured for LO Peg is interesting in the context of other measurements. The value of equator-pole lap time, \\hbox{$\\delta$ =  \\mbox{$181 \\pm 35$} d}, lies on a trend (albeit sparsely populated) with spectral type. \\citet{donati00rxj1508} found \\hbox{$\\delta$ = 50 $\\pm$ 10 d} for the young G2 pre-main-sequence star RXJ 1508.6-4423, while \\citet{marsden04} found \\hbox{$\\delta$ = 66 $\\pm$} 14 d for the IC2602 G2 star, R58. Over a period of 8 years, \\citet{cameron02twist} found that the K0 dwarf AB Dor exhibits systematic changes in differential rotation rate from one season to the next, with \\hbox{$\\delta$ =  70 d to 140 d}. At approximately the same spectral type, the near-one-day period PZ Tel \\citep{barnes00pztel} exhibited an equator-pole lap time of 80 $\\pm$ 15 d. The latest spectral type star for which we have a differential rotation measurement \\citep{barnes04hkaqr} is the dM1-2 object, HK Aqr which exhibited an equator-pole lap time { formally consistent with zero} (i.e. between \\hbox{-1449 d and +448 d.}) Although the sense of the differential rotation was not clear, the magnitude is essentially very small.  LO Peg clearly fits into this trend of generally decreasing differential rotation as we descend the main sequence from early G through to early M. The changing AB Dor differential rotation however shows that there can be significant change, possibly dependent on the phase of a magnetic cycle. In all cases, the magnitude of differential rotation appears, to first order at least, to be independent of rotation rate. The Sun for instance has an axial rotation period of \\hbox{27 d} and exhibits an equator-pole lap time of 120 d, whereas this value falls only by a factor of two (twice the rate of differential rotation) for rapidly rotating objects with 1/50th the solar period. These two predictions of weak rotation dependence and a decrease with rotation rate are predicted by the fluid circulation models of \\citet{kitchatinov99drot}, although the exact magnitude of the predicted differential rotation does not match the observation. Additionally, for a \\hbox{1 d} axial rotation period K5V star, these models predict a poleward meridional flow rate of $\\sim$ 10 \\ms, an order of magnitude lower than those required by \\citet{mackay04} in their simulations. Clearly it would be desirable to determine an estimate of the magnitude and direction of flows observationally."
    },
    "0410/astro-ph0410630.txt": {
        "abstract": "We present here a mid-infrared imaging survey of 26 sites of water maser emission. Observations were obtained at the InfraRed Telescope Facility 3-m telescope with the University of Florida mid-infrared imager/spectrometer OSCIR, and the JPL mid-infrared camera MIRLIN. The main purpose of the survey was to explore the relationship between water masers and the massive star formation process. It is generally believed that water masers predominantly trace outflows and embedded massive stellar objects, but may also exist in circumstellar disks around young stars. We investigate each of these possibilities in light of our mid-infrared imaging. We find that mid-infrared emission seems to be more closely associated with water and OH maser emission than cm radio continuum emission from UC \\ion{H}{2} regions. We also find from the sample of sources in our survey that, like groups of methanol masers, both water and OH masers have a proclivity for grouping into linear or elongated distributions. We conclude that the vast majority of linearly distributed masers are not tracing circumstellar disks, but outflows and shocks instead. ",
        "introduction": "The discovery of the first astronomical maser was the hydroxyl (OH) maser by Weaver et al. (1965), followed shortly thereafter by the first detection of an astronomical water maser by Cheung et al. (1969). In the years following these major discoveries, astronomers found that these maser species were closely associated with phenomena related to the formation of massive stars, in particular they appeared to be directly associated with regions of cm radio continuum and far infrared emission. The extremely reddened far infrared colors of these regions as seen with IRAS seemed to confirm the idea that the masers are located in the hot, dusty environments of star formation containing copious amounts of radio emission. Masers, therefore, became known as signposts of massive star formation. However, advances in higher resolution radio imaging and accurate astrometry in recent decades have led to observations that show that not all masers are directly coincident with cm radio continuum emission. Genzel \\& Downes (1977) conducted the first water maser study toward known \\ion{H}{2} regions and discovered that while OH masers tend to be in or projected on \\ion{H}{2} regions, water masers were typically offset. Other observations followed (e.g. Forster \\& Caswell 1989; Hofner \\& Churchwell 1996), weakening the direct physical link between water masers and \\ion{H}{2} regions. If masers are close to, but not directly coincident with young massive stars, what phenomena or processes do these masers trace?  Masers in general are excited to emit from both radiative and collisional processes. In the star forming environment there are several possible processes and locations that have been suggested where masers can exist. It has been suggested (e.g., Elitzur 1992) that the cool, dense layer of gas between the ionization and shock fronts in the expanding \\ion{H}{2} regions around young massive stars may provide a habitable zone for masers (OH masers, in particular). The bulk motion and relatively high density of molecular material caught up in a well-collimated bipolar outflow or jet from a young star may, in principle, be a good location for maser emission (e.g., Torrelles et al. 1997). Even if masers are not taking part in the outflow from a young stellar object, the shock created by an outflow as it impinges on the ambient medium or on knots of material in the immediate vicinity of an outflowing star, also seem to be good locations for masers (water masers, in particular). The idea that masers are excited by embedded protostellar objects was originally suggested by Mezger and Robinson (1968) for OH masers. Specifically, these masers would exist in the accreting envelopes of massive protostars and be excited by the energy from accretion shocks at stages of early evolution \\textit{before} the onset of an \\ion{H}{2} or ``ultra-compact \\ion{H}{2}'' (UC \\ion{H}{2})  region (Forster \\& Caswell 2000). Water and OH masers have also been suggested to exist in the warm and dense environment of circumstellar disks (e.g., Torrelles et al. 2002, and references therein).  It is difficult to link any of these phenomena of massive star formation to the maser emission when it is not known where the locations of all the associated young stellar objects are in a particular region with respect to the masers. Clearly, traditional methods of searching for massive stars by imaging in the cm radio continuum will not reveal pre-ionizing protostars in embedded envelopes (so-called ``high mass protostellar objects'' or HMPOs). Furthermore the jets, outflows, and winds from lower mass non-ionizing stars have been found to be associated with water maser emission (e.g., Wilking \\& Classen 1987). Since massive stars are believed to form in clusters, one would expect there to be stars of a range of masses and phases of early stellar evolution present. Consequently, these regions need to be imaged at wavelengths other than in the cm continuum to find all the stellar sources within the regions of maser emission. However, massive star forming regions are generally located deep within giant molecular clouds, and are thus obscured at visible wavelengths. Dust absorbs and scatters visible light in these regions giving rise to significant extinction. However, infrared radiation is much less affected by extinction than visible radiation, and thus infrared imaging can probe through the cool obscuring dust enshrouding the massive stellar environment.  Infrared observations of these regions must be performed at moderate to high spatial ($\\lesssim$1$\\arcsec$) resolutions. Given that the average distances to massive star forming regions are several kiloparsecs away, and that massive stars tend to form in a highly clustered way, resolution is an issue when trying to determine which young stellar object is most closely associated with the maser emission. However, large ground-based telescopes with mid-infrared detectors yield both the required resolution and the ability to penetrate the significant obscuration in these regions. Mid-infrared radiation traces the warm dust close to young stellar sources, allowing one to observe the spatial relationship between masers and young stars. Combined with comparable spatial resolution cm radio continuum maps of these regions, mid-infrared images will allow identification of the vast majority of stellar and protostellar sources in each region of maser emission with none of the effects of confusion from foreground or background objects.  We present here a mid-infrared imaging survey of 26 water maser sites taken mostly from a list of sources from Forster \\& Caswell (1989) and Hofner\\& Churchwell (1996). Most of these sources are imaged here at $\\sim$1$\\arcsec$ resolution in the mid-infrared for the first time. The general goal of this work is to try to determine the relationship between water masers and massive star formation processes. A small subset of this survey was already published by De Buizer et al. (2003), which was concentrated on specific mid-infrared sources believed to be associated with water maser emission from HMPO candidates. This paper is a broader and larger-scale mid-infrared survey designed to study the full array of phenomena water (and to a small extent OH masers) may trace in these massive star forming regions. This mid-infrared survey of water maser sites is meant to be complementary to our results from our mid-infrared (De Buizer, Pi\\~{n}a, \\& Telesco 2000) and near-infrared (De Buizer 2003) surveys of massive stars associated with methanol maser emission.  In \\S2 we will discuss the observations, and in \\S3 the data reduction. We will discuss what is already known about each of the individual regions of the survey, and what new things we have learned in light of our mid-infrared observations in \\S4. In \\S5 we will discuss what we have learned from this survey as a whole, and we will compare and contrast this to the results of De Buizer, Pi\\~{n}a, \\& Telesco (2000) for massive star forming regions associated with methanol masers. We will end with our conclusions in \\S6. ",
        "conclusions": "We observed 26 regions of massive star formation in the mid-infrared that contain water and OH maser emission in order to study the relationship between the stellar sources, dust, and maser emission. From our observations we draw the following main conclusions:  \\begin{enumerate}  \\item Water masers are found to be associated in some cases with young embedded stars. This was seen in De Buizer et al. (2003) and there are a handful of other candidate embedded sources found in this larger survey. These embedded sources may be HMPOs or perhaps even young embedded intermediate mass stars, given their mid-infrared luminosities.  \\item Some distributions of water masers with respect to the mid-infrared and/or radio sources on the field would seem to indicate that some water masers are indeed associated with outflows from these objects.  \\item IRAS fluxes are not a good indicator of the spectral energy distribution of the sources of maser emission. Since nearly two-thirds of the fields where mid-infrared sources were detected exhibit multiple mid-infrared sources (14/20), and the remaining third, several have nearby sources seen at other wavelengths (e.g., radio or sub-mm sources), it can be concluded that massive stars are in general gregarious by nature and form in a clustered way. This means that even the angular resolutions of the Midcourse Space Experiment (MSX) satellite ($\\sim$5$\\arcsec$ at 8 $\\mu$m, $\\sim$13$\\arcsec$ at 21 $\\mu$m), though many times better than IRAS, will suffer confusion when observing these regions. Even the Spitzer Space Telescope ($\\sim$2$\\arcsec$ at 8 $\\mu$m, $\\sim$6$\\arcsec$ at 24 $\\mu$m) will have confusion problems at longer wavelengths.  \\item For the handful of sources where near-infrared observations were made by Testi et al. (1994, 1998) we conclude based upon the near- to mid-infrared slope of their spectral energy distribution that they are very young embedded objects (Class I or 0). We will explore the nature of the near-infrared emission from these regions using our own observations as well as those from 2MASS in a following paper.  \\item Mid-infrared emission seems to be more closely associated with water and OH maser emission than cm radio continuum emission from UC \\ion{H}{2} regions, and to a lesser extent, near-infrared emission.  \\item Methanol masers are not the only maser population with linearly distributed maser groups. We find that the percentage of linearly distributed water masers is the same (38\\%) as the percentage of linearly distributed methanol masers, and that OH masers in our survey are even more prone to linear distributions (50\\%). Given the mid-infrared morphologies of the sources associated with the masers, and the outflow observations of De Buizer (2003), we conclude that the vast majority of linearly distributed masers are not tracing circumstellar disks.  \\item This mid-infrared survey and that of De Buizer, Pi{\\~ n}a, \\& Telesco(2000) have confirmed that maser emission in general can trace a variety of phenomena associated with massive stars including shocks, outflows, infall and circumstellar disks. No one maser species is linked exclusively to one particular process or phenomenon. This implies that the conditions for maser excitation are not very stringent, given that each of the phenomena just mentioned would be associated with different densities, temperatures, and molecular abundances.  \\item The three different masers types (water, OH, and methanol) do not seem to be associated with different early evolutionary stages of massive stars. Instead it appears that they all trace a variety of stellar phenomena throughout many early stages of massive stellar evolution. This conclusion is based on the fact that the maser sources exhibit the presence or absence of emission at near-infrared, mid-infrared, submillimeter, and radio wavelengths in a variety of combinations and degrees. We also argue that the masers with no associated radio continuum emission could be accounted for by excitation from a small subset of the population of young intermediate mass stars.  \\end{enumerate}"
    },
    "0410/astro-ph0410582_arXiv.txt": {
        "abstract": "The gravitational collapse of a spherical cloud core is investigated by numerical calculations.  The initial conditions of the core lie close to the critical Bonnor-Ebert sphere with a central density of $\\sim 10^4$ cm$^{-3}$ in one model ($\\alpha=1.1$), while gravity overwhelms pressure in the other ($\\alpha=4.0$), where $\\alpha$ is the internal gravity-to-pressure ratio. The $\\alpha=1.1$ model shows reasonable agreement with the observed velocity field in prestellar cores.  Molecular distributions in cores are calculated by solving a chemical reaction network that includes both gas-phase and grain-surface reactions.  When the central density of the core reaches $10^5$ cm$^{-3}$, carbon-bearing species are significantly depleted in the central region of the $\\alpha=1.1$ model, while the depletion is only marginal in the other model.  The two different approaches encompass the observed variations of molecular distributions in different prestellar cores, suggesting that molecular distributions can be probes of contraction or accumulation time scales of cores. The central enhancement of the NH$_3$/N$_2$H$^+$ ratio, which is observed in some prestellar cores, can be reproduced under certain conditions by adopting recently measured branching fractions for N$_2$H$^+$ recombination. Various molecular species, such as CH$_3$OH and CO$_2$, are produced by grain-surface reactions. The ice composition depends sensitively on the assumed temperature. Multi-deuterated species are included in our most recent gas-grain chemical network.  The deuterated isotopomers of H$_3^+$ are useful as probes of the central regions of evolved cores, in which gas-phase species with heavy elements are strongly depleted. At 10 K, our model can reproduce the observed abundance ratio of ND$_3$/NH$_3$, but underestimates the isotopic ratios of deuterated to normal methanol. ",
        "introduction": "Theoretical models show that molecular abundances in molecular clouds are time-dependent \\citep[e.g.][]{ph80}. Taking advantage of this dependence, the ages of cloud cores have been estimated. For example, cores with a high abundance of CCS and carbon chain species are considered to be young, while those with large abundances of NH$_3$ are considered old \\citep{su92}. There have been arguments, however, on the reliability of such ``chemical clocks'', because molecular abundances depend not only on the elapsed time but also on density, temperature, and the C/O ratio in the gas phase \\citep[e.g.][]{prat1997,th1998}. Moreover, single-dish observations, by which radio astronomers used to estimate molecular abundances, average cloud structure and abundances over a beam size.  During the last decade, independent and spatially-resolved estimates of molecular abundance, temperature, and density have become available.  Observations of sub-millimeter continua and extinctions of background stars have revealed density distributions in a variety of starless dense cores \\citep{wt1994,all2001}, showing that the density is almost constant inside a radius of several thousand AU, and decreases outside of this radius as $r^{-p}$, where $p$ is $2.0-2.5$ and $r$ is the distance from the core center.  Spatial distributions of molecular line intensities have been obtained by interferometric observations and detailed mapping via single dish telescopes.  Intensity ratios of different transitions show that the starless cores are almost isothermal with $T\\sim 10$ K \\citep{taf1998,taf2002,taf2004}, although a slight deviation of a few K may exist \\citep{ev2001, gal2002}.  Spatial distributions of molecular abundances within selected cores have been obtained by analyses that include radiative transfer, showing that CO, CS and CCS are depleted at the core center, while NH$_3$ and N$_2$H$^+$ maintain relatively constant abundances \\citep{cas1999,oha1999,bgn2001,cas2002a, taf2002, taf2004}. The central depletion indicates that adsorption of gas-phase species onto grain occurs, a process that had long been anticipated theoretically \\citep{w76, bh1983, ta87}.  Since adsorption lowers gas-phase molecular abundances as a function of time, it can be used as another indication of the history of the source \\citep{bl1997}.  Considering that dense starless cores should be formed from lower density gas at $10^3-10^4$ cm$^{-3}$, we can estimate the contraction time scale via chemical models that include adsorption. \\citet{aik2001} and \\citet{aik2003} investigated the distribution of molecules in contracting cores by adopting a Larson-Penston (L-P) flow and analogues with slower rates of contraction.  They found the central depletion to be greater in models with slower rates of contraction.  The L-P model better reproduces observed molecular column densities, while the radial distributions of molecules are better reproduced by a contraction that is three-fold slower.  Since predicted molecular column densities vary as contraction proceeds, \\citet{aik2003} were able to use their results to sort observed cores into an evolutionary sequence.  Mono-deuterated species were also included in their chemical model.  They obtained very high D/H ratios such as H$_2$D$^+$/H$_3^+ > 1$, which is consistent with the recent observation of H$_{2}$D$^{+}$ in the core L1544 \\citep{cas2003}.  The L-P flow, however, is an asymptotic solution, which appears in the central region of a collapsing core.  The collapse velocity at outer radii, which is of interest in the comparison with observation, is overestimated in the L-P flow. In their models with slower rates of contraction, \\citet{aik2003} adopted a constant factor $f$ to decrease the contraction speed.  In reality,  the factor $f$ varies with time. To avoid this problem, \\citet{li2002} and \\citet{sh2003} directly coupled calculations of molecular evolution with magnetized core collapse.  They showed that the molecular column densities in L1544 are well reproduced by the magnetized model.  Although the infall velocities obtained in their models are generally faster ($\\gtrsim 0.2$ km s$^{-1}$) than observed in L1544 ($\\sim 0.1$ km s$^{-1}$), the magnetized model gives a somewhat smaller, and thus better, infall velocity than the non-magnetized model. They also suggested that the molecular line widths will vary depending on the spatial distribution of the abundances and radial velocity.  The present paper is an update of \\citet{aik2003}. We now calculate molecular distributions while solving for the contraction of a spherical core numerically, rather than by adopting an L-P flow, so that our infall velocities can be compared with observed values.  For simplicity, we restrict ourselves to the non-magnetized case; this simplification enables us to separate the dynamics and chemistry and thus to utilize a large chemical network. The infall speed is dependent on the ratio between the gravitational and pressure forces, for which we consider two different cases.  We will show that observed infall velocities are reproduced, without the help of a magnetic field, if the initial condition of the core is close to the critical Bonnor-Ebert sphere \\citep{bon1956} with central density $\\sim 10^4$ cm$^{-3}$. On the other hand, magnetic fields are probably needed in more massive or less pressurized cores to maintain the infall speed close to observed values. Our newly calculated molecular distributions are presented in selected stages of contraction and compared with recent observations of prestellar cores such as L1498, L1517B, L1521E, and L1544. As shall be shown, our model reproduces some observed features including an enhancement of the NH$_3$/N$_2$H$^+$ ratio in the central regions, which was not explained in previous studies.  Our basic chemical model contains a large gas-phase network \\citep[e.g.][]{th1998} that includes mono-deuterated species, and grain-surface reactions. In addition, we have extended the network for some calculations to include multi-deuterated species \\citep{rhm2003}. In recent years several multi-deuterated species have been detected in prestellar and protostellar cores, and their relative abundances to normal isotopomers, known as D/H ratios, found to be very high ($10^{-3}-10^{-1}$) \\citep{vdt2002, lis2002, par2002, bac2003, par2004}. Since the D/H ratio increases significantly as density increases and as gas-phase CO is depleted, it is of interest to see if the abundances of multi-deuterated species can be reproduced in our contracting core models.  The remainder of the paper is organized as follows.  A description of the model is given in \\S 2. Numerical results are presented and compared with observations in \\S 3. In \\S 4, we summarize the chemical and kinematic properties of some prestellar cores, and briefly discuss the formation of cores and the role of turbulence. A summary is contained in \\S 5. ",
        "conclusions": "\\subsection{Chemical and Kinetic Properties of Observed Prestellar Cores}  We now revisit the observed properties of the prestellar cores L1498, L1517B, L1521E, and L1544, and discuss which model (i.e. $\\alpha=1.1$ or 4.0) is preferable for each object. The upper part of Table \\ref{comp} summarizes chemical properties that are described in \\S 3.2. The values in brackets are the preferred $\\alpha$ for each property and object.  Starting with depletion, we see that $\\alpha=1.1$ is preferable for reproducing the heavy depletion of CO and CS in L1498, L1517B, and L1544, while the model with $\\alpha=4.0$ accounts for no depletion in L1521E. The abundance ratio of NH$_3$/N$_2$H$^+$ is enhanced towards the core center by a factor of $\\sim 15$ in L1498 and L1517B and by a factor of $\\sim 3$ in 1544. Considering their central densities, $\\alpha=1.1$ is preferable for L1498 and L1517B, while both models ($\\alpha=1.1$ and 4.0) can reproduce the modest enhancement in L1544.  Comparison via molecular column densities are not listed, since these distinctions are less obvious than the spatial distribution; the preferable $\\alpha$ varies with species for each object. The molecular column density is directly affected by uncertainties in both observation (e.g.  assumed geometry, velocity field and temperature in the analysis) and reaction rate coefficients (\\S 3.3).  The kinematic properties of the objects are listed in the lower part of Table \\ref{comp}. Subtracting the thermal component from the intrinsic line width, the non-thermal components of the N$_2$H$^+$ line width are calculated to be 0.17 km s$^{-1}$ in L1498 and L1517B \\citep{taf2002}, and 0.27 km s$^{-1}$ in L1521E. \\citet{hir2002} also obtained a relatively large intrinsic line width of $0.3-0.5$ km s$^{-1}$ for species such as C$_3$S ($J=4-3$) and H$^{13}$CO$^+$ ($J=1-0$) in L1521E. Infall velocities in L1498 and L1544 are estimated from CS and N$_2$H$^+$ lines by \\citet{lee2001}. \\citet{oha1999} observed CCS in L1544, and derived a larger infall velocity (0.1 km s$^{-1}$) than \\citet{lee2001}. While $\\alpha=1.1$ is preferable for reproducing the narrow line width and small infall velocities in L1498, L1517B, and L1544, the large widths may be an indication of faster contraction (i.e. $\\alpha=4.0$) in L1521E.  Combining the chemical and kinematic properties, the model with $\\alpha=1.1$ is consistent with L1498, L1517B, and L1544, while the model with $\\alpha=4.0$ is preferable for L1521E. Although it is not listed in Table \\ref{comp}, L1689B could be another example which is fit better by $\\alpha=4.0$. \\citet{les2003} found that the molecular depletion in L1689B is much less significant than that of L1544, in spite of the fact that their central densities are similar. Molecular line widths in L1689B are also broader, indicating more active kinematics than those in L1544 \\citep{les2003}.    \\begin{deluxetable}{l c c c c} \\tabletypesize{\\scriptsize} \\tablecaption{Properties of the Observed Prestellar Cores \\label{comp}} \\tablewidth{0pt} \\tablehead{ Parameter &\\colhead{L1498} & \\colhead{L1517B}   & \\colhead{L1521E} & \\colhead{L1544} }  \\startdata \\multicolumn{5}{c}{Chemical Properties}\\\\ \\hline CO depletion & heavy & heavy & no & heavy \\\\ & [1.1]\\tablenotemark{a} & [1.1] & [4.0] & [1.1] \\\\ CS depletion & heavy & heavy & no & heavy \\\\ & [1.1] & [1.1] & [4.0] & [1.1] \\\\ NH$_3$/N$_2$H$^+$ central enhancement  & yes & yes & ... & yes \\\\ & [1.1] & [1.1] & ... & [either] \\\\ \\hline \\multicolumn{5}{c}{Kinematics}\\\\ \\hline Non-thermal Line width [km s$^{-1}$] & 0.17\\tablenotemark{b} & 0.17\\tablenotemark{b} & 0.27\\tablenotemark{c} & 0.2\\tablenotemark{b} \\\\ Infall Velocity [km s$^{-1}$] & 0.03-0.044\\tablenotemark{d} & & & 0.021-0.1\\tablenotemark{d,e} \\\\ & [1.1] & [1.1] & [4.0] & [1.1] \\\\ \\enddata  \\tablenotetext{a}{A preferred value of $\\alpha$} \\tablenotetext{b}{\\citet{taf2002}} \\tablenotetext{c}{\\citet{ts2004}} \\tablenotetext{d}{\\citet{lee2001}} \\tablenotetext{e}{\\citet{oha1999}}  \\end{deluxetable}   \\subsection{Formation of Cores and Role of Turbulence} The line widths of molecular cloud tracers are typically significantly larger than the thermal line width, which suggests that the clouds are dominated by non-thermal motions often identified as turbulence \\citep[e.g.][]{mg1988, fg1992, gb1998, bb2000}. The existence and dissipation of turbulence are considered to be one of the keys controlling star formation \\citep[e.g.][]{bkv2003, mk2004}. In our present paper, we do not consider turbulence. Since observed molecular line widths in many prestellar cores are dominated by thermal motions \\citep{taf2002, tuk2004}, this simplification is justified at least for those objects.  However, the different values of $\\alpha$ in our models can tell us something about the time scale for the ambient material to be accumulated onto the core and for turbulence to dissipate. Although the origin of turbulence is not well understood, turbulence will decay within a sound-crossing time scale, without excitation sources \\citep[e.g.][]{sog1998}. If the accumulation of interstellar gas is efficient and a core becomes too massive to be supported by thermal pressure before turbulence dissipates, it collapses dynamically after the dissipation (i.e. resulting in large $\\alpha$). On the other hand, if the accumulation is rather slow and turbulence dissipates before the core mass reaches the critical value, the core becomes a stable Bonnor-Ebert sphere. Additional accumulation will achieve a critical Bonnor-Ebert core with $\\alpha\\sim 1$.  As discussed in the previous subsection (see also \\S 3.1 and \\S 3.2), both the line width and molecular distributions obtained in our model with $\\alpha=1.1$ are consistent with the observations in L1517B and L1498. On the other hand, molecular distributions in L1521E are better reproduced by the $\\alpha=4.0$ model. Since observed line widths are relatively large, as predicted in the model, L1521E could  indeed be a core with large $\\alpha$. However, a clear infall signature is not observed in this object. Based on  the above core  formation processes, one  possibility is that  L1521E  has  accumulated  enough  mass to  be  recognized  as  a prestellar  core before the dissipation of the turbulence. In this case, the turbulence accounts for the relatively large line width, and the relatively fast accumulation (before the turbulence decays) accounts for the {\\it young} ($\\lesssim 10^5$ yr) molecular abundances with little or no depletion. Detailed observation of velocity fields in chemically young cores such as L1521E and studies on the origin of turbulence in molecular clouds are highly desirable in order to understand the formation of cores."
    },
    "0410/astro-ph0410061_arXiv.txt": {
        "abstract": "Present and future observations of supernova relic neutrinos (SRNs), i.e., a cosmological neutrino background from past core-collapse supernova explosions, potentially give us useful information concerning various fields of astrophysics, cosmology and particle physics. We review recent progress of theoretical and observational studies of SRNs, particularly focusing on the detectability and also on implications for cosmic star formation history and neutrino physics. ",
        "introduction": "\\label{sec:Introduction}  A core-collapse supernova explosion is one of the most spectacular events in astrophysics, and it attracts a great deal of attention from many physicists and astronomers. It also produces a number of neutrinos and 99\\% of its gravitational binding energy is transformed to neutrinos; detection of the galactic supernova neutrino burst by ground-based large water \\v Cerenkov detectors, such as Super-Kamiokande (SK) and Sudbury Neutrino Observatory (SNO), would provide valuable information on the nature of neutrinos as well as supernova physics. In addition, because supernova explosions have occurred very commonly in both the past and present universe, tracing the cosmic star formation rate (SFR), they should have emitted a great number of neutrinos, which now make a diffuse background, i.e., supernova relic neutrinos (SRNs). Involved physics in SRN ranges quite widely---from cosmic SFR and supernova physics to neutrino properties as elementary particles. Therefore, detecting SRNs or even setting limits on their flux can give us quite useful and unique implications for various fields of astrophysics, cosmology and particle physics. Detectability of SRNs in various detectors and its implications have been discussed in many theoretical papers \\cite{Bisnovatyi-Kogan84}--\\cite{Cocco04} from various points of view, about which we discuss in detail in the following part of the present paper.  Flux estimation requires models of neutrino spectrum emitted from each supernova explosion and cosmic SFR. Furthermore, it is experimentally established that neutrinos mix among different flavours, altering the neutrino spectrum at detectors from the original one after neutrinos propagate inside a supernova envelope; this effect should now be taken into account appropriately. In addition to a flux estimation as precise as possible, a detailed discussion of background events, which hinder the SRN detection, is essential; it has been thoroughly studied by Ando, Sato and Totani \\cite{Ando03a} and we follow their discussion later in this paper. A stringent observational upper limit on the SRN flux is obtained by the SK group \\cite{Malek03} and its value is only a factor of 3--6 larger than several theoretical predictions \\cite{Totani96a}--\\cite{Hartmann97}, \\cite{Ando03a}, being consistent with (and more stringent than) the theoretical upper limit given by Kaplinghat, Steigman and Walker \\cite{Kaplinghat00} using conservative models. In order to make the SRN detection more likely, a promising method was proposed by Beacom and Vagins \\cite{Beacom04a}. Their basic idea is to dissolve gadolinium trichloride (GdCl$_3$) into the water \\v Cerenkov detectors, which greatly reduces the background events if it is applied to the currently working or proposed future detectors such as SK, Hyper-Kamiokande (HK) and Underground Nucleon Decay and Neutrino Observatory (UNO). Therefore, we are now at an exciting stage, where SRNs would soon be detected actually, and be used to obtain several implications for various fields of astrophysics as a unique and complementary method to usual observations of the light.  Since we have a promising prospects of the SRN detection, it is timely to discuss the SRN potential for probing the universe or particle physics. Cosmic SFR is one of such possibilities. The most popular method inferring SFR is to use the galaxy luminosity function of rest-frame ultraviolet (UV) radiation \\cite{Lilly96}--\\cite{Steidel99}, in addition to far-infrared (FIR)/sub-millimeter dust emission \\cite{Hughes98,Flores99} and near-infrared (NIR) H$\\alpha$ line emission \\cite{Gallego95}--\\cite{Glazebrook99}. In these traditional approaches, however, there are a fair number of ambiguities when the actual observables are converted into the cosmic SFR \\cite{Somerville01}. The most serious problem is that the effects of dust extinction are non-negligible, especially for rest-frame UV observations. For the approach using SRNs, on the other hand, we are not troubled by such a problem, because neutrinos are completely free of dust extinction. This point is the same with observations in the sub-millimeter wave band; however, neutrinos are emitted directly from stars, whereas sub-millimeter radiation comes from dust and is an indirect process. Another advantage is that supernovae are directly connected with the death of massive stars with $M \\gtrsim 8M_\\odot$, whose lifetime is expected to be very short compared with the Hubble time-scale $H_0^{-1}$. It enables direct inference of the cosmic SFR assuming the initial mass function (IMF), not bothered with uncertainty concerning the galaxy luminosity fuction, which sometimes causes difficulty in the case of rest-frame UV observation \\cite{Somerville01}. Motivated by all of these reasons, many researchers have investigated the dependence of the SRN flux on various models of cosmic SFR or to what extent the SFR can be probed from the future SRN observations. Totani, Sato and Yoshii \\cite{Totani96a} gave flux estimation using several theoretical models of galaxy evolution. After the bursting release of the observational SFR data points since the pioneering study by Madau \\etal \\cite{Madau96a}, the SRN flux has been estimated using these SFR results \\cite{Malaney97,Hartmann97,Ando03a}. More recently, theoretical SRN flux calculations have been compared with the observational upper limit by SK \\cite{Malek03}. Fukugita and Kawasaki \\cite{Fukugita03} used the limit to obtain constraints on the cosmic SFR; Strigari \\etal \\cite{Strigari04} adopted the latest SFR model based on the Sloan Digital Sky Survey \\cite{Glazebrook03} and concluded that their median SRN flux is slightly below the current SK upper limit. Following the proposal to dissolve Gd into detectors \\cite{Beacom04a}, Ando \\cite{Ando04a} investigated potential performance of the Gd-loaded water \\v Cerenkov detectors such as SK, HK and UNO, using the Monte Carlo (MC) simulation, especially focusing on how reliably an assumed SFR model can be reproduced from the SRN observation.  Neutrino properties can also be probed by the SRN observation in principle. It is pointed out that the resulting SRN signal depends on neutrino oscillation models \\cite{Ando03c}; particularly if the mass hierarchy is inverted, i.e., the first mass eigenstate that most strongly couples to electron flavour is heavier than the third state, $m_1>m_3$, and the value of $\\theta_{13}$, which has not been well constrained yet, is sufficiently large, then the flavour conversion during propagation inside the supernova envelope could considerably change the neutrino signal. Another interesting possibility is that a stringent constraint can be set on neutrino decay models from the SRN observation as first pointed out by Ando \\cite{Ando03f} and succesively studied by Fogli \\etal \\cite{Fogli04}. Non-radiative neutrino decay can be induced by the interaction between neutrinos and massless or very light particles such as Majoron. The strongest lower limit to the neutrino lifetime-to-mass ratio is obtained from the solar neutrino observation \\cite{Beacom02}--\\cite{Eguchi03b} and meson decay \\cite{Barger82}--\\cite{Britton94} to be $\\tau /m \\gtrsim 10^{-4}$ s/eV. SRN observation, on the other hand, is sensitive to $\\sim 10^{10}$ s/eV, which is many orders of magnitude larger than the current limit. This is because of the much longer baseline between the Earth and cosmological supernovae, from which SRNs are emitted, than 1 AU in the case of solar neutrino observations. Neutrino decay could also change the detected signal from high-energy astrophysical objects \\cite{Beacom03} or the galactic supernova explosions \\cite{Frieman88}--\\cite{Ando04c}, and could alter usual discussions on the early universe and structure formation \\cite{Beacom04b} as well as on supernova coolings \\cite{Fuller88}--\\cite{Farzan03}; but we stress that in principle we can obtain the most stringent constraints on the neutrino lifetime from the SRN observations.  This paper is organized as follows. We first introduce formulation for calculating the SRN flux in \\sref{sec:Formulation and models}. Models of cosmic SFR, supernova neutrino spectrum calculated by various groups are given in the same section. Neutrino oscillation inside the supernova envelope is briefly described there. SRN flux and event rate at ground-based detectors calculated with our reference models are given in \\sref{sec:Flux and event rate of supernova relic neutrinos}. \\Sref{sec:Detectability and observational upper limit} is devoted to a detailed discussion on background events against the SRN detection and detectability at various detectors. Current observational limit by SK, which is obtained by a statistical argument including background events, is briefly summarized also in the same section. We discuss implications of the SRN observation for cosmic SFR and neutrino properties such as oscillation and decay in sections \\ref{sec:Implication for cosmic star formation history} and \\ref{sec:Implication for neutrino properties}, respectively. Finally, we conclude this paper by giving brief summary in \\sref{sec:Conclusions}. Throughout this paper, we only consider electron anti-neutrinos ($\\bar\\nu_\\rme$) at detectors because this kind of flavour is most easily detected at the water \\v Cerenkov detectors, on which we mainly focus. ",
        "conclusions": "\\label{sec:Conclusions}  As stressed several times in this paper, we are at an exciting era of the SRN observation. This is because the current SK upper limit \\cite{Malek03} is {\\it just above} several theoretical predictions using realistic models for the cosmic SFR and supernova neutrino spectrum, and also because of the promising technique proposed in \\cite{Beacom04a} could greatly improve performance of the presently working or planned future detectors, especially for detecting SRNs. Being at such a situation, we believe that it is timely to review recent progress of theoretical and observational researches concerning SRNs from various points of view. Many physical and astronomical consequences can be derived from the constraints on the SRNs as we have reviewed through this paper.  Basics of SRN calculation were given in sections \\ref{sec:Formulation and models} and \\ref{sec:Flux and event rate of supernova relic neutrinos}. Models involved in the SRN calculations are those of the cosmic SFR and original supernova neutrino spectrum. Although the SFR-$z$ relation has been intensively studied in recent years, there remains a fair amount of uncertainty; in fact, the best estimate of the local SFR density ranges fairly widely as (0.5--2.9)$\\times 10^{-2}h_{70}$ yr$^{-1}$ Mpc$^{-3}$ \\cite{Baldry03}. In order to take this uncertainty into account, we have introduced the correction factor $f_\\ast$ in equation \\eref{eq:SFR}, for which a wide range (0.7--4.2) is still allowed from the SFR observation. As for the original supernova spectrum, contrary to the traditional approach using the Fermi-Dirac distribution, we followed the approach taken by our group \\cite{Ando03a,Ando04a}, i.e., adopting the results of numerical simulations. Especially following \\cite{Ando04a}, we used three neutrino spectra by different groups, i.e., LL \\cite{Totani98a}, TBP \\cite{Thompson03} and KRJ \\cite{Keil03}, and performed comparison among models. Recent progress of neutrino experiments has proven that neutrinos are massive and mix among different flavours. This effect was also taken into account appropriately. As the result, it was found that the uncertainty concerning the original neutrino spectrum gives difference in the resulting SRN flux by at least a factor of 2--4. The expected event rate at a water \\v Cerenkov detector with a fiducial volume of 22.5 kton was found to be 0.25--1 $f_\\ast$ yr$^{-1}$ (for $E_\\rme > 19.3$ MeV) and 1--2 $f_\\ast$ yr$^{-1}$ (for $E_\\rme > 10$ MeV), depending on these supernova models (see \\tref{table:event rate}).  Besides the flux estimation, a careful discussion about background events is definitely necessary in order to investigate the detectability. This has been thoroughly argued in \\cite{Ando03a} and we followed them in \\sref{sec:Detectability and observational upper limit}. The most serious background comes from decay products of invisible muons, which are produced almost at rest inside the detector and that is why ``invisible.'' Because of the background, the detection is severely restricted; for a pure-water \\v Cerenkov detector with the size of SK, it would take about (maybe more than) 10 years in order to reach the signal-to-noise ratio of a few (see \\fref{fig:SN_ratio}). We also discussed performance of other detectors such as SNO and KamLAND in the same section, by mainly following \\cite{Ando03a,Strigari04}. The current observational upper limit by SK \\cite{Malek03} and the proposed technique that potentially improves the performance of water \\v Cerenkov detectors \\cite{Beacom04a} were briefly introduced there. In addition, we showed that the SK limit on the SRN flux can be fairly well reproduced by a simple statistical argument using equation \\eref{eq:SN ratio}.  \\Sref{sec:Implication for cosmic star formation history} was devoted to the discussion on the current status and future prospects for obtaining limits on the cosmic SFR from the SRN observations. We first showed that the current SRN limit from SK \\cite{Malek03} already sets a stringent limit on the local SFR value, as discussed in \\cite{Fukugita03}. Surprisingly, it might rule out a part of local SFR value allowed by the current observations using the light as shown in \\tref{table:constraints}. Performance of future detectors loaded with Gd as proposed in \\cite{Beacom04a} was investigated in detail using the MC simulation, following discussion in \\cite{Ando04a}. The distribution of the best-fit values of two parameters concerning SFR \\eref{eq:parameterization} was obtained by the generated $10^3$ MC data and by accompanying likelihood analyses. If one of two free parameters was fixed somehow in advance using other observations, then the SRN observation can constrain the remaining parameter by accuracy of $\\sim 0.3 (f_\\ast V_{\\rm eff}/22.5 ~{\\rm kton} \\cdot 5 ~{\\rm yr})^{-1/2}$; this number would soon become quite significant if the future mega-ton class detectors such as HK or UNO loaded with Gd started data-taking.  Finally in \\sref{sec:Implication for neutrino properties}, SRN constraints on particle physics models of neutrinos were discussed. Neutrino oscillations in the case of inverted mass hierarchy \\cite{Ando03c}, resonant spin-flavour conversions induced by the neutrino magnetic moment (see \\cite{Ando03g} and references therein) and neutrino decay \\cite{Ando03f,Fogli04} are expected to considerably change the SRN signal at the detectors. In particular for neutrino decay model, the SRN observation potentially gives the most stringent constraint on the neutrino lifetime, compared with a galactic supernova \\cite{Frieman88}--\\cite{Ando04c}, high-energy astrophysical objects \\cite{Beacom03} as well as the most reliable (but weak) limit at present by the solar neutrino observations \\cite{Beacom02}--\\cite{Eguchi03b}.   \\ack SA's work is supported by a Grant-in-Aid for JSPS Fellows. KS's work is supported in part by a Grant-in-Aid for Scientific Research provided by the Ministry of Education, Culture, Sports, Science and Technology of Japan through Research Grant No S14102004, a Grant-in-Aid for Scientific Research on priority Areas No 14079202."
    },
    "0410/astro-ph0410127_arXiv.txt": {
        "abstract": "We use the KPNO International Spectroscopic Survey (KISS) for emission-line galaxies to identify and describe a sample of local analogues to the luminous compact blue galaxies (LCBGs) that are observed to be abundant at intermediate and high redshift.  The sample is selected using criteria believed effective at isolating true examples of LCBGs: $SB_{e}$(B-band) $< 21.0$ mag arcsec$^{-2}$, $M_{B}< -18.5$ (for $H_0= 75$ km s$^{-1}$ Mpc$^{-1}$), and $B-V < 0.6$.  Additionally, all LCBG candidates presented are selected to have star-formation as their dominant form of activity.  We examine the properties of our LCBGs and compare them to those of other KISS star-forming galaxies of the same absolute magnitude range. We find that the KISS LCBGs lie on the extreme end of a fairly continuous distribution of ``normal'' star-forming galaxies in the plane of surface brightness versus color.  This result differs from the results of previous studies that show LCBGs at higher-z to be more separate from the ``normal'' (usually non-active) galaxies they are compared against.  On average, LCBGs have a higher tendency to emit detectable flux in the radio continuum, have higher H$\\alpha$ luminosities by a factor of 1.6, indicating strong star-formation activity, and have slightly lower than expected metal abundances based on the luminosity-metallicity relation for KISS galaxies.  We calculate the volume density of our low-$z$ ($z<0.045$) sample to be $5.4\\times10^{-4}h_{75}^{3}Mpc^{-3}$, approximately 4 times lower than the volume density of the LCBGs at 0.4 $<$ z $<$ 0.7 and $\\sim$10 times lower than the volume density of the population at 0.7 $<z<$ 1.0. ",
        "introduction": "\\label{intro}  Common at intermediate and high redshifts, yet rare in the local universe, luminous compact blue galaxies (LCBGs) typically have luminosities equal to or greater than that of the Milky Way but are considerably smaller in size \\citep{koo94,phil97,guz97,ostlin,cfrs}. As evidenced by such properties, they are among the most extreme galaxies known in the universe. Although these galaxies have been the subject of numerous comprehensive studies in the past decade, their evolution and nature, to a large extent, remains in dispute \\citep{koo95,phil97,hammer,ostlin2,barton, pisano,guz03}. At the heart of the issue lies the difficulty of identifying nearby analogues that would allow a more complete exploration of the mechanisms responsible for their extreme properties at high redshifts, as well as their evolutionary paths.  Characterized by high surface brightnesses, small half-light radii, and vigorous star-formation, LCBGs were initially identified as stellar objects from 4-m plate surveys for QSO candidates \\citep{koo92,koo94,guz96,guz98}. The formal definition of LCBGs is developed by \\cite{anna} from  examining regions of six-dimensional parameter space (color, luminosity, asymmetry, concentration, size, and surface brightness) occupied by three classes of compact galaxies chosen to represent the population of luminous, blue, compact galaxies as a whole. The three classes are: (1) Compact Narrow Emission-Line Galaxies (CNELGs, Koo \\etal 1994, 1995; Guzm\\'{a}n 1996, 1998), (2) a subset of galaxies very similar to CNELGs studied by \\cite{koo95,phil97} and \\cite{guz97} called faint blue galaxies, and (3) higher redshift Blue Nucleated Galaxies (BNGs, Schade \\etal 1995, 1996). While the classification of an LCBG relies on the galaxy's falling within three of five defined regions in the aforementioned parameter space, recently groups have adopted three cut-offs that roughly correspond to the criteria laid out by \\cite{anna}: $B-V < 0.5-0.6$, $SB_e < 21-21.5$ mag arcsec$^{-2}$, and $M_B < -18.5$, with $H_0= 70$ km s$^{-1}$ Mpc$^{-1}$ \\citep{garland,pisano}. These selection criteria will be further examined in subsequent sections.   With a volume density that drops off substantially from $z=1$ to the present (previously estimated to be a change of roughly a factor of ten), LCBGs appear to be the most strongly evolving galaxy population \\citep{lilly96,phil97,guz97,lilly98,cfrs}. \\cite{guz97} estimate that they contribute nearly $40\\%$ of the fractional increase in the SFR density between $z=0$ and $z=1$. Furthermore, luminous, compact star-forming galaxies appear to dominate high-redshift galaxy samples. \\cite{low97} have suggested that the Lyman Break Galaxies at $z=3$, due to their very compact cores and high surface brightnesses, may be the high-redshift counterparts of the intermediate-z LCBGs. \\cite{smail98}, in a sub-millimeter galaxy survey, found a $z\\sim4$ galaxy population that may contribute a substantial fraction of the SFR density in the early universe, nearly half of which are very luminous and compact. Evidently, LCBGs play a prominent, if not paramount, role in the formation and evolution of galaxies in the young universe.  A drove of unanswered questions hovers around the subject of LCBGs, mainly owing to the lack of a representative local sample. One of these questions involves the mechanism(s) responsible for the observed intense starbursts in LCBGs. In theory, these triggering mechanisms could be determined by examining the H$\\alpha$ velocity fields and morphologies of local LCBG examples \\citep{ostlin}. This sort of study relies on the assumption that one can select a sample of local LCBGs to have the same properties as the more distant LCBGs. \\cite{ostlin} use a sample of six galaxies with $M_{B}$ between -17.5 and -20 selected from local blue compact galaxies (BCGs) that are bright in H$\\alpha$ to infer certain kinematic properties of the higher-$z$ LCBGs. They conclude that, in most cases, mergers or strong interactions trigger the intense starbursts of LCBGs. However, one should use caution in interpreting these results given the size and properties of their sample. Indeed, a statistically complete sample of low-z LCBGs will provide more clues as to the the origin of their activity.  Another crucial topic in the study of LCBGs has been the question of what sort(s) of galaxy they become as they evolve. The determination of reliable mass estimates for the intermediate-$z$ LCBGs is key to establishing these evolutionary connections to present-day galaxies. However, accurate stellar mass estimates are difficult to determine and measurements of emission-line widths used to indicate the virial mass may underrepresent the gravitational potential for the intermediate-redshift LCBGs \\citep{guz01,guz03,pisano}.  As a consequence of these difficulties, there is not good agreement for the mass estimates of LCBGs and there is a corresponding uncertainty as to their proposed evolutionary path(s).  \\cite{phil97}, \\cite{barton}, and \\cite{hammer} argue for an evolutionary scenario in which LCBGs represent the formation of the bulges of objects that will eventually evolve into massive spiral disks of today. Koo \\etal(1995), and Guzm\\'{a}n \\etal (1996, 1997, 1998) maintain that while LCBGs represent a heterogeneous class of galaxies, the majority will evolve into local low-mass dwarf elliptical galaxies.  A possible solution to the mass debate may lie in the HI line widths of local LCBGs. These line widths can be used to infer the total dynamical mass.  Accordingly, \\cite{pisano} set out to measure the HI line widths of a nearby sample of galaxies they believe are representative of the more distant LCBGs. They infer the masses of the intermediate-$z$ galaxies from these local examples, and find that LCBGs have a variety of dynamical masses, and thus a variety of evolutionary paths. Yet, they recognize that their local sample may not have been truly representative of the the high-$z$ LCBG population. Of their 11 local galaxies, only three meet the criteria proposed by \\cite{anna}. In a study of 20 local LCBG candidates selected to more nearly meet the the criteria of \\cite{anna}, \\cite{garland} found results similar to those of \\cite{pisano}. Specifically, they show that while most LCBGs are nearly ten times less massive than local galaxies of the same luminosity, some are just as massive. The implication is that LCBGs are a heterogeneous class of galaxy.  Clearly, defining reliable local samples of LCBGs is central to the process of understanding their nature and evolutionary paths.  The \\cite{pisano} and \\cite{garland} samples are recent attempts to construct samples of nearby LCBGs.  The latter study is based on a sample of LCBGs discovered by \\cite{castander} who analyzed data from the Sloan Digital Sky Survey.  Another recent study by \\cite{drinkwater} utilized the complete spectroscopic study of a 12 deg$^2$ area centered on the Fornax Cluster to identify a sample of 13 LCBGs with redshifts less than 0.21.  In the current study, we take a slightly different approach.  Using the KPNO International Spectroscopic Survey (KISS, Salzer \\etal 2000), we aim to identify and describe a statistically complete local sample of star-forming galaxies that best resemble those LCBGs observed in the more distant universe.  Our low-z sample will provide a basis for future work aimed at resolving the debates over the evolutionary histories and kinematics of LCBGs. Due to the limited resolution of the survey images, the classification scheme developed by \\cite{anna}, that uses morphological parameters, cannot be employed. Instead, we apply selection criteria consistent with the parameters of \\cite{anna} as a way to isolate the LCBGs from other KISS star-forming galaxies. While sharp boundaries in parameter space do not always effectively isolate one type of galaxy from another, LCBGs in the color-surface brightness plane occupy a well-defined region of parameter-space \\citep{anna}. Therefore, we employ criteria in \\emph{B}$-$\\emph{V} color, in surface-brightness, and in luminosity to select a sample of local LCBG candidates.  Still, none of the selection criteria we use can guarantee that the KISS LCBG candidates, like other putative local samples, will contain \\emph{bona fide} LCBGs; they only insure a similarity. Regardless, we believe that the sample we define will provide an adequate starting point for future investigation. In $\\S$ 2, we describe KISS, our measurements, and selection criteria. An analysis of the properties of our local LCBGs along with the volume density calculation appears in $\\S$ 3.  Section 4 contains a discussion of the implications our sample has for high-$z$ studies and potential future work on these galaxies beyond KISS. In $\\S$ 5, we summarize our results. Throughout this paper we adopt $H_0= 75$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "In this paper, we present a sample of local analogues to the intermediate- and high-$z$ LCBGs. We selected our sample from only those KISS galaxies that are actively forming stars. Due to resolution effects, we limit our sample to only those galaxies within 185 Mpc of the Milky Way.  The sample of LCBGs is defined by the following criteria:  a surface brightness limit of $SB_e <  21.0$ mag arcsec$^{-2}$, an absolute magnitude limit of $M_B < -18.5$, and a color cut-off of $B-V < 0.6$ with $H_0= 75$ km s$^{-1}$ Mpc$^{-1}$. A total of 16 KISS SFGs meet these criteria.  We have compared some of the properties of the KISS LCBG candidates to those of a comparison sample of KISS SFGs, and a comparison sample of Century Survey (CS) galaxies that are supposed to represent ``normal galaxies.'' On average, the KISS LCBG candidates have a higher tendency to emit detectable radio continuum flux, have higher H$\\alpha$ luminosities indicating strong star-formation properties, have slightly lower than expected metal abundances for KISS galaxies of their luminosity, and considerably smaller size than those galaxies in both KISS and CS comparison samples. We have calculated the volume density for the sample of local LCBG candidates to be $5.4\\times10^{-4}h_{75}^{3}Mpc^{-3}$.  This value indicates a substantial density evolution of LCBGs, which is in general agreement with the predictions made by \\cite{lilly96,phil97,lilly98} and \\cite{cfrs}.  However, precise estimates of the density evolution of the LCBGs are difficult, due to the large uncertainties in the values of the volume density for higher-redshift populations.  We have utilized the CS database in an attempt to select a sample of local non-active LCBGs in order to constrain the density of non-star-forming LCBGs.  We found no additional non-star-forming LCBGs, although one weakly star-forming (possibly post-starburst) galaxy was recognized.  We postulate that LCBGs are structurally predisposed for star-formation, and that their compact nature may be indicative of a high central gas concentration that drives ongoing star-formation.  The need for a local sample of LCBGs is well documented. Much of the published work on these galaxies speaks to the need for a representative local comparison sample that will help settle the debates over their size and mass, morphology, and low- redshift counterparts. We conclude that our sample of local LCBG candidates contains galaxies that are likely to be local versions of the more distant LCBGs. This local sample provides the starting point in our ongoing study of the properties and evolution of LCBGs."
    },
    "0410/astro-ph0410311_arXiv.txt": {
        "abstract": "We present evidence for a highly obscured X-ray source in one of two giant Ly$\\alpha$ emission nebulae in the $z=3.09$ proto-cluster region SSA22. Neither Ly$\\alpha$ nebula is associated with significant radio emission. While one has a significant submillimeter detection and is undetected in the X-ray, the other is a factor of 2-10 times less submillimeter bright and appears to contain a hard-band X-ray source. We discuss our analysis and techniques for assessing the X-ray properties of this source and suggest that we have detected an embedded AGN source in one of these nebulae which may be at least partially responsible for exciting the Ly$\\alpha$ emission through a mechanism that is essentially decoupled from the radio, submillimeter, or optical luminosities. We also present an upper limit on the mean X-ray emission from 10 other extended Ly$\\alpha$ objects in the SSA22 region. ",
        "introduction": "The standard paradigm of structure formation in a universe dominated by cold dark matter suggests that the high-density regions of the high-redshift universe that later collapse into clusters should also be host to the progenitors of massive galaxies. Identifying both proto-cluster regions and proto-elliptical galaxies therefore offers an opportunity to directly witness the astrophysical phenomena that likely set many of the characteristics of clusters and cluster galaxies seen today.  The identification of protocluster regions at $z>2$ remains a challenging task. Among the approaches used are the localization of spatial excesses of objects such as Lyman break galaxies \\citep{ste98}, Lyman emission line objects (e.g., \\citet{hu98}), and submillimeter sources (e.g., \\citet{ivi00}), and the apparent correlation of QSO absorption features along adjacent lines of sight \\citep{fra96}. A slightly different approach involves the use of signpost objects, such as luminous radio galaxies, which are expected to form in the highest density regions at high redshift \\citep{ste03}. All of the above methods have yielded positive results with varying degrees of success.  In a survey of continuum-selected Lyman break galaxies \\citet{ste98,ste00} have found evidence for a highly significant overdensity of objects at $z=3.09$ (a factor of $\\sim 5$ higher number density than the average equivalent volume at this redshift). In subsequent narrowband Ly$\\alpha$ imaging of this field, two very bright, large, diffuse Ly$\\alpha$-emitting nebulae (or ``halos'' or ``blobs'') were discovered. Named ``blob 1'' (B1) and ``blob 2'' (B2) by \\citet{ste00} these objects have sizes of 17'' and 15'' (B1, B2), Ly$\\alpha$ luminosities of $10^{44}$  ergs s$^{-1}$(B1) and $9\\times 10^{43}$ ergs s$^{-1}$ (B2), and rest equivalent widths of {\\em at least} $370$ \\AA. Both blobs have radio continuum fluxes at 1.4 GHz  at least  3 orders of magnitude fainter than those of typical luminous high-redshift radio galaxies (B1 has a tentative radio identification with a 44 $\\mu$Jy [1.4 GHz] source; \\citet{cha04}). High-$z$ Ly$\\alpha$ halos certainly exist with very powerful radio sources and jets (e.g., \\citet{bre03,reu03}); however these have been preselected by their radio luminosity, and so it is not clear whether low radio activity is common or not in such systems. submillimeter observations of the blobs have detected sources that (given the $\\sim 15''$ half-power beamwidth at 850 $\\mu$m) are potentially counterparts to B1 and B2 \\citep{cha01,cha04}. However, while the B1 counterpart was detected as a $20.1\\pm3.3$ mJy source at 850 $\\mu$m, for B2 the submillimeter emission was significantly less at $3.3\\pm1.2$ mJy. Thus, the B1 submillimeter counterpart is one of the most luminous submillimeter high-$z$ sources known, while the B2 counterpart is of low significance.  In recent follow-up work on B1, \\citet{cha04} have further identified several optical continuum components (to a limit $R\\sim 28.6$) that may be associated with the core of the submillimeter source, and that appear within, and are potentially associated with, the 17$''$ ($\\sim 140$ kpc) Ly$\\alpha$ blob. They also obtain a marginal detection of a CO(4-3) molecular line and estimate a molecular gas mass of $2\\times 10^{10}$M$_{\\odot}$. Furthermore, \\citet{cha04} find no significant X-ray emission associated with B1 using deep {\\sl Chandra} data, which we further analyze below.  The current physical interpretation of the B1 structure is that it is consistent with an undetected, dust-obscured active galactic nucleus (AGN), together with a starburst region embedded at the  core of  the  Ly$\\alpha$  nebula. Indeed, \\citet{cha04} speculate that a highly obscured  AGN could still be responsible for a jet  structure inclined  well  away from  our  line of sight that is inducing star  formation with the  observed spatial morphology. Ionizing flux from the AGN in such a geometry could then also assist in lighting up the Ly$\\alpha$ cloud. The only obvious difficulty  with this interpretation is the apparent  absence of extended radio structure. We will return to these scenarios in the Discussion below. ",
        "conclusions": "The evidence presented above for the presence of obscured AGN emission of $L_x\\sim 10^{44}$ erg s$^{-1}$ embedded in the Ly$\\alpha$ cloud/B2 of SSA22 provides the first direct indication of AGN activity in this possibly unusual, radio-quiet, object.  This detection provides significant clues as to what is going on in its ``companion'' Ly$\\alpha$ cloud, B1, and possibly to the generation of the Ly$\\alpha$ cloud emission itself.  It is useful to compare and contrast B1 and B2. These objects have comparable Ly$\\alpha$ luminosities, both are detected in the submillimeter, while B2 is a factor of $\\sim 10$ less luminous, and while B2 has an X-ray detectable, albeit highly obscured, AGN, B1 does not to the current sensitivity limits.  Various scenarios can be constructed from these observations. For example, {\\em if} the observed submillimeter luminosity is directly indicative of the mass of dust in a system (assuming similar temperatures) then it is tempting to conclude that B2 is less enshrouded than B1 and is therefore more transparent to high-energy X-rays (8-33 keV rest frame). While this may be correct, it is also possible that the obscuration of the AGN (in B2 for certain) is due to a much smaller scale molecular/dust torus surrounding the SMBH. The similarity in Ly$\\alpha$ luminosity is intriguing. If the ionizing source for the Ly$\\alpha$ cloud were directly connected with that responsible for heating the submillimeter dust (i.e. starburst activity) then some amount of coincidental tuning would be needed to produce a similar level of emission between B1 and B2.  Finally, we  offer a different, albeit speculative, mechanism that could explain many,  if  not   all,  of  the  above  features  of the  B1  and  B2 systems. Recently, evidence has  been found for extended X-ray emission around  high-$z$  radio-loud   galaxies  due to  the  inverse Compton scattering of  cosmic microwave background (CMB) photons by  a population of relativistic particles originating from the central SMBH in these objects \\citep{fab03,sch03}.  Owing to  the $(1+z)^{4}$  growth  of  the  CMB  energy density,  this  emission  is increasingly luminous at high-$z$, and  can readily exceed that of the central AGN \\citep{sch02,sch03}. In at  least one case \\citep{fab03}, much of the X-ray emission appears  to be due to an {\\em  older} relativistic population of much greater spatial extent than the currently detected radio synchrotron emission. The presence of a spatially extended, highly luminous, X-ray emission offers a ready mechanism for photo-ionization of cooler gas. Indeed, \\citet{sch03} speculate that it alone could sustain a $\\sim 100$ kpc Ly$\\alpha$ emission cloud, while not over-ionizing observed species such as those of oxygen. If B2 (and perhaps B1) contains an AGN, then it is plausible that a population of older relativistic particles exist on scales of $\\sim 100$ kpc, unseen via radio synchrotron since the typical emission frequencies will be less than 100 MHz (rest frame) but still actively upscattering CMB photons to the X-ray. Some of the ``triangular shaped'' or hourglass, morphology of the Ly$\\alpha$ clouds reported in B1 \\citep{cha04} and seen in B2 \\citep{ste00} matches extremely well the morphologies seen in the X-ray surrounding 3C 294 ($z=1.89$; \\citet{fab03}) and in {\\em both} the X-ray and Ly$\\alpha$ seen in 4C 41.17 ($z=3.79$; \\citet{sch03}).  The relatively low X-ray surface brightness of such emission ($\\sim 8\\times 10^{-17}$ ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$) would indicate that it can readily escape direct detection even in a 78 ks ACIS-S observation - although as discussed in Section 2, there is some suggestion that the  B2 X-ray emission could be extended. The phenomenon described above offers an interpretation if this is the case.  In such a picture, the relative submillimeter, radio, and optical luminosities associated with the Ly$\\alpha$ clouds are essentially decoupled from the mechanism responsible for exciting their fluorescent emission - which is driven by a periodic injection of relativistic particles (as seems to be a standard AGN property), and the enhanced energy density of the CMB (a factor of 280 at $z=3.09$)."
    },
    "0410/astro-ph0410591.txt": {
        "abstract": "{ Weakly Interacting Massive Particles are often said to be the best Dark Matter candidates. Studies have shown that large Dark Matter-photon or  Dark Matter-baryon interactions could be allowed by cosmology. Here we address the question of the role of the Dark Matter interactions in more detail to determine at which extent Dark Matter has to be necessarily weakly interacting. To this purpose, we compute the collisional damping (and free-streaming) scales of generic interacting  Dark Matter candidates  and investigate the effects on structure formation. Our calculations are valid provided the Dark Matter particles have experienced a phase of statistical equilibrium at some stage during their evolution. By comparing these damping lengths to the scale of the smallest primordial structures known to exist in the Universe, we obtain  necessary conditions that any candidate must satisfy. These conditions are  expressed in terms of the   Dark Matter particles' mass and either the total  Dark Matter interaction rate or the interaction rate of Dark Matter with a specific species. The case of Dark Matter interacting with neutrinos or photons is considered in full detail. Our results are valid even for energy dependent cross-sections and for any possible initial fluctuations spectrum. We point out the existence of new Dark Matter scenarios and exhibit new damping regimes. For example, an interacting candidate may bear a similar damping than that of collisionless Warm Dark Matter particles. The main difference is due to the Dark Matter coupling to interacting (or even freely-propagating) species. Our approach yields a general classification of  Dark Matter  candidates which extends the definitions of the usual Cold, Warm and Hot   Dark Matter   scenarios when interactions, weak or strong, are considered.} ",
        "introduction": "The damping of primordial fluctuations has been considered more than 30 years ago (\\cite{joe67,joe68, misner}) when baryonic matter was popular. In the early eighties, when the existence of   Dark Matter   became a serious possibility, it was quickly understood that fluctuations made of Weakly Interacting Massive Particles will render galaxy formation much easier (\\cite{gunn, peebles82, bond83}). If the   Dark Matter   interactions are weak enough and the   Dark Matter   mass heavier than a few keV, one indeed expects no collisional damping effects, and no free-streaming on scales relevant for galaxy formation. Such particles (\\eg neutralinos) are now embedded in the well-known Cold Dark Matter (CDM) scenario. As a limiting case, particles which marginally comply with the free-streaming constraints are usually referred to as Warm   Dark Matter (WDM) candidates, the archetype  being (\\cite{davis, blumen}) particles with a mass of $\\sim 1$ keV whose damping length is about $100$ kpc.    While both the CDM and WDM initial conditions reproduce rather well the mass distribution of the observed large scale structures, the CDM scenario is in danger to yield an excess of small galaxies ( \\cite{rijoe1,cdmcrisis,cdmcrisisb}). The WDM model, known to reproduce better the observed faint end of the galaxy multiplicity function (\\cite{ rijoe2}), was revived (\\textit{e.g.} \\cite{wdm1}) as a possible solution to the discrepancy between observations and  CDM numerical simulations. However, the collisionless WDM scenario does not seem to be the solution to all discrepancies (\\cite{julien}, \\cite{ostriker}), especially if one considers the constraints from the late reionization signal claimed to be found by WMAP( \\cite{WMAP,yoshidareio}).  Some of the potential problems of the CDM scenario may possibly be alleviated by invoking astrophysical processes (\\textit{e.g.} \\cite{chiu,stoehr,deckel}) or perhaps by redoing data analysis of the dwarf and Low Surface Brightness rotation curves (\\textit{e.g.} \\cite{swaters,autres}). Also, the Dark halo profile of the Milky Way may (\\cite{MWnew}) be more cuspy than previously found (\\cite{binney}). Should the still remaining discrepancy between numerical simulations and observations be due to the nature of Dark Matter or not, a thorough investigation of the   Dark Matter   properties seems worthwhile.  In particular, one may wonder whether taking into account the   Dark Matter interactions  -- generally considered to be negligible -- would help to solve the Dark Matter crisis at small scales.   This is discussed by \\cite{sperstein} who consider strongly self-interacting   Dark Matter   to comply with the shape of the galaxy rotation curves (\\cite{moore}). However, according to further investigations (\\cite{yoshida}, \\cite{clusters}), this suggestion, at least in its original form, seems not to be the answer.  One can then address a more general question to determine which kind of interaction  is allowed for   Dark Matter, with which particles, and with which strength.  The original idea that led to the introduction of weakly interacting Dark Matter particles (that are now embedded in the Cold Dark Matter scenario) was to avoid  Silk damping as well as prohibitive free-streaming. In the present paper, we extend this approach to other Dark Matter interactions, say interactions of Dark Matter with any other species, including itself and seek constraints on the allowed Dark Matter particles' parameter space. First insights were given in \\cite{bfsL,brhs,chen}. We here  undertake an exhaustive and extensive analysis by considering also new aspects, considering in particular the way the   Dark Matter particles achieve their relic density.     We start by relaxing the assumption of collisionless   Dark Matter   to estimate the   Dark Matter   collisional and free-streaming damping effects. Requiring that these two damping mechanisms do not prevent the formation of observed structures allows us to define the range of   Dark Matter   mass and interaction rates that are compatible with observations. This is independent of any particle physics. We  then classify all generic   Dark Matter candidates in a 2-dimensional parameter space (interaction rate versus mass) and determine what regions of this parameter space are allowed by the present observations.  Our calculations are valid for massive neutrinos, provided their mass is within the known experimental and observational limits. For the sake of comparison with the general case we treat here, we use at various places a \"reference\", the cosmology in which i) the radiation has two components (namely photons and massless neutrinos, both having the standard temperature) and ii) the matter is made of two species, the baryons and the Dark Matter.   As pointed out in \\cite{bfsL}, the collisional damping scale of the Dark Matter fluctuations can be expressed in terms of two separate contributions, namely the \\emph{\\,self-damping\\,} and \\emph{\\,induced-damping\\,} contributions.  The self-damping is related to the collisional propagation of the   Dark Matter   particles. It may be directly compared with their free-streaming scale.  The induced-damping corresponds to the (collisional or free-streaming) damping acquired by a species $i$ and transmitted to the   Dark Matter   fluctuations as long as the coupling between   Dark Matter   and this species is large enough. The induced-damping effect can  be seen as a generalization of the Silk damping (of baryons coupled to photons), here applied to the case of   Dark Matter   particles coupled to any possible species $i$. We  especially focus our attention on the coupling to neutrinos and  photons, seen to provide a quite large damping and  yield rather stringent  constraints.  Our results are given in two separate papers. The present one (referred to as paper I) is dedicated to the calculation of the free-streaming and collisional damping scales associated with any interacting  Dark Matter   particles. We then derive the mass and interaction rates of the   Dark Matter   candidates allowed by structure formation. This is a rather technical paper. The second one (paper II), deals with the physical and astrophysical relevance of the candidates whose damping scales, calculated here, are not prohibitive. These two papers have been written in a self-contained way and can be read independently.  The present paper is organized as follows: in Section \\ref{sec:transp}, we recall the expressions of the different transport coefficients for an imperfect fluid made of several species (relativistic or not), and take this opportunity to discuss issues for which there is still some confusion in the litterature. In Section \\ref{sec:dampscales}, we derive the expressions of the self-damping and induced-damping lengths. We also point out the existence of a new damping effect, that we call \\textit{mixed-damping}. The latter is a special case of induced damping. It describes the damping acquired by Dark Matter particles that would be coupled to a species which is free-streaming. In Section \\ref{sec:sdfsdamp}, we derive the constraints on the Dark Matter properties from the calculation of the free-streaming and self-damping scales.  For the  specific case of Dark Matter coupled to neutrinos or photons, the calculation of the induced-damping lengths and the derivation of the associated constraints on interaction rates and cross-sections are done in Sections \\ref{sec:idnudamp} and  \\ref{sec:idphotdamp} respectively.  The definitions of the quantities used in this paper are given in Appendix \\ref{app:def}.  We derive in Appendix \\ref{app:denssym} expressions for the evolution of the Dark Matter density, slightly simplified to be systematically used in our analytical calculations. The interaction rates appearing in the present work are given in the same appendix.  The regions of the Dark Matter parameter space (Dark Matter particles' mass and interaction rates) where the various damping lengths take different forms are explicitly given in Appendix \\ref{app:bord} . Finally, in Appendix \\ref{app:anndamp}, we display analytical expressions for the damping scales and for the associated limits on the Dark Matter interaction rates and cross-sections. ",
        "conclusions": "In this general evaluation of damping properties, we have developed an approach which allows to discuss systematically all possible kinds of   Dark Matter.  We have written the transport coefficients for a composite fluid of species in terms of the interaction rate of each of these species with the medium. This is rather straightforward for the shear viscosity and can be readily extended to heat conduction. We have taken some care to discuss bulk viscosity along the same lines, although the latter turns out for the sake of the present paper  not to generate stronger constraints on the collision rates and cross-sections than shear viscosity. This provides a new  insight into the physical processes at work, and has been used as an opportunity to correct some inaccuracies and wrong statements in earlier papers. We finally wrote all the transport coefficients of a composite fluid as a sum of  contributions from the individual components, taking the coupling to the whole medium into account. This is the key step  which allows us to {\\it discuss the damping of all possible kinds of   Dark Matter}.  We are led to separate, following \\cite{bfsL}, the contribution due to \\emph{self-damping}, that is the transport by the   Dark Matter   itself, and  the one due to the \\emph{induced-damping} which is the transport by other species effective for damping the   Dark Matter   fluctuations. The classical examples of which being neutrinos and photons.  This work aims to put our previous results on a general and firm basis.     Writing that each of these contributions alone must not produce prohibitive damping, we were able to get \\emph{necessary} conditions separately for  the   Dark Matter   interaction rate with all coupled species, as well as for its specific interaction with each of the species it is coupled to. More generally, we identify the parameter space of relevant, namely interaction rate and   Dark Matter   particles' mass. Another parameter, which is seen to appear at the proper places, comes into play. The latter ($\\kappa_{dm}$ in the text) measures the possible offset of the   Dark Matter   and photon temperatures. It is in practice of order unity or somewhat larger, but it must be stressed that in special circumstances (e.g. axion-like   Dark Matter) this parameter may get overwhelmingly large and provide the dominant physics.     As expected, the bounds due to self-damping are closely related to the damping due to   Dark Matter   free-streaming. Both bear similar properties: the velocity of the Dark Matter particles, that is their mass, is the most important parameter. However, there are regions of parameter space where the Dark Matter interaction plays a role. As a result, in these regions, the expression of the free-streaming length (expressed in term of the Dark Matter mass), which is generally used for WDM candidates, does not hold anymore. Instead the free-streaming length depends on both the mass and interaction rate of the Dark Matter candidate. Said differently, the well-known requirement of having   Dark Matter   particles with a mass above $\\sim 1 keV$ is valid only in a restricted region of parameter space, which covers only a  fraction of the Dark Matter   models, where the Dark Matter interactions are really negligible. In regions where the Dark Matter interaction rate is not too large, the smallest acceptable value of the Dark Matter mass is above $1 MeV$. In other parts, where the interaction rate turns out to be very large,  extremely small masses (notably in the $ eV$ range) are allowed.  The bound due to the damping induced by a given species merely depends on the interaction rate of the   Dark Matter with this species, as well as on the species' particle velocity. Relativistic species are  the more efficient, damping the   Dark Matter   fluctuations nearly over the horizon size.  The neutrino induced-damping yields quite a rich  parameter-space diagram (Dark-Matter-neutrino interaction rate/ Dark Matter   particles' mass). The bounds due to the possible coupling of the   Dark Matter   to the neutrinos turn out to be quite novel. Indeed, one in particular refers to Dark Matter that is coupled to neutrinos which are already free-streaming! We have evaluated analytically this effect in what we have called the \\emph{mixed-damping}. The difference with the \\emph{collisional bound}, which is only a lower bound to the damping length but takes only into account the effect of standard collisions, is large. Our mixed-damping estimate is still rather rough, although quite natural a result. It definitely would deserve a check by numerical simulations. The difference with the {collisional bound} is so large that there should be a substantial effect even in a realistic calculation, and  the damping of a collisional fluid coupled to the free-streaming neutrinos should be the dominant contribution as we have stressed here. This is important, especially in the context where Warm   Dark Matter  scenarios. Indeed, we saw that, when our constraints are marginally satisfied, the damping induced by the neutrinos could yield a collisional Dark Matter candidate.   The bound due to a possible coupling of   Dark Matter   with photons (this may be the case through higher order reactions involving charged particles even if the   Dark Matter   is neutral), in the (Dark-Matter-photon interaction rate/ Dark Matter   particles' mass) parameter space, are somewhat less constraining. They are nevertheless relevant for a large fraction of the parameter space. Again, the possibility for Warm   Dark Matter scenarios, where the damping is not due to free-streaming, but to interactions, exists.  The major (and nearly the only) hypothesis is that the   Dark Matter   went through a phase of statistical equilibrium, at some stage of its evolution. There is also the (milder) assumption  that a sudden decrease by many orders of magnitude in the   Dark Matter particle density is not possible elsewhere than at its non-relativistic transition (for instance a phenomenon like the recombination for the electrons -where the latter are substracted from the medium at a temperature bearing no relation with the electron mass- is assumed not to occur for   Dark Matter  ).  Our bounds on the   Dark Matter   interactions with the most relevant species have been calculated  without any assumption on the form of the primordial spectrum. The calculations are valid for any Dark Matter interaction rate and any Dark Matter particle mass, the values of practical interest covering 20 orders of magnitude for the interaction rates or cross-sections and much over 13 orders of magnitude for the Dark Matter particles' mass.  Our estimates are valid within factors of order unity, a procedure which provides important simplifications in the classification. A thorough discussion of all possibilities allows to place any scenario within the relevant parameter space, with a precise statement on the important parameters and the actual physical process at work. The reader interested in a specific scenario can easily work out the relevant bound with more accuracy. He/she will find here a useful estimate of the result and a guide on how to proceed. The present results provide, we believe, a new starting basis for a systematic discussion on the possible nature of   Dark Matter. The astrophysical relevance of these findings, including other constraints than just the damping condition, or more restrictive assumptions on the initial conditions, in particular those of (\\cite{brhs})  which specifically restrict themselves to CDM initial conditions in order to yield more stringent limits, will be thoroughly discussed in Paper II."
    },
    "0410/astro-ph0410194_arXiv.txt": {
        "abstract": "The efficiency of a rapidly spinning Kerr black hole to turn accretion power into observable power can attain 32 percent for the photon emission from the disk, as is well known, following the work of Novikov-Page-Thorne. But many accretion disks are now understood to be underluminous ($L<L_{Edd}$), while still putting large amounts of energy into the jet. In this case, the apparent efficiency of jets driven by the innermost accretion disk of a highly rotating Kerr black hole ($a_*=0.999999$) can reach 96 percent. ",
        "introduction": "The idea of extracting energy from black holes has been proposed by Penrose thirtyfive years ago and followed by a cascade of energy extraction models. Nowadays it is known that from almost all Active Galactic Nuclei (AGNs) which harbour a supermassive ($M>10^5\\times M_{\\odot}$) Kerr black hole, focused jets of hot gas shoot into space at relativistic speed. \"How much energy can these jets get from the black hole or its accretion disk?\" is a question which arises by itself. ",
        "conclusions": "In this short paper we have tried to describe the energy, mass and angular momentum of a black hole - accretion disk - jet system. We find that the efficiency of such a system putting energy from the accretion into the jet is limited by Thorne's (Thorne 1974) bound of 42 percent, but including the possible transfer of angular momentum and energy from the black hole to the inner disk and then to the jets can temporarily increase the apparent efficiency to near 100 percent. Thus, observed jets might represent such stages of apparent high efficiency."
    },
    "0410/astro-ph0410477_arXiv.txt": {
        "abstract": "Cosmological shock waves are a ubiquitous consequence of cosmic structure formation. They play a major role in heating baryonic gas in the large scale structure of the universe. In this contribution we study the shock-heated gas in connection with shocks themselves, using a set of N-body/hydrodynamic simulation data of a $\\Lambda$CDM universe. The distributions of shock speed and temperature of shock-hated gas both should reflect the depth of gravitational potential of associated nonlinear structures, so their morphology. From their distributions we find, in addition to hot gas in and around clusters/groups and warm-hot intergalactic medium (WHIM) with $T=10^5-10^7 \\K$ mostly in filaments, there is a significant amount of low temperature WHIM with $T < 10^5\\K$ distributed mostly as sheet-like structures. The low temperature WHIM was heated and collisionally ionized mainly by shocks with $v_{sh} \\la 150 \\kms$, while photo-ionization by the UV and X-ray background radiations is important for metal ions. Not only the WHIM with $T=10^5-10^7 \\K$ but also the WHIM with $T < 10^5\\K$ make up a significant fraction of gas mass, implying the low temperature WHIM could be important in resolving the missing baryon problem. The shock-heated gas in filaments and sheets are manifested best through emissions and absorptions in soft X-ray and far UV. We confirm that the WHIM with $T=10^5-10^7 \\K$ makes significant contributions to the soft X-ray background, absorptions of highly ionized species such as OVII and OVIII in AGN spectra, and line emissions from OVII and OVIII ions, as pointed by previous studies. However, the WHIM with $T < 10^5\\K$ is the major contributor to absorptions of lower ionized species such as OV and OVI, because these photo-ionized ions are most abundant in sheets of low density and temperature. On the contrary, lines of OV (630 \\AA) and OVI (1032 \\AA) are emitted mostly from the WHIM with $10^5 < T < 10^6\\K$, because they are from collisionally excited ions. ",
        "introduction": "In the current paradigm of the cold dark matter cosmology, the large scale structure of the universe has formed through hierarchical clustering of matter. Deepening of gravitational potential wells caused baryonic matter to accrete onto nonlinear structures. It was postulated that the baryonic matter behaved like a gas and ``collisionless shocks'' formed in the tenuous, magnetized, cosmic plasma via collective electromagnetic viscosities \\citep{krj96,qis98,mrk00,gp03,rkhj03}. In fact, the gravitational energy of the baryonic gas associated with structure formation was dissipated into heat primarily via such shocks. The dissipated energy, then, governs the thermal history and nature of the gas in the universe, and manifests the large scale structure as well as its dynamics through emitted radiation \\citep{kcor94,ckor95,co99a,dcob01,vss02,fl03}.  In a recent work, \\citet{rkhj03} studied the properties of ``cosmological shock waves'' and their roles on thermalization of gas and acceleration of nonthermal particles in the large scale structure, using cosmological N-body/hydrodynamic simulations. The cosmological shocks were classified into two types. ``External shocks'' form around outermost surfaces that encompass nonlinear structures, so they are by nature accretion shocks that decelerate the never-shocked intergalactic gas infalling toward sheets, filaments and knots. ``Internal shocks'' are produced within those nonlinear structures by infall of previously shocked gas to filaments and knots, during subclump mergers, as well as by chaotic flow motions. External shocks have high Mach numbers of up to $M \\sim 100$ due to low temperature of the accreting gas. Internal shocks, on the other hand, have mainly low Mach numbers of $M \\sim$ a few, because the gas inside nonlinear structures has been previously heated by shocks (external and/or internal) and so has high temperature. However, internal shocks are more important in energetics, because of higher preshock density. As a result, thermalization of gas and acceleration of cosmic rays occurred mostly at internal shocks.  The intergalactic gas was heated mostly by cosmological shock waves. The shock-heated gas with $T=10^5-10^7 \\K$ was studied previously, utilizing cosmological hydrodynamic simulations \\citep{co99a,dcob01}. This component, which is called a warm-hot intergalactic medium (WHIM), was estimated to contain $\\sim 30 \\%$ of baryon mass at the present epoch and resides mostly in filaments with a median overdensity of $\\sim 10-30$. Such WHIM can count for a significant fraction of the missing baryons at low redshifts \\citep{fhp98,fuk03}. X-ray emission from the WHIM may contribute to a fraction of the soft X-ray background radiation (XBR) according to \\citet{poc01}. While there were recent reports of observations of the WHIM around clusters \\citep[see, \\eg][]{fbh03,zmmg04b}, detection of such emission still remains a challenging task due to low surface brightness. On the other hand, the WHIM may be better detected through absorption systems and emission lines in soft X-ray and far UV. For instance, the possibility of probing the WHIM with OVI absorbers was explored by \\citet{cen01} and \\citet{fb01}. Absorption systems due to other Oxygen and metal ions in the intervening WHIM were studied by \\citet{hgm98}, \\citet{cwkd03} and \\citet{vbcm03}. Detections of such absorbers in X-ray observations including Chandra were reported \\citep[see, \\eg][among many]{nica03,mymg03}. Emission lines from metal ions in the WHIM were studied by \\citet{yyso03}, \\citet{fssh03}, and \\citet{fcsh03}. A possible detection of such emission lines from a filament around Coma was reported \\citep{fbh03}.  In this paper, we study the shock-heated gas in the large scale structure of the universe and evaluate its observational manifestations with data from a set of cosmological hydrodynamic simulations where shock waves were identified \\citep[see][]{rkhj03}. Specifically, the properties of shock-heated gas are analyzed in connection with those of cosmological shock waves. From the shock speed and gas temperature distributions, we see that 1) the shocks with $v_{sh} \\ga 700 \\kms$ and the hot gas with $T > 10^7\\K$ are distributed in and around clusters/groups, 2) the shocks with $150 \\la v_{sh} \\la 700 \\kms$ and the WHIM with $T=10^5-10^7 \\K$ are mostly in filamentary structures, and 3) the shocks with $v_{sh} \\la 150 \\kms$ and the WHIM with $T < 10^5\\K$ are mostly in sheet-like structures. Here we refer the uncollpased, collisionally ionized, shock-heated gas as the WHIM and extend the low temperature bound of the WHIM to well below $10^5$K. While the WHIM with $T=10^5-10^7$ K has been studied extensively in many previous works, the ``low temperature WHIM'' with $T<10^5$K has not received much attention and will be the main subject of this study.  In the next section, simulation data are described along with numerical details. The properties of shock-heated gas are described in \\S 3, and its observational manifestations are discussed in \\S 4. Summary is followed in \\S 5. ",
        "conclusions": "Hierarchical clustering induces cosmological shock waves in the course of large scale structure formation in the universe \\citep{rkhj03}. The intergalactic gas was heated mostly by such shocks \\citep{co99a,dcob01,vss02}. In this paper, we studied the properties, spatial distribution, and possible observational manifestations of the shock-heated gas in filaments and sheets with a set of simulation data of a $\\Lambda$CDM universe from a grid-based N-body/hydrodynamic code \\citep{rokc93}. The nature of fixed grid-spacing makes such data suitable to the study of the shock-heated gas in filaments and sheets. Shock-heating was focused, so feedbacks from stars, galaxies and AGNs, and photo-ionization heating were ignored, although the photo-ionization of metal ions by the background radiation were taken into account in post-processing data analyses. We expect that exclusion of such processes should not have weakened significantly the main results of this work, since the gas in filaments and sheets was heated mostly by cosmological shocks.  The speed of cosmological shock waves reflects the depth of gravitational potential of the associated nonlinear structures, so their morphology. We saw that 1) the shocks with $v_{sh} \\ga 700 \\kms$ are distributed mostly around and inside clusters/groups, 2) the shocks with $150\\kms \\la v_{sh} \\la 700 \\kms$ are mostly around and inside filamentary structures, and 3) the shocks with $v_{sh} \\la 150 \\kms$ are mostly around and inside sheet-like structures. The distribution of the shock-heated gas should be closely related to that of shock waves. We found that the WHIM with $T=10^5-10^7 \\K$ is distributed mostly in filamentary structures, while the low temperature WHIM with $T < 10^5$K is mostly in sheet-like structures. The hot gas with $T > 10^7$K resides on knot-like structures as the intra-cluster/intra-group medium and the low density medium around them. The amount of shock-heated gas was estimated as follows: the hot component with $T > 10^7 \\K$ (excluding the gas inside clusters/groups) accounts for $\\sim 2.2\\%$ of total gas mass, the WHIM with $T=10^5-10^7 \\K$ for $\\sim 24\\%$, and the WHIM with $T < 10^5 \\K$ for $\\sim 13\\%$ in our highest resolution simulation. Thus our results indicate the existence of the WHIM with $T < 10^5 \\K$, which was heated by shocks of low speed with $v_{sh} \\la 150 \\kms$ and is distributed as sheet-like structures. We suggest that {\\it the low temperature WHIM with $T < 10^5 \\K$ contributes to a significant fraction of the missing baryons at low redshifts} \\citep{fhp98,fuk03}.  It has been suggested in previous studies (see \\S 1 for references) that the shock-heated gas in filaments and sheets can be manifested through emissions and absorptions in soft X-ray and far UV. 1) We found that the soft X-ray emission from the shock-heated gas in our highest resolution data contributes to $\\sim 30 \\%$ of the observed XBR in the $0.5-1$ keV range. This is somewhat larger than that previously suggested by \\citet{poc01}. 2) Column densities of Oxygen ions along randomly selected line paths in the computational box were calculated. Such column densities can produce the absorption systems or X-ray forests in the spectra of distant quasars and AGNs \\citep[see, \\eg][]{hgm98}. The photo-ionization of Oxygen ions by the background UV and X-ray radiations is important for the gas in filaments and sheets due to low density and temperature. As a result, the fractions of Oxygen ions depend strongly on the local gas density. We found that, for OVII and OVIII, the absorption systems associated with the WHIM with $T=10^5-10^7 \\K$ account for $70 - 80 \\%$ of the identified absorption systems, while the low temperature WHIM with $T < 10^5\\K$ contributes $\\sim 55 - 60 \\%$ of the OVI and OV absorption systems. 3) Finally, the emission lines from Oxygen ions were calculated. They are emitted mostly from collisionally excited ions. We estimated that $\\sim 95 \\%$ of the OVII line (574 eV) and OVIII line (653 eV) emissions come from the WHIM with $10^6 < T < 10^7 \\K$, while $\\sim 75 \\%$ of the OVI line (1032 \\AA) and $\\sim 80 \\%$ of the OV line (630 \\AA) are emitted by the WHIM with $10^5 < T < 10^6 \\K$. Hence, we conclude that {\\it absorption systems in UV due to species such as OV and OVI provide the best chance to detect the low temperature WHIM with $T < 10^5\\K$.}"
    },
    "0410/astro-ph0410688_arXiv.txt": {
        "abstract": "{ Spatially resolved images of the galactic supernova remnant G78.2+2.1 ($\\gamma$-Cygni) in hard X-ray energy bands from 25 keV to 120 keV are obtained with the {\\it IBIS-ISGRI} imager aboard the International Gamma-Ray Astrophysics Laboratory {\\it INTEGRAL}. The images are dominated by localized clumps of about ten arcmin in size. The flux of the most prominent North-Western (NW) clump is (1.7$\\pm$0.4)$\\times$10$^{-11} \\enf$ in the 25-40 keV band. The observed X-ray fluxes are in agreement with extrapolations of soft X-ray imaging observations of \\gc by {\\it ASCA GIS} and spatially unresolved {\\it RXTE PCA} data. The positions of the hard X-ray clumps correlate with bright patches of optical line emission, possibly indicating the presence of radiative shock waves in a shocked cloud. The observed spatial structure and spectra are consistent with model predictions of hard X-ray emission from nonthermal electrons accelerated by a radiative shock in a supernova interacting with an interstellar cloud, but the powerful stellar wind  of the O9V star HD 193322  is a plausible candidate for the NW source as well. ",
        "introduction": "The supernova remnant (SNR) G78.2+2.1 ($\\gamma$-Cygni) is a degree size extended source which has been imaged in radio waves to $\\gamma$-rays.   % The SNR is located in the complex Cygnus X region of  massive gas and dust complexes and close to the most powerful Cyg OB2 association. Radio observations by Higgs \\etal (1977) established a shell-like structure of the remnant, suggested also by Lozinskaya (1977). Multi-frequency radio observations by Zhang \\etal (1997) revealed a patchy structure of the radio spectral-index distribution. The integrated radio spectral index $\\alpha$ = 0.54 $\\pm$ 0.02, with variations of $\\pm$ 0.15 within the remnant. Prominent radio brightness enhancements are present in the South-East and the North-West of \\gc (e.g. Zhang \\etal 1997). A distance to \\gc of $\\sim$ 1.5 kpc was estimated by Landecker \\etal (1980) from radio HI observations.  Optical images of \\gc in H$_{\\alpha}$+[NII], [SII] and [OIII] filters, recently presented by Mavromatakis (2003), clearly show a patchy structure with bright spots of some ten arcminutes scale and line emission fluxes of a few times $\\,10^{-15}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$ for H$_{\\alpha}$ filters and a few times $\\,10^{-16}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$ for the [OIII] filter.  Archival {\\it ROSAT} and {\\it ASCA} observations of G78.2+2.1 were analyzed by Lozinskaya \\etal (2000). They pointed out a complex structure of the remnant with {\\it ROSAT} emission extending well beyond the apparent SNR radio shell, possibly indicating expansion into a progenitor star wind cavity. With archival {\\it ASCA} observations, Uchiyama \\etal (2002) found X-ray emission above 4 keV to be dominated by several localized clumps, mostly in the Northern part of the remnant. The 4-10 keV emission of the  clumps has a hard photon index of 0.8--1.5.  A high-energy \\gr source, 2CG 078+2, was discovered in the field of \\gc with the {\\it COS B} satellite (e.g. Swanenburg \\etal 1981). {\\it CGRO-EGRET} confirmed this source (2EG J2020+4026/3EG J2020+4017), which is one of the brightest steady-state unidentified sources in the {\\it EGRET} catalogue with a flux of $\\sim$ 1.2 $\\times 10^{-6}\\fl$ above 100 MeV (Sturner and Dermer 1995; Esposito \\etal 1996). {\\it Whipple} \\gr telescope observations (e.g. Buckley \\etal 1998) established an upper limit of 2.2 $\\times 10^{-11}\\fl$ for the flux above 300 GeV, indicating a break in the high-energy emission spectrum above a few GeV. The \\gr emission might be produced by a pulsar, but can be attributed to interactions of accelerated energetic particles with ambient matter and radiation as well (e.g. Sturner \\etal 1997; Gaisser \\etal 1998; Bykov \\etal 2000). Solving this problem of the origin of the \\gr emission requires the analysis of multiwavelength imaging observations. Hard X-ray imaging provides a crucial tool to distinguish between leptonic and hadronic contributions to the \\gr emission. We present below the first hard X-ray imaging observation of \\gc.   \\begin{figure} \\includegraphics[width=0.5\\textwidth]{f1.ps} \\caption{The {\\it ISGRI} 25--40 keV broad view of the \\gc region (linearly scaled S/N map of the sources above  2$\\sigma$). The solid circle is the same as in Fig.~\\ref{fig:A_R_I} . } \\label{fig:broad_view} \\end{figure} ",
        "conclusions": "The locations of the hard X-ray clumps detected by {\\it INTEGRAL ISGRI} correlate with dim regions of high \\nh $\\gsim 5\\times 10^{21} \\cms$ in the {\\it ROSAT PSPC-B} 0.4-2.0 keV greyscale map in Fig.~\\ref{fig:A_R_I}. The clumps adjoin the patches of bright optical H$_{\\alpha}$+[NII], [SII] and [OIII] line emission (of $\\gsim$ 10\\arcmin\\ scale size) observed by Mavromatakis (2003). The line emission may indicate an interaction of the SNR with a nearby cloud and the presence of a radiative shock wave.  In evolved supernova remnants interacting with interstellar clouds, such as \\gc\\ and IC~443, a highly inhomogeneous structure consisting of a forward shock of a moderate Mach number, a cooling layer, a dense radiative shell, and an interior region filled with hot tenuous plasma is expected. A SNR evolving in the inter-clump medium of a molecular cloud of density $n_a \\sim 25$ H atoms cm$^{-3}$  becomes radiative at radii $\\sim 6$ pc (Chevalier 1999). A model of nonthermal electron acceleration and interactions in such SNRs by Bykov \\etal (2000) predicted that these SNRs are efficient electron accelerators and sources of hard X- and \\gr emission. A distinctive feature of the model is the presence of emitting clumps with very flat hard X-ray spectra of photon indexes $\\alpha \\sim 1$. The photon break energy and the maximum energy of accelerated electrons depend on the ambient density. A radiative shock propagating through the interclump medium could accelerate radio-emitting electrons in the GeV regime, also producing hard X-ray and \\gr {\\it EGRET}-regime emission. A slower MHD shock in a dense molecular clump could produce bright hard X-ray emission without prominent radio and \\gr counterparts.  Bremsstrahlung radiation from the nonthermal electrons has a hard X-ray photon spectrum. The nonthermal X-ray flux  $F_x$ from the electrons accelerated by a MHD shock of a velocity $v_{s7}$ (in 100 $\\kms$) and of an apparent angular size $\\theta$ (in arcmin) can be estimated as a fraction of the shock ram pressure: \\begin{equation} F_x \\approx 9.2\\times 10^{-12}\\, \\eta\\,\\, \\frac{n_a\\, v_{s7}^3}{2}\\,\\, \\theta^2 \\,\\, \\enf. \\end{equation} The efficiency of X-ray bremsstrahlung emission $\\eta$ is relatively low due to the Coulomb losses of fast particles, providing  powerful optical and IR diffuse emission from the interclump medium. Having in mind that $\\eta < 10^{-5}$ in the ambient medium [of solar abundance] for the energy band below 100 keV, one may conclude that $n_a \\cdot v_{s7}^3 \\cdot \\theta^2 > 10^5$. That is a rather stringent requirement implying a high efficiency of electron acceleration by a relatively slow shock with $v_{s7} \\gsim$ 1 and $n_a > 100 \\cmc$ for $\\theta \\sim$ 10. A more detailed confrontation of the observational data and the model predictions, including GeV  \\gr emission, will be described elsewhere.  An alternative interpretation of the observed hard X-ray sources could be a fast massive ballistically moving ejecta fragment interacting with a cloud. The fragment could drive a shock of velocity about 1,000 $\\kms$ in the inter-clump medium providing a powerful source of X-ray emission both in continuum and lines (Bykov 2003). The fragments being enriched with metals (of an average nuclear charge $\\langle Z\\rangle$) have higher radiation efficiency, since $\\eta \\propto\\, \\langle Z\\rangle$. The energetic problem of the bremsstrahlung emission models $\\langle Z\\rangle \\cdot n_a \\cdot v_{s7}^3 \\cdot \\theta^2 > 10^5$ is alleviated. The line signature around 6.2~keV required to fit the {\\it RXTE PCA} data could be naturally explained in that case as a Fe K line complex predicted for the ejecta fragment model, though a contribution from hidden accreting sources (AGNs or X-ray binaries) cannot be excluded.  A powerful wind of the early type O9V star HD 193322 located just 7\\arcmin\\ from the NW bright clump (see Fig.~\\ref{fig:A_R_I}) could also be considered as a candidate source for the observed emission. The star is the central object of the open cluster Collinder 419 located at an estimated distance 1.4 kpc (e.g. McKibben et al., 1998) which is close to the \\gc distance estimations. The trigonometric parallax of HD 193322 provided by the {\\it Hipparcos} catalogue is (2.1 $\\pm$ 0.61) mas marginally consistent with the distance. Interaction of the wind of HD 193322 with G78.2+2.1 would be a plausible source of particle acceleration and hard X-ray emission.   A deeper exposure of the field with {\\it ISGRI}\\, is required to detect [or to place meaningful upper limits] on the hard emission from the NW, NE, SE regions up to 120 keV. The presence of a few clumps of hard emission at a level above 10$^{-11} \\enf$ at the borders of \\gc would support the hypothesis of SNR origin of the nonthermal clumps with important implications for particle acceleration mechanisms in SNRs."
    },
    "0410/astro-ph0410641_arXiv.txt": {
        "abstract": "We present 0.5--8.0~keV \\chandra\\ observations of six bright quasars that represent extrema in quasar emission-line properties --- three quasars each with small and large blueshifts of the \\CIV\\ emission line with respect to the systemic redshift of the quasars. Supplemented with seven archival \\chandra\\ observations of quasars that met our selection criteria, we investigate the origin of this emission-line phenomenon in the general context of the structure of quasars. We find that the quasars with the largest \\CIV\\ blueshifts show evidence, from joint-spectral fitting, for intrinsic X-ray absorption (\\nh$\\sim10^{22}$\\cmsq).  Given the lack of accompanying \\CIV\\ absorption, this gas is likely to be highly ionized, and may be identified with the shielding gas in the disk-wind paradigm. Furthermore, we find evidence for a correlation of \\alphauv, the ultraviolet spectral index, with the hardness of the X-ray continuum; an analysis of independent Bright Quasar Survey data from the literature supports this conclusion. This result points to intrinsically red quasars having systematically flatter hard X-ray continua without evidence for X-ray absorption. We speculate on the origins of these correlations of X-ray properties with both \\CIV\\ blueshift and \\alphauv\\ and discuss the implications for models of quasar structure. ",
        "introduction": "Considering the dynamic range of black hole masses and luminosities, quasar emission-line phenomenology is remarkably consistent.  At the same time, the physical drivers for the known range of emission-line properties are poorly understood.  One of the reasons the emission-line region is not better understood is because of the general similarity of quasar spectra \\citep{rfs+01} and the lack of sufficient emission-line distinctions to test competing models.  That said, there are a few well-known differences.  For example, it is generally believed that the widths of the Balmer lines (and probably \\ion{Mg}{2}) are related to the mass of a quasar's central black hole \\citep[e.g.,][]{ves02} --- although the presumed disk-like configuration of this gas means that the line widths will be affected by projection effects \\citep[e.g.,][]{kro01}.  There is also the observation that more luminous quasars tend to have \\CIV\\ emission lines with smaller equivalent widths, otherwise known as the \\citet{bal77} effect.  Recently, \\citet{bl04} have argued that the Baldwin effect and the dynamics of the broad emission-line region (BELR) in general are driven by differences in the Eddington ratio, $L/L_{\\rm Edd}$, among quasars.  Similarly, \\citet{bg92} and others have proposed that $L/L_{\\rm Edd}$ drives a third trend, mainly the anti-correlation between the strengths of \\OIII\\ and \\ion{Fe}{2} emission in the optical part of the spectrum. A fourth effect seen in quasar emission lines is emission-line blueshifting \\citep[e.g.,][]{gas82,wil84,tf92}, where higher ionization lines yield emission-line redshifts that are systematically too small --- as if the lines had been shifted blueward.  \\citet{ric+02} recently presented a summary of the emission-line blueshift\\footnote{The blueshift is defined as the velocity offset between redshifts derived separately from \\CIV\\ and \\MgII, given in velocity units with more positive velocity indicating larger emission-line blueshifts.} effect among a sample of 3814 quasars from the Sloan Digital Sky Survey (SDSS; \\citealt{yor+00}).  They showed that these emission-line shifts are more ubiquitous than previously thought, with the average quasar having a \\CIV\\ emission line that yields a redshift that is too small with respect to \\MgII\\ by \\mbox{$\\sim800$\\kms}. Furthermore, quasar composite spectra created from samples binned by \\CIV\\ blueshift showed differences in their ultraviolet continua and other emission-line properties.  Specifically, the composite spectrum of the largest-blueshift quasars had a significantly bluer than average ultraviolet continuum and weaker \\CIV\\ emission, while the smallest-blueshift composite was notably redder with stronger \\CIV\\ emission.  \\citet{ric+02} showed that this difference resulted  from the intrinsic continuum color of the quasars, rather than as an effect of reddening.  The distribution of \\CIV\\ blueshifts ranges over 3000\\kms, making this effect an excellent candidate for further investigation into the structure of the BELR. \\citet{ric+02} hypothesized that the range of \\CIV\\ blueshift and ultraviolet continuum color was a consequence of orientation. Orientation could be understood in one of two ways, either with respect to the accretion disk, or with respect to the wind, which may have a range of opening angles (for example, see Figure~7 in \\citealt{elv00} and Figure~5 in \\citealt{ric+04}). In both cases, understanding the BELR might be aided by probing smaller scales. X-ray emission is believed to arise on physical scales comparable to and smaller than the optical-UV continuum. Therefore, X-rays are not only sensitive to physical conditions in the immediate vicinity of the black hole, they can also probe the line of sight to the observer from the inner regions of the accretion flow. X-ray studies thus offer the potential to tie BELR phenomena (on scales of light days) to fundamental properties of the accretion disk.  For example, if the range of BELR phenomenology results purely from orientation effects, one might expect the intrinsic X-ray continuum properties to be unrelated.  To investigate this avenue, we undertook an exploratory \\chandra\\ \\citep{chandra_ref} survey to study the X-ray properties of extreme examples of the \\CIV\\ blueshift distribution. We were awarded six observations in Cycle~4; these observations have been supplemented with seven from the archive. In \\S\\ref{sec:data}, we present the target selection process and describe the initial \\chandra\\ data analysis.  We study the trends in these data in \\S\\ref{sec:trends}.   The results from joint spectral fitting of the \\chandra\\ spectra and the connection to optical/UV emission-line and continuum properties are presented in \\S\\ref{sec:joint}. In \\S\\ref{sec:s}, we summarize the results from our X-ray analysis, and introduce comparison data from the literature on Bright Quasar Survey \\citep[BQS;][]{bqs} objects. In \\S\\ref{sec:d}, we discuss the possible physical origins of our observed correlations and comment on the relative importance of external orientation (tilt of the accretion disk with respect to our line of sight), $L/L_{\\rm Edd}$, and internal orientation effects (such as the disk-wind opening angle --- possibly driven by $L/L_{\\rm Edd}$).  The final section briefly outlines our conclusions.  The cosmology assumed throughout has $H_{\\rm 0}=70$\\kms\\,Mpc$^{-1}$, $\\Omega_{\\rm m}=0.3$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "  \\begin{itemize} \\item{Those quasars in our sample with \\CIV\\ blueshifts $\\gtrsim1100$\\kms\\ show evidence from joint-fitting of the X-ray spectra for intrinsic X-ray absorption with \\nh\\,$\\sim10^{22}$\\cmsq. We interpret the presence of X-ray absorption in the large-blueshift sample as support for the orientation interpretation of the \\CIV\\ blueshift put forth by \\citet{ric+02} whereby large-blueshift quasars are seen at inclination angles close to the line of sight through the wind (which may have a range of opening angles).  This result is broadly consistent with the disk-wind model of quasar broad-line regions of \\citet{mcgv95}.}  \\item{We find that there is a trend of steeper hard-band X-ray continua with bluer \\alphauv\\ spectral index in our sample; this result is supported by a complementary analysis of independent Bright Quasar Survey data from the literature.  We find the Eddington accretion rate, $L/L_{\\rm Edd}$, to be a likely candidate for the primary physical driver of this trend. }  \\end{itemize}  Extending this study to larger samples with higher-quality X-ray data is certainly warranted to test these claims. Specifically, given the lack of \\CIV\\ absorption in the SDSS spectra of the large-blueshift sample, we predict that high quality X-ray spectroscopy will reveal this gas to be ionized with $\\xi>8$\\,erg\\,cm\\persec. With the large effective area and 0.3--10.0~keV bandpass, sensitive \\xmm\\ observations could significantly constrain both the column density and ionization state of the absorption for each quasar in the large-blueshift sample. Furthermore, high signal-to-noise ratio X-ray observations of a quasar sample tuned to isolate the relationship between UV/optical and X-ray continua hold promise for understanding the effect of Eddington accretion rate on quasar spectral energy distributions."
    },
    "0410/astro-ph0410531_arXiv.txt": {
        "abstract": "Young star clusters are often born with such high stellar densities that stellar collisions play an important role in their further evolution. In such environments the same star may participate in several tens to hundreds of collisions ultimately leading to the collapse of the star to a black hole of intermediate mass. At later time the black hole may acquire a companion star by tidal capture or by dynamical --3-body-- capture. When this companion star evolves it will fill its Roche-lobe and transfer mass to its accompanying black hole. This then leads to a bright phase of X-ray emission but the binary may also be visible as a source of gravitational wave radiation in the LISA band. If the star captured by the intermediate mass black hole is relatively low mass ($\\aplt 2$\\,\\msun) the binary will be visible as a bright source in gravitational waves and X-rays for the entire lifetime of the binary.  The majority of compact binaries which formed from the population of primordial binaries in young and dense star clusters do not lead to detectable gravitational wave sources. ",
        "introduction": "\\label{Sect:introduction}   Dense and young star clusters are rather rare in our Galaxy.  The six such star clusters that are currently known are: Arches \\cite{2002ApJ...581..258F}, IRS13 \\cite{2004A&A...423..155M} and IRS16 \\cite{2003ApJ...593..352P,2004ApJ...602..760E}, NGC3603 \\cite{1999A&A...352L..69B}, Quintuplet \\cite{1999ApJ...514..202F} and Westerlund~1 \\cite{2000ApJ...533L..17V}.  The two IRS sources, however, can hardly be called a cluster as they look more like the debris of disrupted star cluster \\cite{2000ApJ...545..301K,2003ApJ...596..314M}.  Young and dense star clusters (hereafter YDC, sometimes pronounced as YoDeC) are characterized by a stellar population which is younger than about 10\\,Myr.  The fact that these clusters are young means that stars of all masses are still present, offering critical insights into the stellar initial mass function and cluster structural properties at formation.  The term ``dense'' means that dynamical evolution and physical collisional processes can operate fast enough to compete with and even overwhelm stellar evolutionary timescales; dense stellar systems are places where wholly new stellar evolution channels can occur, allowing the formation of stellar species completely inaccessible by standard stellar and binary evolutionary pathways.  There are no examples of the older ($\\apgt10$\\,Myr) siblings of these clusters in the Milky Way Galaxy. For the Arches and Quintuplet this may not be so surprising, as they dissolve in the tidal field of the Galaxy in at most a few tens of million years \\cite{2000ApJ...545..301K,2002ApJ...565..265P}, and even if they survive longer they will become hard to detect against the dense background stellar population \\cite{2001ApJ...546L.101P}. For the two rather isolated clusters, NGC\\,3603 and Westerlund~1, it is rather curious that there are no examples of their 10 to 100\\,Myr descendants in the Milky Way.  The other characteristic of YDCs is the high density. The few clusters near the Galactic center must have a high density, as otherwise they would be disrupted easily by the strong tidal field. The clusters further out however, do not have this constraint and it is interesting to note that clearly these clusters ware born without much knowledge of the local potential of the Galaxy, i.e: they behave as isolated clusters.  In this paper, we discuss the formation and evolution of gravitational wave sources in the form of compact binaries in YDCs, but for convenience we extend our definition to star clusters of 100\\,Myr. The basis of this study is the star cluster MGG11 \\cite{2003ApJ...596..240M} at about 200\\,pc from the nucleus of the starburst galaxy M82 at a distance of about 3.6\\,Mpc.  The star cluster MGG-11 has an age of about 7--12\\,Myr \\cite{2003ApJ...596..240M}, a line-of-sight velocity dispersion of $\\sigma_r = 11.4\\pm0.8\\,\\kms$ and a projected half-light radii, $r = 1.2$\\,pc. The cluster mass then totals $\\sim 3.5 \\times 10^5 \\msun$. The cluster mass function seems to be deficient of stars below about 1\\,\\msun, but follow a Salpeter slope for the higher masses.  We simulate this cluster by integrating the equations of motion for all stars and follow the evolution of stars and binaries in the cluster using the {\\tt starlab}\\footnote{{\\tt http://www.ids.ias.edu/$\\sim$starlab/}} \\cite{2001MNRAS.321..199P} direct N-body package.  We start our simulations with 144179 stars, 13107 of which are in binaries totaling a binary fraction of about 10\\%.  The binary parameters ware selected as follows: first we chose a random binding energy between a minimum of 10kT and a maximum, and we chose the orbital eccentricity from the thermal distribution. The maximum binding energy was selected such that the distance at pericenter exceeded four times the radius of the primary star. For the initial density profile we adopt a King model with $W_0=12$ \\cite{1966AJ.....71...64K}. We ignored an external tidal field, but stars are removed from the simulation if they are more than 10 tidal radii (about 80\\,pc) away from the density center of the cluster. ",
        "conclusions": "Young and dense star clusters are promising sources for gravitational wave radiation. Especially when they get older than $\\apgt 50$\\,Myr as white dwarf binaries start to form. But also at younger age the presence of an intermediate mass black hole may provide an important source of gravitational waves.  An interesting possibility is provided by a $\\sim 2$\\,\\msun\\, donor which starts to transfer mass to an intermediate mass black hole at birth. Such a binary may be visible as a bright transient X-ray source and simultaneously as gravitational wave source in the {\\em LISA} band to a distance of a few kpc."
    },
    "0410/astro-ph0410412_arXiv.txt": {
        "abstract": "We review the properties of the currently known CV population from an accountants point of view. In particular, it is examined to what extent different selection mechanisms (variability, X-ray emission, colours, spectroscopic properties) affect the observed CV orbital period distribution. Large-scale surveys provide homogeneously selected CV samples, for which observational biases can in principle be modelled. Such samples will eventually provide the essential observational input for an improved understanding of CV evolution, and we highlight the impact of several completed and on-going surveys. ",
        "introduction": "Historically, cataclysmic variables (CVs) were discovered because of their \\textit{variability}, predominantly being large-amplitude variables such as dwarf novae or classical novae, e.g. U\\,Geminorum \\citep{hind56-1}, but also some irregular low-amplitude systems such as the polar prototype AM\\,Herculis \\citep{wolf24-1}. In 20th century, the rate of discovery of CVs was rather slow until the mid-seventies, \\citet{warner76-1} discussed the properties of some 27 systems with known orbital periods, and noted the following points: a strong preference for short-period CVs and a deficiency of CVs in the orbital period range $1.5-3.25$\\,h. It became rapidly clear that the orbital period distribution of CVs is closely related to the evolution of these systems, more specifically to the rate at which CVs lose orbital angular momentum \\citep[e.g.][]{eggleton76-1}. To our knowledge, the first orbital period histogram was published by \\citet{whyte+eggleton80-1}, based on data for 33 CVs, clearly showing the presence of a ``period gap'' between $2-3$\\,h. Both the number of discovered CVs and the amount of detailed follow-up studies have rapidly increased since 1980, stimulating growing activity on the topic of CV evolution. \\textit{Disrupted magnetic braking} became the standard scenario of CV evolution explaining the presence of the period gap in the context of an abrupt change in the rate of orbital angular momentum loss at $P_\\mathrm{orb}\\simeq3$\\,h (\\citealt{rappaportetal83-1, paczynski+sienkiewicz83-1, spruit+ritter83-1}, see also \\citealt{king88-1, howelletal01-1}).  While no globally accepted alternative to the standard model exists so far, it has been criticised on the base that (a) many of its predictions are in disagreement with the observations and (b) it is based on an ad-hoc assumption about the mechanism of orbital angular momentum loss. The unsettled nature of the field of CV evolution is reflected in the number of alternatives/modifications to the standard model that have been published in the recent years \\citep[e.g.][]{clemensetal98-1, mccormick+frank98-1, schenkeretal98-1, kolb+baraffe99-1, spruit+taam01-1, taam+spruit01-1, kolbetal01-2, kingetal02-1, king+schenker02-1, schenker+king02-1, schenkeretal02-1, andronovetal03-1, barker+kolb03-1, taametal03-1, ivanova+taam04-1}.  The purpose of this paper is to review the properties of the known CV population that currently provides the observational input for CV evolution theory, and to assess the impact that several large-scale surveys had/have on this field.  \\begin{figure} \\includegraphics[angle=-90,width=0.5\\textwidth]{gaensicke_f1a.ps} \\includegraphics[angle=-90,width=0.5\\textwidth]{gaensicke_f1b.ps} \\caption{\\label{f-allcvs}\\textit{Left:} The distribution of the known CVs from \\citet{downesetal01-1} in galactic coordinates. The galactic centre wraps around at the right/left edge of the plot, the solid line indicates declination $\\delta=0$. CVs with a known orbital period are shown as black dots, those without as grey dots. \\textit{Right:} The period distribution of the 531 CVs listed by \\citet{ritter+kolb03-1}, V7.3. The $2-3$\\,h period gap is indicated in grey.} \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410138_arXiv.txt": {
        "abstract": "We present the results of microlens ray-tracing simulations showing the effect of absorbing material between a source quasar and a lensing galaxy in a gravitational lens system. We find that, in addition to brightness fluctuations due to microlensing, the strength of the absorption line relative to the continuum varies with time, with the properties of the variations depending on the structure of the absorbing material. We conclude that such variations will be measurable via UV spectroscopy of image A of the gravitationally lensed quasar Q2237+0305 if the Lyman Alpha clouds between the quasar and the lensing galaxy possess structure on scales smaller than $\\sim 0.1$ pc. The time scale for the variations is on the order of order years to decades, although very short term variability can occur. While the Lyman alpha lines may not be accessible at all wavelengths, this approach is applicable to any absorption system, including metal lines. ",
        "introduction": "The Lyman Alpha (Ly$_\\alpha$) forest is an absorption phenomenon which is observed in the spectra of many high redshift quasars. It arises due to absorption by galactic and intergalactic neutral hydrogen between a quasar and an observer [see \\citet{1998ARA&A..36..267R} for a recent review]. As the universe expands, photons from the quasar are redshifted, and therefore absorption by clouds at different cosmological distances cause absorption at different wavelengths in the spectrum. The resulting spectrum contains many absorption lines blueward of the Ly$_\\alpha$ emission line of the quasar, and the structure of the absorbing material along the line of sight is imprinted on the spectrum \\citep[e.g.][]{1996ApJ...472..509L}.  Numerical simulations of cosmological structure formation predict the large scale structure of the gas which can then be compared with observations of quasar spectra \\citep{1999ApJ...511..521D}. The conclusions of these studies are that the Ly$_\\alpha$ ``clouds\" are mostly intergalactic structures, with the denser regions corresponding to the locations where galaxies form. The absorbing gas along the line of sight to a single quasar can be probed by comparing the relative positions and strengths of the absorption lines in the Ly$_\\alpha$ forest in the spectrum of the quasar \\citep{1980ApJS...42...41S}. The structure across the line of sight is harder to measure, although if two quasars that are nearby on the sky are observed, the fraction of Ly$_\\alpha$ forest lines that are common to the two quasars can provide a measure of the extent of the clouds across the sky \\citep{1982ApJ...256..374S, 2003MNRAS.341.1279R}.  Traditionally, the large scale structure of the Ly$_\\alpha$ Forest has been the main focus of research, because of its cosmological significance \\citep[e.g.][]{1999ApJ...518...24M}. Recently, the small scale structure (on kiloparsec scales) has also been investigated by several methods \\citep{1999ApJ...515..500R, 2001ApJ...554..823R, 2001ApJ...562...76R, 2002ApJ...576...45R}; potentially revealing structure in and about the interstellar medium in young galaxies.  Gravitational lensing has previously been used to probe the size of absorbing clouds [note that if the mass associated with a Ly$_\\alpha$ cloud is sufficiently large, it too may result in gravitational lensing effects, see \\citet{2004astroph}]. When a background quasar (regarded as a point source) is lensed by a foreground galaxy, multiple images of the quasar may be obtained. The light rays responsible for the two different images have traversed slightly different regions of space in their journey to the observer. Hence, a measure of the size of the Ly$_\\alpha$ absorbing clouds (down to kiloparsec scales) can be made by cross-correlating the the absorption lines in two images of the same quasar \\citep{2004A&A...414...79E,2003AAS...20311301B}. This method is analogous to using a true pair of quasars, however the lensing supplies a ``pair'' of quasars that are close together on the sky (separated by $\\sim$ arc seconds), allowing smaller scales to be investigated. \\citet{1999ApJ...515..500R} used this method in the multiply imaged quasar Q1422+231 to trace MgII and CIV absorption clouds between the lens and the quasar down to scales of $\\sim$ 26 pc, which is the smallest scale so far considered.  The very small scale ($\\lesssim$ 1 parsec) structure of the Ly$_\\alpha$ clouds is completely unknown (due to limits in observational techniques, and the fact that the resolution of Ly$_\\alpha$ forest formation simulations is too coarse), and would be useful for deciding what the local equivalents of the Ly$_\\alpha$ clouds are. Some (but not all) low redshift Ly$_\\alpha$ clouds have been found to be associated with galaxies \\citep{1995ApJ...441...51S}, suggesting that the structure of the Ly$_\\alpha$clouds might mimic that of the interstellar medium. This is known to have structure on many scales, exhibiting fractal structure due to turbulence \\citep{1997ApJ...477..196E}. In this paper we present simulations demonstrating how the very small, sub parsec scales of Ly$_\\alpha$ clouds can be measured using a quasar that is subject to gravitational microlensing by a foreground galaxy. ",
        "conclusions": "The cosmological distribution of neutral hydrogen gas leaves an imprint on the spectrum of background quasars, in the form of a series of absorption lines. This phenomenon is thought to be related to the large scale flow of gas and formation of galaxies and clusters. While strong constraints have been placed on the larger scale properties of the Ly$_\\alpha$ clouds, virtually nothing is known on smaller scales, where gas may possess similar turbulent properties to the local interstellar medium, or be collapsing to form stars. In this paper, we have examined the influence of gravitational microlensing on the paths taken by light rays impinging on a distribution of Ly$_\\alpha$ clouds lying between the lens and the source. For a single macro-image, the population of micro-images samples differing absorption regions, with the net absorption seen (in a single line) being the average of that experienced by each ray. As the stars move, the configuration of micro-images changes, sampling different regions of the absorption clouds and causing temporal variability of the absorption line strength.  In this study we considered clouds of three different sizes and at three different locations between the lens and the source. Quantitatively, it was shown that the fastest variations in the absorption line strength occur for the smallest clouds, and for clouds close to the source. We also found that, due to the geometry of the universe, the variations in the absorption lines in the spectrum of Q2237+0305 due to clouds at redshifts $\\gtrsim 0.5$ would have approximately the same effect as identical clouds which are just in front of the source. At all redshifts, the caustic structure evident in the microlensing magnification map is also present in the absorption line strength map. These discontinuities can result in very rapid changes in the absorption line strength, with a corresponding change in the continuum brightness. However, a more detailed examination of the of the relationship between the continuum magnification and the line strength shows no apparent correlation between the two.  An analysis of the variability timescales reveals that significant changes in line strength due to the effects of gravitational microlensing would be visible on time scales ranging from years to decades, dependent on the size on the size of the absorption clouds and their location behind the lens; this is at least an order of magnitude faster than clouds simply drifting across the line of sight to a distant quasar. Hence, in order to measure this effect, long term spectroscopic monitoring of the individual images of microlensed quasars is required. Such observations have already been proposed for studies of quasar structure \\citep{2002ApJ...576..640A, 2002ApJ...579..127S, 2004MNRAS.348...24L}.  This study has only considered Ly$_\\alpha$ clouds in image A of Q2237+0305, which is a positive parity image. Other images of Q2237+0305 have negative parity, corresponding to one of the curves in Figure~\\ref{size} passing through zero. Furthermore, this study has assumed that the absorbing clouds possess no kinematic substructure, and hence the resultant variability applies to the total strength of the line; for this a simple equivalent width measure in relatively low resulution spectra would suffice. Nevertheless, any real distribution of gas will exhibit some kinematic substructure in the form of rotation, infall, or simple clumping. This would result in not only variability in the total strength of an absorption line, but also in its profile [see \\citep{1998MNRAS.297...69L}]. To detect line profile variability, higher resolution spectra would be required. We will discuss these issues in future contributions."
    },
    "0410/astro-ph0410624_arXiv.txt": {
        "abstract": "We have argued (e.g. \\cite{1}) that the well-known `ankle' in the cosmic ray energy spectrum, at logE (eV) $\\sim$ 18.7-19.0, marks the transition from mainly Galactic sources at lower energies to mainly extragalactic above.  Recently, however, there have been claims for lower transitional energies, specifically from logE (eV) $\\sim$ 17.0 \\cite{2} via $17.2-17.8$ \\cite{3} to 18.0 \\cite{4}.  In our model the ankle arises naturally from the sum of simple power law-spectra with slopes differing by $\\Delta \\gamma \\sim 1.8$; from differential slope $\\gamma = -3.8$ for Galactic particles (near logE = 19) to $\\gamma \\sim -2.0$ for extragalactic sources. In the other models, on the other hand, the ankle is intrinsic to the extragalactic component alone, and arises from the shape of the rate of energy loss versus energy for the (assumed) protons interacting with the cosmic microwave background (CMB).\\\\ Our detailed analysis of the world's data on the ultra-high energy spectrum shows that taken together, or separately, the resulting mean sharpness of the ankle (second difference of the log(intensity$\\times$E$^3$) with respect to logE) is consistent with our `mixed' model.  For explanation in terms of extragalactic particles alone, however, the ankle will be at the wrong energy -- for reasonable production models and of insufficient magnitude if, as seems likely, there is still a significant fraction of heavy nuclei at the ankle energy. ",
        "introduction": "A key property of the cosmic radiation is the division between those particles of Galactic origin and those which are extragalactic.  It is virtually certain that those below about logE (eV) = 17 are Galactic (G) and that the very highest energy particles are extragalactic (EG), but the energy (E$_{0.5}$) at which EG starts to dominate is debatable.  It is this question that is the subject of this paper.  In every model the Galactic magnetic field progressively loses its trapping power as the rigidity increases so that the mean charge of the detected Galactic particles increases. Ideally, measurements of the mass composition should be possible and these would appear to indicate E$_{0.5}$ but there is the distinct likelihood of a significant fraction of heavy nuclei in the EG beam.  Here, we start by considering the measured spectral shape and then examine the possibility of an explanation in terms of primary protons.  This is followed by an analysis of the mass composition and its relevance to the Galactic/Extragalactic boundary. ",
        "conclusions": "A variety of evidence leads us to believe that the ankle in the primary energy spectrum is due to the rapid transition from Galactic to Extragalactic particles and is not due to Extragalactic particles alone, which are required to predominate above logE = 18.0.  In our model, at logE=19.0, 30\\%-- 40\\% of the particles are still Galactic, although this fraction is falling rapidly with increasing energy.  In `competing models' this fraction is 10\\% for the Hi-Res case \\cite{2}, 2\\% in the `Hillas-model' \\cite{4} and zero in the Berezinsky et al. model \\cite{3}."
    },
    "0410/astro-ph0410554_arXiv.txt": {
        "abstract": "We present observations of CLASS B2108+213, the widest separation gravitational lens system discovered by the Cosmic Lens All-Sky Survey. Radio imaging using the VLA at 8.46~GHz and MERLIN at 5~GHz shows two compact components separated by 4.56~arcsec with a faint third component in between which we believe is emission from a lensing galaxy. 5-GHz VLBA observations reveal milliarcsecond-scale structure in the two lensed images that is consistent with gravitational lensing. Optical emission from the two lensed images and two lensing galaxies within the Einstein radius is detected in {\\it Hubble Space Telescope} imaging. Furthermore, an optical gravitational arc, associated with the strongest lensed component, has been detected. Surrounding the system are a number of faint galaxies which may help explain the wide image separation. A plausible mass distribution model for CLASS B2108+213 is also presented. ",
        "introduction": "\\label{intro}  The Cosmic Lens All-Sky Survey (CLASS; \\citealt{myers02,browne02}) is the largest statistically complete gravitational lens survey ever undertaken. A complete sample of 11~685 flat-spectrum radio sources ($S_{\\nu} \\propto \\nu^{\\alpha}$ where $\\alpha \\geq -0.5$ between 1.4 and 4.85~GHz) was observed with the Very Large Array (VLA) in A configuration at 8.46~GHz between 1994 and 1999. The 0.2-arcsec angular resolution provided by the VLA allowed the straight forward identification of galaxy-scale gravitational lenses, which typically have image separations of $\\sim 1$~arcsec. CLASS sources were identified as promising gravitational lens candidates if the VLA 8.46~GHz image contained multiple compact components with Gaussian FWHM $\\leq 170$~mas, the image separations were between $0.3 \\leq \\Delta\\theta \\leq 15$~arcsec, the total integrated flux-density was $S_{8.46} \\geq 20$~mJy and the component flux-density ratio was $\\leq 10:1$. The detection of compact structure and extended emission in the multiple components at higher resolutions with MERLIN (Multi-Element Radio Linked Interferometer Network) and the VLBA (Very Long Baseline Array) confirmed the lensing hypothesis for 21 gravitational lenses during the first three phases of CLASS (see \\citealt{browne02} and references therein for a discussion of all the gravitational lenses discovered by CLASS).  In this paper, we present observations of the gravitational lens candidate CLASS B2108+213, which was discovered during the fourth and final phase of CLASS. In Section \\ref{radio} we describe the VLA, MERLIN and VLBA radio observations which led to the discovery of the gravitational lens candidate CLASS B2108+213. Follow-up optical and infrared imaging of the system and the surrounding field, taken with the {\\it Hubble Space Telescope}, is presented in Section \\ref{photometry}. Using the available observational data, the lensing hypothesis, along with an analysis of the lensing potential, is discussed in Section \\ref{disc}. Finally, in Section \\ref{conc} we summarise our findings and outline the future work we intend to carry out on this intriguing new gravitational lens from CLASS. For all calculations we adopt an $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$ flat-universe, with a Hubble parameter of $H_{0} = 70$~km~s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{conc}  We have presented radio and optical observations of the new gravitational lens CLASS B2108+213. This system comprises two lensed images, which are separated by 4.56~arcsec, thus making CLASS B2108+213 the widest separation gravitational lens discovered by CLASS. Optical imaging found two lensing galaxies within the Einstein radius of the system. Intriguingly, we also detected radio emission from a third component which was coincident with the optical emission from the primary lensing galaxy. The extension of the third radio component and its different radio spectrum from the two lensed images suggests that it is unlikely to be a third lensed image. However, it is too early to completely discount the possibility that scatter broadening effects might have been responsible for modifying the properties of a third image. The positions and flux-densities of the lensed images were successfully reproduced by a model containing two singular isothermal spheres and an external shear. However, further observational constraints will need to be obtained for a detailed analysis of the CLASS B2108+213 lensing potential to be carried out.  The wide image separation of CLASS B2108+213 implies a group/cluster assisted lensing potential. Therefore, CLASS B2108+213 provides an excellent opportunity to study the structure and mass distribution of galaxy groups beyond the local universe ($z > 0.2$). We have begun a spectroscopic programme to determine the redshifts of the lensing galaxies and lensed source in CLASS B2108+213, and to identify which nearby galaxies are part of the lensing group. Thus far we have found four nearby galaxies at similar redshifts to G1 ($z_{G1} = 0.365$; McKean et al., in preparation). Furthermore, our spectrum of G1 has sufficient signal-to-noise for the stellar velocity dispersion to be measured. These observations will provide additional constraints for a future mass model.  The other interesting feature of CLASS B2108+213 is the possible detection of a third lensed image. Although the observations presented here show that this is almost certainly not the case, this conclusion could be confirmed by measuring the flux-density, polarization and size of the radio emission from each component as a function of frequency. If scattering of a third lensed image is occurring then we would expect the high frequency spectrum of each radio component to be identical.  \\subsection*"
    },
    "0410/astro-ph0410248_arXiv.txt": {
        "abstract": "The accretion induced neutron star magnetic field evolution is studied through considering the accretion flow to drag the field lines aside and dilute the  polar field strength, and as a result the equatorial field strength increases, which  is buried inside the crust on account of the accretion induced global compression of star crust.  The  main conclusions of model are as follows: (i) the polar  field decays with increasing the accreted mass; (ii) The bottom magnetic field strength of about $10^8$ G  can occur when neutron star magnetosphere radius approaches  the star radius,  and it  depends on  the  accretion rate as $\\mdot^{1/2}$; (iii) The neutron star magnetosphere radius decreases with accretion until it reaches the star radius, and  its evolution is little influenced by the initial field and the accretion rate after  accreting $\\sim 0.01 \\ms$, which implies that the magnetosphere radii  of neutron stars in LMXBs would be homogeneous  if they accreted the comparable masses. As an extension, the physical  effects of the possible strong magnetic zone in the X-ray neutron stars and recycled pulsars are discussed. Moreover,  the strong magnetic fields  in the binary pulsars  PSR 1831-00 and PSR 1718-19 after accreting  about half solar mass in the binary accretion phase, $8.7\\times10^{10}$ G and $1.28\\times10^{12}$ G, respectively, can be explained  through  considering the incomplete frozen flow in the polar zone. As a model's expectation, the existence of the low magnetic field ($\\sim 3\\times 10^{7}$ G) neutron stars or millisecond pulsars is suggested. ",
        "introduction": "Recently, the newly discovered double pulsars J0737-3039A and  J0737-3039B  in  the binary system (Burgay et al 2004; Lyne et al 2004; van den Heuvel 2004; Lorimer 2004) have  shown that the millisecond pulsar (MSP) possesses the low field  and short   spin period ($B=7\\times10^{9}$ G, P=22.7 milliseconds) and the normal pulsar possesses high field and long spin period ($B=6\\times10^{12}$ G, P=2.77 seconds), which convinces our previous  point of view that MSP is recycled in the binary system where the accreted  matter weakened its magnetic field and accelerated its  spin (see, e.g., van den Heuvel  2004). Moreover, with the launch of  RXTE satellite,  the accretion-powered X-ray pulsar SAX J 1808.4-3658 (P=2.49 milliseconds), as well as  other five examples (see, e.g., Chakrabarty 2004; van der Klis 2004; Wijnands \\& van der Klis 1998), and 11 type-I X-ray burst frequency,  confirmed as stellar  spin frequency, sources have been discovered (see, e.g., Muno 2004; Chakrabarty 2004; van der Klis 2004, 2000), which exhibits the strong evidence that neutron star (NS) in low mass X-ray binary (LMXB) is the progenitor of MSP with spin frequency $\\sim $ 400 Hz and magnetic field strength $\\sim 10^{8-9}$ G because of the accretion (see, e.g., Wijnands \\& van der Klis  1998; Chakrabarty 2004).   The origin, structure and evolution of NS  magnetic field is still an open problem (see, e.g., Bhattacharya \\& van den Heuvel 1991; Phinney \\& Kulkarni 1994; Bhattacharya \\& Srinivasen 1995; Colpi et al 2001; Melatos \\& Phinney 2001; Cumming 2004; Blandford et al  1983;  Blondin \\& Freese 1986). In the early stage of discovery of NS, it is believed that NS magnetic field decays  due to Ohmic dissipation  in the crust. But later calculations of Ohmic dissipation suggest that isolated NS magnetic field  may not decay significantly (Sang \\& Chanmugam  1987), if the field occupies the entire crust, and remains large for  more than  Hubble time. Although many instructive mechanisms on the isolated NS magnetic decay  have been proposed, such as,  the crustal plate tectonics model by Ruderman (1991) and Chen \\& Ruderman 1993, the Hall-drift dominated field evolution (see, e.g., Rheinhardt et al 2004; Jones 2004; Hollerbach \\& Ruediger 2002;  Geppert \\& Rheinhardt 2002; Rheinhardt \\& Geppert 2002; Naito \\& Kojima 1994), the spin-evolution  induced magnetic field decay (see, e.g., Konar \\& Bhattacharya 1997, 1999; Ding et al 1993; Jahan Miri \\& Bhattacharya 1994; Ruderman et al 1998), etc., there has not yet been a commonly accepted idea on such issue.    The accretion induced NS magnetic field decay  has not been paid much attention  until 1980's when Taam and van den Heuvel (1986) presented the NS magnetic field evolution associated with the accretion  in the binary system, based on which Shibazaki et al (1989) concluded the simple empirical  formula of the field deduction versus the accreted mass. However, it is indicated that the significant decay of magnetic field is achieved only if the neutron star experiences the interacting binary (see, e.g., Verbunt \\& van den Heuvel 1995; van den Heuvel 1995 and references therein).  Furthermore,  van den Heuvel and Bitzaraki (1995), from the statistical analysis of 24 binary radio pulsars with nearly circular orbits and low mass companions, discovered a clear correlation between spin period and orbital period, as well as the magnetic field and orbital period. These relations strongly suggest that the increased amount of accreted mass  leads to the decay of NS  polar magnetic field, and the `bottom' field strength  $\\sim 10^{8}$ G is also implied, which means that NS magnetic field evolution will stop at some minimum  value, $\\sim10^{8}$ G (see also, e.g., Burderi et al 1996; Burderi \\& D'Amico 1997).   In order to understand the accretion induced field decay, some suggestions and models have been proposed. Based on the accretion buried or screened NS magnetic field guess by  Bisnovati-Kogan and Komberg (1974), Zhang et al (1994, 2000) proposed the ferromagnetic crust screen model   to interpret the simple inverse correlation between the field deduction  and the accreted mass, as  declaimed by Shibazaki et al (1989). Moreover, the accretion induced flow and thermal effects to account for the speed-up Ohmic dissipation of NS crust currents are also proposed  and studied by several authors (see, e.g., Romani 1990;  Geppert \\& Urpin 1994; Urpin \\& Geppert 1995). Later on, the diamagnetic screening of NS magnetic field by freshly accreted material is suggested (see e.g. Lovelace et al. 2005), which would be effective above a critical accretion rate $\\mdot$ of a few percent of the Eddington rate (see, e.g., Cumming et al 2001). More recently, it is calculated that the polar cap widening or the equator-ward hydro-magnetic spreading   by the accreted material flowing accounts for the NS magnetic field decay (see, e.g., Konar \\& Choudhury 2004, references therein; Payne 2005; Payne \\&  Melatos 2004;  Melatos \\& Phinney 2001).  Furthermore, the magnetic burial mechanism has been studied in detail analytically and numerically  by Payne \\&  Melatos (2004) and Payne (2005) under the assumptions of  no instabilities or Ohmic diffusion, where the structure of highly distorted magnetic field and hence the magnetic  moment as a function of accreted mass is computed. They find  that   the deduction of the magnetic moment is inversely related to how much mass accreted and about $\\sim 10^{-5}$\\ms is required  to reduce significantly the magnetic dipole moment.    In this paper, we follow   the idea,  firstly proposed  by van den Heuvel and Bitzaraki (1995),  and also by Burderi and D'Amico (1997),  that    the accretion matter is channeled by the strong magnetic field initially into the two polar patches, corresponding to NS magnetosphere radius of several hundreds of stellar radius, pushes the field lines aside and thus dilutes the polar field strength due to the flow of  accreted  materials from the polar to the equator. The bottom field should be reached when the accretion becomes all over star, which corresponds to that the NS magnetosphere radius matches the star radius, and in turn it  infers the bottom  magnetic field to be $\\sim 10^{8}$ G.  The paper is organized  as follows: in the next section the model is described and the magnetic field evolution  equation  is derived. Section 3  outlines the applications of the model, including how the NS magnetic field and magnetosphere evolve with the accreted mass, the initial magnetic field and accretion rate, etc.. The discussions  are drawn in Section 4, where we summarize the conclusions of model and the  observational consequences. ",
        "conclusions": "In summary, the main conclusions of model are listed in the following: (1) The NS polar magnetic field decays in the binary accretion phase, and its evolution experiences the following processes. Roughly speaking,  $B\\propto \\dm^{-1}$ is the initial field dependent if $\\dm \\sim 10^{-4}\\ms$, and then    $B\\propto \\dm^{-1.75}\\bff$ is independent of the initial field after  $\\dm \\sim 10^{-2}\\ms$, until the final stage the field B  remaining  stable at the bottom value of about $B\\sim 10^{8}$ G after  $\\dm\\sim 0.2 \\ms$. (2) The bottom field is the accretion rate (or luminosity) dependent as $B\\propto \\mdot^{1/2}$, which concludes that some  Z (Atoll) sources with the Eddington  luminosity (1\\% Eddinngton luminosity)  may possess the magnetic fields of about $\\sim 3\\times 10^{8}$ ($\\sim 3\\times 10^{7}$) G. Therefore,  the existence of the low magnetic field ($\\sim 3\\times 10^{7}$ G) neutron stars or millisecond pulsars is suggested, which needs the confirmation from the future observations. (3) The appearance of the bottom field is  nothing to do with the initial field strength,  as shown in Fig.2 and Fig.3. (4) On the  NS magnetosphere evolution, $\\rrm \\propto \\dm^{-1}$ is the initial field dependent if $\\dm \\sim 10^{-4}\\ms$ , and then    $\\rrm \\propto \\dm^{-1}$ is independent of the initial field and the luminosity after $\\dm \\sim 10^{-2}\\ms$, until the final stage when the magnetosphere radius reaches the star radius. (5) The equator magnetic field increases with the accretion and is stronger than the initial  field strength, which is buried inside the crust. Moreover, as known, the occurrence of thermonuclear X-ray bursts in LMXBs requires the NSs to have the fields $B<10^{10}$ G (see Joss \\& Li 1980), so we may expect the NS in LMXB to accrete $\\dm>0.01\\ms$ and to enter into the homogeneous state, while the field B {\\it forgets} the initial field $\\bo$ and only {\\it remembers} the bottom field $\\bff$. Also, the magnetic field strength distribution of the isolated non-recycled NSs is from $\\sim 10^{10}$ G to $\\sim 10^{14}$ G (see, e.g., Vranesevic et al 2004; the highest magnetic field can be extended to  $10^{15}$ G if the magnetars are included, see Thompson \\& Duncan  1995), but most MSPs locate in a narrow domain of B field distribution, from $10^{8}$ G to $10^{9}$ G, which seems to support the fact that the NS final magnetic field  is not  scaled by the initial field $\\bo$ but scaled by the bottom  field $\\bff$. On the characteristic  accreted mass $\\mb\\sim 10^{-5}$ -- $10^{-3}$ \\ms for the significant deduction of the magnetic field, our result is very close to that obtained by Payne \\& Melatos (2004) and Payne (2005), however the physical meaning of latter is at where the hydrostatic pressure at the base of accretion column overcomes the magnetic tension and the matter spreads over the stellar surface to drag the polar  field lines toward the equator. In our model, $\\mb$ represents the fraction of the crust mass scaled by the ratio of the stellar radius to the initial magnetosphere radius, then the relation  of both characteristic masses is still unclear.  Finally, although our model has given the observational consistent results, it is a very simple and crude theoretical framework to account  for the accretion induced polar field decay. We think that the   basic physical  idea is clear and has been proposed by a couple of people (see, e.g., van den Heuvel and Bitzaraki 1995; Burderi et al 1996; Burderi \\& D'Amico 1997) that the transformation from the systematic channeled accretion to the spherical accretion yields the power to decrease the NS  polar field strength. However, many physical   details have  been neglected in the model, such as, the detail of accretion flow and its reaction with the field lines, the plasma instability in the NS melted surface, the details of how the  magnetic field is buried in the equatorial zone, the Ohmic dissipation, as well as the Hall-drift effects,  etc., therefore the considerations  of these will construct our future research explorations."
    },
    "0410/astro-ph0410004_arXiv.txt": {
        "abstract": "In the near future the Planck satellite~\\cite{planck} will gather impressive information about the anisotropies of the cosmic microwave background and about the galaxy clusters which perturb that signal. We will here review the ability of Planck to extract information about these galaxy clusters through the Sunyaev-Zeldovich imprint. We will conclude that Planck will provide a catalogue of galaxy clusters, which will be very useful for future targeted observations. We will explain why Planck will not be very good in extracting detailed information about individual clusters, except for the dominating Compton parameter, $y$, which will be measured to a few percent for individual clusters. In this last point I am being rather conservative, but that will leave space for pleasant surprises when the Planck data will be analysed. ",
        "introduction": "The Planck satellite will be a real piece of art and an amazing piece of hardware. It will provide us with information about the cosmic background radiation (CMB) beyond the wildest imagination of most people, surely beyond mine.  The sensitive detector will be sent far into space, over a million kilometers from earth, where it will be better protected from the Sun and Earth radiation. The satellite will be spinning and carefully mapping the whole sky. Thus it will give us accurate data about the full CMB sky, and about all the obstructions and contaminations between us and the last scattering surface of the CMB.  Many research years have been, and will still be, devoted to the analysis and interpretation (and possibly even understanding) of the CMB signal, and there is no question that the main scientific goal of Planck will be to carefully study the CMB. In that sense all the foregrounds, such as galaxy clusters and dust, are just contaminations which will have to be carefully removed. I am going to be discussing if we can use this removed ``contamination'' for anything useful.  It is rather amusing that such beautiful structures as galaxies and galaxy clusters are just considered to be contamination by many CMB people.  The analysis of the Planck CMB data will determine many of the cosmological parameters very accurately. This includes dark matter and dark energy parameters as described by R. Bean and A. Melchiorri in these proceedings. Numerous papers have reviewed the expected precision with which Planck will determine such parameters, including refs.~\\cite{balbi,melch,bowen}.  One important strength of Planck is that it will have 9 observing frequencies, and as we will see below this will allow a separation of the CMB signal from the contaminations. The actual construction and optimization of performance of the Planck detector is a very complicated process, and I refer to the lectures by P. de Bernardis and J. M. Lamarre in these proceedings for all the details. There are also recent nice reviews on the two instruments on-board Planck~\\cite{lfi,hfi}, see also~\\cite{villa,Yurchenko,naselsky}.  Planck will find of the order 10 thousand massive galaxy clusters out to a redshift of $z \\sim 1.5$.  This cluster catalogue will be deeper and larger (all sky survey) than any other existing catalogue. This cluster catalogue will be very useful on its own and for future follow-up observations. It is the goal of this lecture to discuss at an elementary level how this will be achieved, and secondly we will discuss more carefully what Planck will {\\em not} be able to do. Several groups have addressed similar issues already, see e.g. refs.~\\cite{aghanim97,white03,geisbuesch}. ",
        "conclusions": "The Planck satellite is a true piece of art, which will provide impressive information about the CMB sky. The SZ signal can easily be removed, and will not be a problem for CMB parameter extraction.  Planck will also give us a cluster catalogue which is larger, deeper and has more sky-coverage than any other existing catalogue. The Compton parameter for the individual clusters in this catalogue will be measured to $1\\%$ accuracy for the brightest clusters, and to about $10\\%$ for the dimmer clusters. Such a catalogue is very useful for future targeted SZ studies.  I have argued why I personally believe that the systematic errors for the individual clusters are still so hard to assess, that cosmological parameter extraction with this cluster sample will be very difficult. My view is rather conservative, so this leaves space for pleasant surprised when we will analyse the coming Planck data carefully."
    },
    "0410/astro-ph0410718_arXiv.txt": {
        "abstract": "During $\\eta$ Car's spectroscopic event in mid-2003, the stellar wind's bright H$\\alpha$ and H$\\beta$ emission lines temporarily had a distinctive shape unlike that reported on any previous occasion, and particularly unlike the 1997--98 event.  Evidently the structure of the wind changed between 1997 and 2003.  Combining this with other evidence, we suspect that the star may now be passing through a rapid stage in its recovery from the Great Eruption seen 160 years ago.  In any case the data indicate that successive spectroscopic events differ, and the hydrogen line profiles are quantitative clues to the abnormal structure of the wind during a spectroscopic event. ",
        "introduction": "Systematic spectral variations are clues to the nature of $\\eta$ Car --- our most observable very massive star dangerously close to the Eddington limit, and the site of a prototype ``supernova impostor'' eruption seen about 160 years ago.  On several occasions between 1945 and 1990 the general excitation or ionization decreased for a few months, in the stellar wind and also in slow-moving ejecta located within 1000 AU of the star (see \\citet{zanella84,dh97,ad99, kd99a}; and other refs.\\ cited therein). Prompted by a fresh example observed in 1992, \\citet{ad96} noticed that these events recur in a 5.5-year cycle which had also appeared in photometry \\citep{pw94}.  The next two spectroscopic events occurred near 1998.0 and 2003.5, close to the predicted times.   As noted in Section 5 below, at least three rival scenarios may explain an event of this type.    The star may eject a shell, possibly triggered by a hypothetical companion star in a 5.5-year orbit;  the excitation of our side of the wind may temporarily decrease when a hot secondary star moves to the far side of its orbit;  UV radiation from a hot companion may become engulfed in the primary wind near periastron; or some mixture of these effects may occur.  Unfortunately no quantitative model has been developed, and one can imagine other possibilities. A strange detail, highly relevant to the discussion in this paper, is that the 5.5-year cycle probably did not exist before the 1940's -- at least not in its present form; see Section 5 below and references cited there.   Suitable data on these events were scarce until recently. Ground-based spectroscopy  of the star is seriously contaminated by diffuse ejecta located at radii 0.2{\\arcsec} -- 2{\\arcsec}; and, so far as we know, no long-term, spatially consistent ground-based observations of a broad wavelength range were attempted. The Hubble Space Telescope (HST) can attain the required high spatial resolution, but was used for spectroscopy of $\\eta$ Car on only a few occasions from 1991 to 2001 -- fortunately including part of the 1997--98 event (see below).  The fund of available data greatly improved during 2002--03, when we obtained numerous HST observations using the Space Telescope Imaging Spectrograph (STIS) as part of a Hubble Treasury Project to observe the 2003.5 event \\citep{kd02b}.   Here we report a particularly striking result:  H$\\alpha$ and H$\\beta$, the brightest observable emission lines in the stellar wind, briefly developed unusual velocity profiles which were {\\it quite different from their appearance during the previous event in 1997--98.\\/} This development, combined with other circumstances noted in Sections 4 and 5 below, leads us to suspect that $\\eta$ Car may now be passing through an unusually rapid phase in its recovery from the 19th-century Great Eruption.  Independent of this conjecture, our data show that the 5.5-year cycle is not a straightforward clocklike affair; while the emission-line behavior reported here is essential for modeling the spectroscopic event.  Parallel ground-based observations of the polar-latitude spectrum reflected by the ``Homunculus'' ejecta-nebula will be described in a companion paper by \\citet{ws05}. ",
        "conclusions": ""
    },
    "0410/astro-ph0410232_arXiv.txt": {
        "abstract": "We present the first coordinated VLBA/VLTI measurements of the stellar diameter and circumstellar atmosphere of a Mira variable star. Observations of the $v=1, J=1-0$ (43.1 GHz) and $v=2, J=1-0$ (42.8 GHz) SiO maser emission toward the Mira variable S Ori were conducted using the VLBA.  Coordinated near-infrared $K$-band measurements of the stellar diameter were performed using VLTI/VINCI closely spaced in time to the VLBA observations.  Analysis of the SiO maser data recorded at a visual variability phase 0.73 show the average distance of the masers from the center of the distribution to be 9.4~mas for the $v=1$ masers and 8.8~mas for the $v=2$ masers. The velocity structure of the SiO masers appears to be random with no significant indication of global expansion/infall or rotation. The determined near-infrared, $K$-band, uniform disk (UD) diameters decreased from $\\sim$\\,10.5\\,mas at phase 0.80 to $\\sim$10.2\\,mas at phase 0.95.  For the epoch of our VLBA measurements, an extrapolated UD diameter of $\\Theta_\\mathrm{UD}^K=10.8 \\pm 0.3$\\,mas was obtained, corresponding to a linear radius of $R_\\mathrm{UD}^K = 2.3 \\pm 0.5$~AU or $R_\\mathrm{UD}^K =490 \\pm 115~R_\\odot$. Our coordinated VLBA/VLTI measurements show that the masers lie relatively close to the stellar photosphere at a distance of $\\sim$\\,2 photospheric radii, consistent with model estimates.  This result is virtually free of the usual uncertainty inherent in comparing observations of variable stars widely separated in time and stellar phase. ",
        "introduction": "The evolution of cool luminous stars, including Mira variables, is accompanied by significant mass loss to the circumstellar envelope (CSE) with mass-loss rates of up to $10^{-4}$\\,M$_\\odot$/year \\citep[e.g.][]{JK:90}. There are currently a number of tools that can be used to study the stellar surface and the CSE at various wavelengths. Optical/near-infrared long-baseline interferometry has provided information regarding the stellar diameter, effective temperature, center-to-limb intensity variation, and the dependence of these parameters on wavelength and variability phase for a number of Mira variables \\citep[e.g.,][]{HST:95,BDBL:96,PFR:99,YOUNGETAL:00,TCA:02,TCB:02,WED:04}. The structure and dynamics of the CSE of Mira variables and supergiants have been investigated by mapping the SiO maser emission at typical distances of 2--4 stellar radii toward these stars using very long baseline interferometry (VLBI) at radio wavelengths \\citep[e.g.,][]{BDK:97,KD:97,BM:00,H:01,DK:03}. Dust shells outward of the SiO maser emission zone have been studied using mid-infrared interferometry \\citep[e.g.,][]{D:94,G:95,T:03}. Near- and mid-infrared interferometry have also been used to study the warm extended atmosphere of Mira variables containing gaseous H$_2$O, SiO and CO molecules \\citep[e.g.,][]{PFR:99,MPC:02,OBD:04}.  Theoretical models describe and predict photospheric center-to-limb intensity variations including the effect of close molecular layers \\citep[e.g.,][]{Bessel:96, HSW:98, Tej:03, ISW:04}, the location and dynamics of SiO maser emission zones \\citep{Humphreys:96,Humphreys:02}, as well as gas and dust components of the stellar winds \\citep[e.g.,][]{Winters:00,Hoefner:03}.  Multi-wavelength studies of the stellar surface and the CSE, utilizing a combination of the above techniques, are well suited to study the structure of the CSE, the mass-loss process, and the evolution of these stars. It is common for infrared observations of the stellar photosphere or the circumstellar dust to be compared to radio observations of circumstellar masers from the literature and vice versa.  For example, \\citet{D:94} compared the inner radii and extents of dust shells as determined from Infrared Spatial Interferometer (ISI) observations with estimates of the photospheric radii and measurements of the extents of the SiO, H$_2$O, and OH maser shells. Similarly, \\cite{MONNIER:04} compared high-resolution data on the stellar diameters and dust shells of two stars (VX Sgr and NML Cyg) from Keck aperture masking and Infrared Optical Telescope Array (IOTA) interferometry with previously published maps of the SiO maser emission toward these stars.  Such comparisons, however, are somewhat limited due to the inherent variability in the CSE and the star itself. The SiO masers for example have been shown to have proper motions and intensity variations on time scales of a few weeks \\citep{BDK:97,DK:03}.  Similarly, interferometric measurements of the stellar parameters vary with stellar variability phase as discussed above.   This inherent variability of Mira stars and their surroundings necessitates a more coordinated approach to multi-wavelength studies of these objects.  The first such coordinated multi-wavelength study of a late-type variable was performed by \\citet{G:95} for the star VX~Sgr.  Coordinated measurements included both mid-infrared observations to determine the extent of the circumstellar dust shell and VLBI mapping of the SiO maser zone. More recently, \\citet{C:04} compared Very Long Baseline Array (VLBA) images of the SiO maser emission of several Mira stars with near-infrared diameters obtained at the Infrared Optical Telescope Array (IOTA) interferometer \\citep{MPC:02}, and with literature values of the mid-infrared dust shell extensions.  Here, we present the first results from a series of coordinated VLTI/VLBA observations of late-type variable stars, in particular, for the Mira variable S~Ori. Our results include the first radio interferometric images of the SiO masers (VLBA) and the most accurate photospheric diameter measurements using near-infrared $K$-band interferometry (VLTI/VINCI). ",
        "conclusions": "\\subsection{The photospheric diameter \\label{sec:photosphere}}  Figure~\\ref{diamdate} and Tab.~\\ref{tab:vinciresults} show our near-infrared $K$-band uniform disk (UD) diameters of S~Ori determined from our four VLTI/VINCI epochs as a function of observation date and visual variability phase, as discussed in Sect.~\\ref{sec:obsvlti}. The $K$-band UD diameter decreases almost linearly from $\\sim$\\,10.5\\,mas at visual variability phase 0.80 to $\\sim$\\,10.2\\,mas at phase 0.95, i.e. by $\\sim$\\,3\\%. It is not clear whether the local diameter minimum of $\\sim$\\,10.0\\,mas at phase 0.92 in Tab.~\\ref{tab:vinciresults} and Fig.~\\ref{diamdate} is caused by an additional systematic calibration error as discussed in Sect.~\\ref{sec:obsvlti}, or by a real minimum of the stellar size. The minimum of a Mira $K$-band UD diameter is predicted to lie at visual variability phase $\\sim$\\,0.9, according to model predictions by \\cite{ISW:04}. Our coordinated VLBA measurement at visual variability phase 0.73 occurred about one month before the first VLTI epoch at phase 0.8, and is hence close to, but not exactly contemporaneous to the VLTI epochs. Both, observations of the diameter change of the Mira star S\\,Lac by \\cite{TCB:02} as well as Mira star model predictions by \\cite{ISW:04} show an almost linear change of the $K$-band UD diameter between variability phases 0.7 and 0.9, i.e. in the pre-minimum part of the diameter curve and pre-maximum part of the visual light curve. Thus, we can linearly extrapolate our $K$-band UD diameter data at phases 0.8--0.95 to the VLBA phase 0.73. For phase 0.73, we determine an extrapolated $K$-band UD diameter of $\\Theta_\\mathrm{UD}^K (\\mathrm{VLBA\\,epoch, phase=0.73}) = 10.8 \\pm 0.3$\\,mas. This value and its error correspond to the mean and difference of the two linear extrapolations shown in Fig.~\\ref{diamdate}, which used all four points and only the two closest points to the VLBA epoch respectively. This UD diameter corresponds to a linear radius $R_\\mathrm{UD}^K = 2.3 \\pm 0.5$~AU or $R_\\mathrm{UD}^K =490 \\pm 115~R_\\odot$ with our assumed distance to S~Ori.  A UD intensity profile is often not an ideal representation of the true near-infrared center-to-limb intensity variation (CLV) of cool giants in general and of Mira stars in particular. The CLV of cool giants has been studied by, e.g. \\cite{QMB:96,BURNS:97,SCHOLZ:98,HSW:98,WHJ:01,WAK:04,ISW:04, WED:04}.  UD diameters of non-variable giants obtained from visibility measurements in the first lobe are usually transformed into physically more meaningful Rosseland angular diameters using correction factors determined from atmosphere models in the literature \\citep[e.g.][]{C:03}. For the case of Mira variable stars, few atmosphere models are available.  We used the hydrodynamic atmosphere models from \\cite{HSW:98,Tej:03,ISW:04},  to estimate the relationship between the $K$-band UD diameter and the continuum diameter. While the broad-band UD diameter is affected by various molecular bands, the continuum diameter is a better estimate of the real photospheric stellar size. These models have been constructed for the prototype Mira stars $o$\\,Cet and R\\,Leo and have been compared to observations of several Mira stars by, for instance, \\cite{HBSW:01,HOFMANN:02,WED:04}. The general model results are not expected to be dramatically different for other Mira stars such as S Ori (Scholz, M. 2004, private communication). The model predicted difference between the continuum and $K$-band UD diameters is relatively low in the pre-maximum region of the visual variability curve as in the case of our observations. At this phase of 0.73, the continuum diameter may be smaller than the $K$-band UD diameter by about 15\\% \\citep{ISW:04} . With this assumption, the continuum photospheric diameter for the epoch of our VLBA observation would be $\\Theta_\\mathrm{Phot}(\\mathrm{VLBA\\ epoch,\\ phase=0.73})\\approx 9.2$\\,mas. This angular radius corresponds to a  photospheric radius of   $R_\\mathrm{Phot} \\approx 420$\\,R$_\\odot$ or $R_\\mathrm{Phot}\\approx 1.9$\\,AU with our assumed distance to S~Ori. With a bolometric flux of $m_\\mathrm{bol}=3.1 \\pm 0.3$ \\citep{WMF:00}, our value for $\\Theta_\\mathrm{Phot}$ corresponds to an effective temperature $T_\\mathrm{eff}\\approx 2670$\\,K. \\subsection{The $v=1$ and $v=2$ SiO Maser Emission \\label{MASER_RESULTS}}  Our total intensity images of the 43.1 GHz and 42.8 GHz SiO maser emission (Figure~\\ref{SORI_COMB_IMAX}) show a typical clumpy distribution of the maser spots within a ring-like structure.  For the $v=2$, 42.8-GHz transition the ring is relatively sparse with nearly all of the masers concentrated to the NW side of the shell, and a few features to the SE.  The $v=1$, 43.1-GHz masers, however, form a more typical ring-like structure often seen for other stars with SiO masers. Like the $v=2$ masers, the $v=1$ SiO also has a higher concentration of features on the NW side of the shell.  The $v=1$ maser ring appears to be symmetrical, forming a nearly circular shell.  The distribution of the $v=2$ SiO masers is too sparse to allow us to comment on any symmetry for this transition.  In order to identify and parameterize maser components, two-dimensional Gaussian functions were fit to the emission in each spectral (velocity) channel using the AIPS task SAD.  Image quality was assessed using the off-source RMS noise and the deepest negative pixel in the image. A cutoff flux density was conservatively set to the greater of 8$\\sigma_{\\rm RMS}$ or the absolute value of the deepest negative pixel in the plane.  Features with flux densities greater than this cutoff were fit with Gaussians to determine component parameters.  Errors in the positions in right ascension and declination of identified features were computed using the fitted source size divided by twice the signal-to-noise ratio (SNR) in the image and ranged from 1~$\\mu$as for features with high SNR, to 50~$\\mu$as for features with lower SNR.  Since the $\\sim$0.2~km~s$^{-1}$ channel spacing is sufficient to resolve the masers spectrally, features typically appear in multiple adjacent spectral channels.  Positions in right ascension and declination and center velocities for the masers were determined using a flux-density-squared weighted average for features identified in two or more adjacent channels with a spatial coincidence of 0.2~mas ($\\sim$2/3 of the geometric mean of the synthesized beam).  The flux assigned to the maser averages was the maximum single-channel flux density. The maser components identified using this procedure are represented by the circles in Figures \\ref{SORI_43.1_COMPS} and \\ref{SORI_42.8_COMPS}.  In the figures, point sizes are proportional to the logarithm of the fitted flux density which ranged from 0.2 to 4.3~Jy for Figure \\ref{SORI_43.1_COMPS} and 0.1 to 5.3~Jy for Figure \\ref{SORI_42.8_COMPS}.  The line-of-sight (LOS) velocity information for the masers is also represented in Figures \\ref{SORI_43.1_COMPS} and \\ref{SORI_42.8_COMPS}.  The top panel of each figure shows the spectrum of SiO maser emission ranging from 10 to 25~km~s$^{-1}$ color-coded by LOS velocity in increments of 2~km~s$^{-1}$.  The bottom panels of Figures \\ref{SORI_43.1_COMPS} and \\ref{SORI_42.8_COMPS} show the spatial distribution of the SiO masers plotted with the same velocity color-coding as in the top panel, and with the color of the maser representing its corresponding velocity range in the spectrum.  Comparing the maser component maps with the total intensity contour maps, we see that the identified features accurately represent the emission summed over all velocity channels in the image cube.  From the component maps there does not appear to be any coherent velocity structure indicative of global expansion/infall or rotation.  We do note that a group of components on the western side of the shell shows a velocity gradient with velocity decreasing with increasing distance from the star.  Such velocity gradients in the SiO maser emission have been observed previously \\citep[e.g.][]{BM:00, H:01}.  To compare the size of the SiO maser distribution to the photospheric diameter of the star as measured by the near-infrared $K$-band interferometry we determined the average distance of the masers from the center of the distribution.  To accomplish this, we first determined the center of the distribution by performing a least-squares fit of a circle to the combined $v=1$ and $v=2$ maser component data.  This fit produced a common center from which we computed the mean maser angular distance $\\overline{r}_{\\rm SiO}$ and the standard deviation for each transition independently.  The mean angular distances from center for the observed SiO masers at $v=1$, 43.1-GHz and $v=2$, 42.8-GHz are $\\overline{r}_\\mathrm{SiO}=9.4$ and $\\overline{r}_\\mathrm{SiO}=8.8$~mas respectively.  These distances are indicated by a dashed circle and are listed in the lower panel of Figures \\ref{SORI_43.1_COMPS} and \\ref{SORI_42.8_COMPS}.  Standard deviations of the distances are $\\sigma_\\mathrm{SiO}=1.4$ and $\\sigma_\\mathrm{SiO}=1.7$~mas for the $v=1$ and $v=2$ transitions respectively.  The standard deviations provide an indication of the thickness of the shell and are likely dominated by the features on the western side of the shell with a wide range of distances from the center. The mean angular distances we derive are consistent with least-squares circle fits to each distribution independently using the previously mentioned common center.  The least-squares fits to each transition have the added assumption that the distribution is circular, thus we have chosen to report only the mean distance from center. The computed angular distances translate to linear shell distances from center of $4.0 \\pm 0.6$~AU and $3.7 \\pm 0.6$~AU at the assumed distance of $422 \\pm 37$~pc.  The standard deviations are likewise $0.6 \\pm 0.1$~AU and $0.7 \\pm 0.1$~AU for the $v=1$ and $v=2$ masers respectively.  In principal, a comparison of the relative spatial locations of the $v=1$ and $v=2$ masers should allow us to comment on possible pumping mechanisms (i.e. collisional, radiative, or combination) as in \\cite{Desmurs:00}.  However, because the $v=2$, 42.8-GHz masers are mostly confined to a small region on the NW side of the shell, it is difficult to unambiguously determine the relative shift between the two maser shells.  Although the mean distance from center we determine for the $v=2$ masers is slightly smaller that the $v=1$ mean distance, the two values are consistent to within the errors.   \\subsection{Comparison of SiO maser distribution and photospheric stellar diameter}  Figures \\ref{SORI_43.1_COMPS} and \\ref{SORI_42.8_COMPS} also provide a comparison of the distribution of the SiO maser spots with our obtained angular size of the photospheric disk at the same epoch. The average distance of the maser spots from the center of their distribution for the 43.1 GHz and 42.8 GHz transitions ($\\overline{r}_\\mathrm{SiO}=9.4$\\,mas and $\\overline{r}_\\mathrm{SiO}=8.8$\\,mas) appears at 1.7\\,R$_\\star$ and 1.6\\,R$_\\star$ respectively when compared to our estimate of the $K$-band UD diameter.  When compared to our estimate of the continuum diameter these values are slightly larger at 2.0\\,R$_\\star$ and 1.9\\,R$_\\star$ respectively. The widths of the distributions are approximately 0.3\\,R$_\\star$ for the masers in both transitions. As a reference the angular photospheric radius is listed in the lower panel of Figures \\ref{SORI_43.1_COMPS} and \\ref{SORI_42.8_COMPS} and is indicated by a colored circle at the center of the maser distribution. The position of the near-infrared stellar disk as indicated on the plots is assumed to coincide with the center of the SiO maser distribution as discussed in Sect.~\\ref{MASER_RESULTS}.  The true location of the star relative to the masers is still unknown. We note that the above estimates are based on angular sizes derived for the SiO maser emission and the near-infrared photosphere, and hence are independent of the distance to S Ori.  This result is consistent with theoretical estimates by \\cite{Humphreys:02} based on a stellar hydrodynamic pulsation model combined with an SiO maser model.  They obtain values between 1.7\\,R$_\\star$ and 2.1\\,R$_\\star$, depending on the variability phase.  Our result is consistent with a SiO maser ring radius of approximately 2\\,R$_\\star$ as determined by \\cite{C:04} for several Mira variables.  Recently, however, \\cite{MONNIER:04} updated the stellar diameter estimate of \\cite{G:95} for the supergiant VX~Sgr and found that the SiO masers lie at a greater distance from the photosphere, 3.9\\,R$_\\star$, rather than the 1.3\\,R$_\\star$ determined by \\cite{G:95}.  It is striking that the stellar diameter for the other supergiant observed by \\cite{MONNIER:04}, NML Cyg,  also indicates a SiO maser ring distance of $\\sim$4\\,R$_\\star$ when compared to the SiO maser shell diameter obtained by \\cite{BM:00}. It is unclear whether these larger distances are an artifact caused by the non-contemporaneous measurements or, whether they indicate an inherent difference between Miras and supergiants. In earlier studies, measured sizes for SiO maser shells have been compared to stellar sizes based on diameters found in the literature \\cite[e.g.][]{BDK:97,BM:00,H:01}. These diameters are often widely spaced in time and stellar phase from the measurements of the SiO maser shell. Sometimes the stellar diameters have not been measured at all and comparisons are made with even less precise diameter estimates based on the luminosity of the star \\cite[e.g.][]{DIAMOND:94,COLOMER:96,SANCHEZ:02}. These examples demonstrate that it is desirable to compare SiO maser ring diameters with stellar diameter estimates, and highlight the need for contemporaneous multi-wavelength observations."
    },
    "0410/gr-qc0410105_arXiv.txt": {
        "abstract": "\\doublespace We present a physically reasonable source for an static, axially--symmetric solution to the Einstein equations. Arguments are provided, supporting our belief that the exterior space--time produced by such source, describing a quadrupole correction to the Schwarzschild metric, is particularly  suitable (among known solutions of the Weyl family) for discussing the properties of quasi--spherical gravitational fields.  \\vspace*{0.21truein} Key words: Axially--symmetric solutions. ",
        "introduction": "As is well known, Weyl solutions \\cite{2} represent the family of all static and axially--symmetric exterior solutions to the Einstein equations. Since there are as many different Weyl solutions as there are different harmonic functions (see next section), then the  obvious  question arises: what is the exact vacuum solution to the Einstein equations corresponding to a given non--spherical, static axially symmetric source (an ellipsoid, say)?.  If the field is not particularly intense ($r\\gg 2M$) and the deviation from spherical symmetry is sligth, then there is not problem in representing the corresponding field (both inside and outside the source) as a suitable perturbation of the spherically symmetric exact solution. However, as the object becomes more and more compact, such perturbative scheme will eventually fail.  Indeed, as it is well known \\cite{1}, the only static and asymptotically--flat vacuum space--time possessing a regular horizon is the Schwarzschild solution. For all the others Weyl exterior solutions \\cite{2}, the physical components of the Riemann tensor exhibit singularities at $r=2M$. Therefore, it is intuitively clear that as $r$ approaches $2M$ and gravitational field becomes stronger, the properties of sources of Weyl space--time should start drastically to differ from the properties of spherical sources \\cite{Bel}. It is important to keep in mind that this sharp difference in the behaviour of both types of sources (for very high gravitational fields) is independent on how small, multipole moments (higher than monopole) of the Weyl source, are. This is so because, as $r$ approaches $2M$, any finite perturbation of the Schwarzschild space--time becomes fundamentally different from any Weyl solution, even when the latter is characterized by parameters whose values are arbitrarily close to those corresponding to the spherical symmetry. This point has been stressed before \\cite{4}, but usually it has been overlooked.  Furthermore in a recent work \\cite{Herrera} it was shown that for a non--spherical source (even in the case of slight deviations from spherical symmetry), the speed of entering the collapse regime decreases substantially, as compared with the exactly spherically symmetric case. In the same order of ideas it has been shown \\cite{Herreramass} that small departures from sphericity, produce significant decreasing (increasing) in the values of active gravitational mass of collapsing (expanding) spheres, with respect to its value in  equilibrium, enhancing thereby the stability of the system. Also, the sensitivity of the trajectories of test particles in the $\\gamma$ spacetime, to small changes of $\\gamma$, for orbits close to $2M$, has been brought out \\cite{HS}. It is important to stress that all these effects take place for strong gravitational fields, but for $r > 2M$.  These works \\cite{Herrera}, \\cite{Herreramass}, \\cite{HS} were done using as Weyl solution, the so--called gamma metric ($\\gamma $--metric) \\cite{zipo}. This metric, which is also known as Zipoy-Vorhees metric, belongs to the family of Weyl's solutions, and is continuously linked to the Schwarzschild space--time through one of its parameters. The motivation for this choice was that the exterior $\\gamma $--metric corresponds to a solution of the Laplace equation (in cylindrical coordinates) with the same singularity structure as the Schwarzschild solution (a line segment \\cite{zipo}). In this sense the $\\gamma $--metric appears as the ``natural'' generalization of Schwarzschild space-time to the axisymmetric case.  However, the existence of so many different (physically distinguishable \\cite{luisyo}) Weyl solutions gives rise to the question: which among Weyl solutions is better entitled to describe small deviations from spherical symmetry?.  Although it should be obvious that such a question does not have a unique answer (there is an infinite number of ways of being non--spherical, so to speak), we shall invoke a very simple criterion, emerging from Newtonian gravity, in order to choose our solution.   Indeed, in order to answer  the question above, let us recall \\cite{yo} that most known Weyl solutions,  present a drawback  when describing quasi--spherical space--times. It consists of the fact that its multipole structure is such that multipole moments, higher than quadrupole, are of the same order as the quadrupole moment. Instead, as it is intuitively clear, the relevance of such multipole moments should decrease as we move from lower to higher moments, the quadrupole moment being the most relevant for a small departure from sphericity. Thus for example in Newtonian gravity, multipole moments of an ellipsoid of rotation, with homogeneous density, and axes (a,a,b) read: \\[ \\left\\{ \\begin{array}{lll} D_{2n} &= &(-2)^n 3 M a^{2n}\\epsilon^n(1-\\epsilon/2)^n/[(2n+1)(2n+3)], \\,\\,\\, \\epsilon\\equiv (a-b)/a,\\nonumber\\\\ D_{2n+1} &= & 0\\nonumber \\end{array} \\right.,\\] because of the factor $\\epsilon^n$, this equation clearly exhibits the progressive decreasing of the relevance of multipole moments as $n$ increases.  Thus, in order to describe small departures from sphericity, by means of exact solution to the Einstein equations, we  would require an exact solution whose multipole structure shares the property mentioned above. Fortunately enough, such solution exists \\cite{yo}. Indeed, there is one (exact) solution  of the Weyl family, which may be interpreted as a quadrupole correction to the Schwarzschild space--time (see below).  It is for this exterior metric that we are going to construct a source. The motivation for this  is twofold: On the one hand, it is always interesting to propose bounded and physically reasonable sources of gravitational fields, which may serve as models of compact object. On the other hand, spherical symmetry is a common assumption in the study of compact self--gravitating objects (white dwarfs, neutron stars, black holes). Therefore it is pertinent to ask, how do small deviations from this assumption, related to any kind of perturbation  (e.g. fluctuations of the stellar matter, external perturbations, etc.), affect the properties of the system?. However, as mentioned before, for sufficiently strong fields, in order to answer to this question it is necessary to deal with non--spherically symmetric exact solution to the Einstein equations.  For constructing the source we shall follow a prescription given by Hern\\'an\\--dez \\cite{8} allowing to obtain interior solutions of Weyl space--time, from known spherically symmetric interior solutions. Our interior solution will be obtained from the interior Schwarzschild solution (homogeneous density).  The paper is organized as follows. In the next section we review Weyl solutions and the concept of relativistic multipoles. In Section 3 we describe the M--Q solution \\cite{yo} and give its properties. The Hern\\'andez method is applied in Section 4 to obtain an interior to the M--Q solution. Finally, results are discussed in the last Section. ",
        "conclusions": "We have seen that the  M-Q$^{(1)}$ solution satisfies the requested condition to be considered as a  quadrupole correction to  the spherical symmetry, namely: all relativistic moments higher than the quadrupole are of higher order in $q$. Accordingly it represents, among the known members of the Weyl family, a good candidate to describe small departures from spherical symmetry.  Next, we have found a source for that space--time. The interior solution obtained by an application of the Hern\\'andez algorithm, matchs smoothly to the  M-Q$^{(1)}$ metric on the boundary of the matter distribution, is regular  and satisfies all standard physical conditions. In the weak field approximation it represents a quadrupole correction to the incompressible fluid sphere.  As we have mentioned before, a spherical core was introduced in order to assure acceptable physical behaviour at the centre. It consists of a sphere of incompressible fluid with possitive energy density, larger than pressure and which matches smoothly to the outer part of  the source. For the models presented below the spherical core represents about the $0.1\\%$ of the total volume of the source.  Since we are interested in slight deviations from spherical symmetry we shall consider  models with very small values of $q$. Indeed the magnitude of the $q$'s in the models presented are many orders of magnitude smaller than the values corresponding to the earth and the sun, which in our units are approximately $10^{15}$ and $10^5$ respectively \\cite{luisyo}. However for a neutron star  of one solar mass, $10$ km radius and the same eccentricity as the sun, the order of magnitude of $q$ is $10^{-4}$ \\cite{luisyo}, as in one of the models below.  It is also worth noticing that the critical value of alfa ($\\approx 0.7$) is smaller than the corresponding value in the exactly spherically symmetric case (8/9)  We have ran a large number of models for a wide range of values of the parameters and, both, positive and negative $q$'s. Below we show figures corresponding to two models, one with small $\\alpha$ and the other with a value of $\\alpha$ close to its critical value. In both models we considered negative values of $q$ since we are primarily interested  in oblate objects.  Figures 2--5 exhibit the behaviour of physical variables for an small value of $\\alpha$ (weak field limit). Despite the small value of the quadrupole parameter, its contribution is clearly shown in figures 2--4. In figure 5 the numerical error is comparable with the quadrupole contribution, and accordingly the latter is somehow screened by the former.  Figures 6-8 display the physical variables for a large value of $\\alpha$. In this case, the quadrupole correction appears sharply in $T^0_{0}$ and $T^1_{2}$. However it is neglectable in $T^1_1$, $T^2_2$ and $T^3_3$.  Finally, figure 9 display the range of parameters $\\alpha$ and $q$ for which the solution is physically acceptable.  It is our hope that this source as well as  other interiors to some Weyl metrics presented elsewhere (\\cite{7} and \\cite{magli}) could be used as initial configurations to describe the departure from equilibrium of very compact objects endowed with an small but non--vanishing quadrupole structure."
    },
    "0410/astro-ph0410468_arXiv.txt": {
        "abstract": "{The detection of transits is an efficient technique to uncover faint companions around stars. The full characterisation of the companions (M-type stars, brown dwarfs or exoplanets) requires high-resolution spectroscopy to measure properly masses and radii. With the advent of massive variability surveys over wide fields, the large number of possible candidates makes such a full characterisation impractical for all of them. We explore here a fast technique to pre-select the most promising candidates using either near-IR photometry or low resolution spectroscopy. We develop a new method based on the well-calibrated surface brightness relation along with the correlation between mass and luminosity for main sequence stars, so that not only giant stars can be excluded but also accurate effective temperatures and radii can be measured. The main source of uncertainty arises from the unknown dispersion of extinction at a given distance. We apply this technique to our observations of a sample of 34 stars extracted from the 62 low-depth transits identified by OGLE during their survey of some 10$^5$ stars in the Carina fields of the Galactic disc. We infer that at least 78\\% of the companions of the stars which are well characterised in this sample are not exoplanets. Stars \\ogle{105}, \\ogle{109} and \\ogle{111} are the likeliest to host exoplanets and deserve high-resolution follow-up studies. \\ogle{111} was very recently confirmed as an exoplanet with M$_{planet}$ $\\cong$ 0.53 $\\pm$ 0.11 M$_{Jup}$ (Pont et al. 2004), confirming the efficiency of our method in pre-selecting reliable planetary transit candidates. ",
        "introduction": "The search for extra-solar planets (exoplanets) is becoming increasingly important, yielding new insights both in stellar and planetary physics (e.g. Gonzalez 2003). Whilst some studies (Lineweaver \\& Grether 2003) suggest that 100\\% of Sun-like stars may have planets, current surveys  appear to be heavily biased by the detection techniques used. The radial velocity technique has proved to be extremely successful over the past  decade in discovering  exoplanets, even though few planetary characteristics can be derived from the measurements. In contrast, planetary transits may yield many more properties, provided the host stars are well characterised. While transit surveys may monitor many orders of magnitude more stars in a comparatively shorter observing time than radial velocity surveys, they are not  devoid of problems. Blending, grazing eclipses and contamination by giant stars, among other problems, make some detections somewhat ambiguous (Dreizler et al. 2002, Brown et al. 2001, Drake 2003, Drake \\& Cook 2004). Clearly transit surveys have the ability to discover many interesting exoplanet candidates, but non-planetary events must be filtered out. The best way to perform such a separation is to obtain radial velocities at different phases along the orbit. Whilst this approach is becoming very successful (Dreizler et al. 2002, Konacki et al. 2003, Bruntt et al. 2003, Street et al. 2003, Bouchy et al. 2004), precision radial velocities for all transit candidates, especially those at faint magnitudes, would be far too time consuming, and so pre-selections must  be carried out. For instance, Dreizler et al. (2002) made a selection of the best planetary candidates by using low dispersion spectroscopy to refine stellar parameters, while Drake  (2003) used the gravity darkening dependence of the light curves. Another possibility is  near-IR photometry which can also be used to discard giant stars, stars with IR excess, and to better determine stellar parameters (e.g. see Ribas et al. 2003, for the case of bright stars with 2MASS photometry) and thus the radii of the companion objects. Ellipsoidal variability is also a good way to filter out the larger companions, although the correlation in the photometric time series limits in practice its applicability (Sirko \\& Paczy\\'nsky 2003).  The characterisation of the companions and their host stars is thus important for two main reasons. First, it allows a pre-selection of the most promising exoplanet candidates, and filters out giant stars and large companions. Second, the depth of the transits depends upon $R_c/R_*$, the ratio of the radius of the companion $R_c$ to the one from the host star $R_*$. Yet the most important and interesting quantity in the case of exoplanets is the radius $R_c$, as it allows inferences on the planet's composition and evolutionary history. Characterising the properties of the host star, and in particular its radius $R_*$ is therefore very important for transit candidates.  In this work we explore the possibility to characterise, through IR photometry and low dispersion spectroscopy, the properties of low-luminosity stars whose transits were discovered by the OGLE-III collaboration (Udalski et al. 2002b), in order to constrain astrophysical parameters of both stellar and companion objects, some of which may be exoplanets.  In Sections 2 and 3 we present the observations and data reduction. We describe the estimation of stellar parameters in Section 4 through a novel application of the surface brightness method. The inferred properties of the companions are discussed in Section 5. A brief discussion of the stellar populations observed in these fields is provided in Section 6 with the help of a model of the Galaxy. ",
        "conclusions": "In order to make an efficient pre-selection of possible exoplanets, we obtained $K$-band photometry for a sample of transit candidates detected by the OGLE collaboration in the CAR 104 and CAR 105 fields of the Galactic disc,  along with low-dispersion spectroscopy of 11 promising  candidates. By combining these new data with the optical $VI$ photometry from OGLE, we presented a novel application of the surface brightness relations to measure the colour excess of each transit star, its distance, and, more importantly, the effective temperature and radius, and hence to fully caracterise their properties. The method rejects automatically giant stars, and provide accurate solutions for main-sequence stars. A comparison with the temperatures derived from the low-resolution spectra shows an excellent agreement and validates the method. The new measures of stellar radii allow us to further refine  the radii of the companions.  We find that :  1. The combination of optical with near-IR photometry allows a very good estimate the stellar parameters for most stars and therefore to select the best planetary transit candidates in a very effective way, bypassing the need of high-resolution spectroscopy at this stage of the selection process.  2. We discard \\ogle{83}, \\ogle{89}, \\ogle{98},  \\ogle{102}, \\ogle{116}, \\ogle{118} and \\ogle{119} for being very probably giant stars.  3. Over 70\\% of the sub-sample contains low-mass companions whose radii are far too large to be exoplanets.  4. We select \\ogle{105}, \\ogle{109} and \\ogle{111} (full squares in Fig. 1, 13 and 14) as the most likely transit candidates to host exoplanets, as the inferred radii of their companions are below  $R_{c} / R_{\\odot} < 0.14$. As their periods are shorter than 3.5 days, these are new possible ``hot Jupiter''-like planets.  5. Recently, Pont et al. (2004) confirmed \\ogle{111} as an exoplanet with a mass of 0.53 $\\pm$ 0.11 M$_{Jup}$ and a radius of 1.0$^{1.13}_{0.94}$ R$_{Jup}$. Our calculated radius for \\ogle{111} is about 0.91 $\\pm$ 0.03, showing the reliability of our method to reject ``false positives'' and efficiently select promising extrasolar planetary transit candidates. \\\\ \\\\ A comparison with the Besan\\c{c}on model of the Galaxy shows a satisfactory agreement and that these stars are normal and representative of the thin disc population.  In the future it would be useful to obtain high-resolution spectra with UVES/VLT for the three candidates to measure radial velocities and to determine accurately their masses.\\\\"
    },
    "0410/astro-ph0410142_arXiv.txt": {
        "abstract": "The cosmological magnetic field is one of the important physical quantities which affect strongly the cosmic microwave background (CMB) power spectrum. Recent CMB observations have been extended to higher multipoles $l\\gtrsim$1000, and they resultantly exhibit an excess power than the standard model prediction in cosmological theory which best fits the Wilkinson Microwave Anisotropy Probe (WMAP) data at lower multipoles $l\\lesssim$900. We calculate the CMB temperature anisotropies generated by the power-law magnetic field at the last scattering surface (LSS) in order to remove the tension between theory and observation at higher multipoles and also place an upper limit on primordial magnetic field. In our present calculation we take account of the effect of ionization ratio exactly without approximation. This effect is very crucial to precisely estimate the effect of the magnetic field on CMB power spectrum. We consider both effects of the scalar and vector modes of magnetic field on the CMB anisotropies, where current data are known to be insensitive to the tensor mode which we ignore in the present study. In order to constrain the primordial magnetic field, we evaluate likelihood function of the WMAP data in a wide range of parameters of the magnetic field strength $|\\mathbf{B}|_\\lambda$ and the power-law spectral index $n_B$,  along with six cosmological parameters in flat Universe models, using the technique of the Markov Chain Monte Carlo(MCMC) method. We find that the upper limit at $2\\sigma$ C.L. turns out to be $|\\mathbf{B}_\\lambda|\\lesssim 3.9 $nG at 1 Mpc for any $n_B$ values, which is obtained by comparing the calculated result including the Sunyaev-Zeldovich(SZ) effect with recent WMAP data of the CMB anisotropies. ",
        "introduction": "Temperature and polarization anisotropies in CMB provide very precise information of the physical processes in the early Universe. However, recent new CMB data sets from the WMAP (Bennett, et al. 2003), the Arcminute Cosmology Bolometer Array Receiver (ACBAR; Kuo et al. 2004), and the Cosmic Background Imager (CBI; Mason et al. 2003) have indicated a potential discrepancy between theory and observation at higher multipoles $l \\ge 900$. The best-fit cosmological model to the WMAP data predicts the power spectrum which shows appreciable departure from those observed in balloon and interferometer experiments. This discrepancy cannot be explained by taking account of tuned standard cosmological parameters.  One possible interpretation of an excess power at high multipoles is a manifestation of the re-scattering of CMB photons by hot electrons in clusters known as the SZ effect (Sunyaev and Zeldovich 1980). Although there can be no doubt that some contribution from the SZ effect exists in the observed angular power spectrum, it has not yet been established conclusively that this is the only possible interpretation (Aghanim et al. 2001) of the small scale power. Indeed, the best value of the matter fluctuation amplitude to fit the excess power at high multipoles is near the upper end of the range of the values deduced by the other independent method (Bond et al., 2002; Komatsu and Seljak 2002). Toward the solution of this bothersome problem, it has recently been reported that the bump feature in the primordial spectrum gives  the better explanation for both CMB and the matter power spectra at small scales (Mathews et al., 2004).  The inhomogeneous cosmological magnetic field generated before CMB last scattering epoch is also a possible candidate to reconcile the tension in such higher multipoles. It excites an Alfven-wave mode in the baryon-photon plasma in the early Universe and induces small rotational velocity perturbations. Since the mode can survive on scales below those of the Silk damping during recombination (Jedamzik et al., 1998; Subramanian et al., 1998), it could be a new source of the CMB anisotropies in such small scales. Analytic expressions of temperature and polarization angular power spectra in rather larger angular scales ($l \\le 500$) based upon the thin LSS approximation were derived for both vector and tensor modes (Mack et al., 2002). Subramanian and Barrow (2002) considered the vector perturbations in the opposite limit of smaller angular scales. Too strong magnetic field in the early Universe, however, has a possibility of conflicting with the cosmological observations currently available. The  combination of those studies and current observations places a rough bound on the strength of primordial magnetic field $|\\mathbf{B}_{\\lambda}| < (1.0-10)$nG.  In order to compare the theoretical CMB anisotropies calculated by including a primordial magnetic field,  with observations much more precisely, we need to perform numerical calculations of the fully linearized equations with the use of realistic recombination history of the Universe. In particular, we first need to develop a numerical method to predict the theoretical spectrum for intermediate angular scales, in which the analytic approximation becomes inappropriate. The numerical study of only the scalar mode was introduced by Koh and Lee (2000), and they discussed the sensitivity of CMB to the primordial magnetic field strength. Lewis (2004) has recently shown the characteristic of a primordial magnetic field effect on CMB in vector and tensor modes, numerically. However, there is no systematic study to put constraint on parameters of primordial magnetic field by taking  account of both scalar and vector modes simultaneously.  The purpose of this Letter is to explore more completely the effects of a primordial magnetic field on the CMB anisotropies, and to place a new limit on the field strength together with many other cosmological parameters using the MCMC method in the analysis numerically. The present calculation differs from all previous discussions in the following important ways. First, we take account of the effect of ionization ratio exactly without approximation. Second, we consider the effect of scalar mode of magnetic field more precisely than the analysis of Koh and Lee (2000) by introducing the cutoff wave number physically consistent with the magneto-hydrodynamic (MHD) study of damping Alfven wave (Jedamzik et al., 1998; Mack et al., 2002). The effect of scalar mode of primordial magnetic field changes the CMB power spectrum significantly at larger scales for $l<900$ (Koh and Lee, 2000), while the effect of vector mode, which we also consider in the present study, dominates at smaller scales. These effects are very crucial to quantitatively estimate the effects of magnetic field on the CMB power spectrum over a wide range of multipoles $l$ for various scales. Third, it is for the first time to discuss the primordial magnetic field as a possible solution to remove the tension between theory and observation which exhibits in smaller angular scales of the CMB anisotropies. ",
        "conclusions": "Our study confirms that the effect of not only the vector mode but the scalar mode magnetic field plays the important role in order to solve the potential discrepancy of the CMB anisotropies between theory and observation at higher multipoles.  We show the result of our MCMC analysis with WMAP data in the two parameter plane $|\\mathbf{B}_{\\lambda}|$ vs. $n_B$ in Fig.\\ref{fig1}. Although the higher $l$ data from CBI(Mason et al., 2003), ACBAR(Kuo et al., 2004), and others may constrain the magnetic field strongly, we did not include those data in the present analysis because they have too large error bars. Therefore, the result is most sensitive to the WMAP data points at 500$\\le l \\le$900. Shown in Fig.\\ref{fig1} is the excluded region at 2$\\sigma$(95.4\\%) C.L as bounded by the thick solid curve, but we could not find the lower boundary of the allowed region at the same  2$\\sigma$(95.4\\%) C.L. Note that we find a very shallow minimum of the reduced $\\chi^2 \\simeq 1.08$ along the thick dashed line which is almost parallel to the upper 2$\\sigma$ boundary displayed as thick solid line is Fig.\\ref{fig1}. We can thus obtain the strength of primordial magnetic field $|\\mathbf{B}|_\\lambda \\lesssim$3.9nG(2$\\sigma$) at 1Mpc for the power-law spectral index $n_B\\sim1.1$. This upper limit is the most reliable and newest one to the primordial magnetic field because we consider all effects on the CMB anisotropies, i.e. the effect of ionization ratio, the SZ effect, and the both scalar and vector mode effects of the magnetic field, for the first time, in the present estimate of $|\\mathbf{B}_\\lambda|$ and $n_B$. In our numerical estimate of the magnetic field parameters, we continued the MCMC analysis until the other cosmological parameters converge well to the values listed in table 1. The inferred parameter values are not very different from those of Spergel et al. (2002), but the important fact is that we cannot find obvious degeneracies of magnetic field parameters with other cosmological parameters. We understand this for the following reasons. Primordial magnetic field is constrained by the cutoff scale of damping Alfven wave as indicated by Eq.(\\ref{eq:kd}) (Jedamzik et al. 1998; Mack et al. 2002). Since the cutoff scale is small enough compared with the mulitpoles $l \\ll$ 3000 for current available data which we are interested in, the manifestation of this effect is remarkably seen as a monotonically increasing contribution from both scalar and vector perturbations in the mulutipole region higher then $l\\sim$500 as far as $|\\mathbf{B}|_\\lambda \\lesssim$3.9nG. This sensitivity of the CMB power spectrum to the primordial magnetic field differs completely from those to the other cosmological parameters, which helps resolve the degeneracies among them.  There is, however, strong degeneracy between magnetic field strength $|\\mathbf{B}_\\lambda|$ and power spectral index $n_B$. At this moment, we can not resolve this strong degeneracy, because the WMAP data are constrained to lower $l\\le$ 900. The CBI and ACBAR data are in an interesting higher multipole region, but the error bars are too larger to resolve the degeneracy. Precise data for higher $l$, where the effect of the primordial magnetic field is especially strong, are highly desirable.  Let us shortly discuss the consistency of our upper limit $|\\mathbf{B}_\\lambda|\\lesssim 3.9nG(2\\sigma)$ with the other observational constraint on the magnetic field of the cluster of galaxies. Many astronomical observations indicate that the magnetic field strength in the present day cluster of galaxies is $\\sim (0.1-1)\\mu$G. Assuming isotropic collapse for cluster formation, one can straightforwardly estimate the magnetic field strength about $(1-10)$nG at the epoch of LSS. If the field strength is extremely larger than this critical value $\\sim (1-10)$nG, we need some damping process, and if it is smaller than $\\sim (1-10)$nG, on the other hand, we need amplification process. We do not know any viable physical process of damping the magnetic field, but there are several amplification processes proposed in literature. We found wide allowed region in the two parameter plane ($|\\mathbf{B}_{\\lambda}| , n$) in Fig.\\ref{fig1} for the primordial magnetic field which best fits the current CMB data. Their parameter values satisfy the observational constraint form the cluster of galaxies $|\\mathbf{B}_{\\lambda}|\\sim (1-10)$nG at the epoch of LSS.  To summarize, we studied the CMB fluctuation power spectrum by taking account of the scalar and vector modes from primordial magnetic field and the SZ effect. The likelihood analysis of the WMAP data indicates that the upper limit of magnetic field strength is $|\\mathbf{B}_\\lambda| \\lesssim 3.9$ nG (2$\\sigma$) at 1 Mpc for any power spectral indices $n_B$ and reasonable values of the other cosmological parameters."
    },
    "0410/astro-ph0410697_arXiv.txt": {
        "abstract": "{We present the first non-LTE atmosphere models for WR\\,stars that incorporate a self-consistent solution of the hydrodynamic equations.  The models account for iron-group line-blanketing and clumping, and compute the hydrodynamic structure of a radiatively driven wind consistently with the non-LTE radiation transport in the co-moving frame. We construct a self-consistent wind model that reproduces all observed properties of an early-type WC\\,star (WC5).  We find that the WR-type mass-loss is initiated at high optical depth by the so-called `Hot Iron Bump' opacities ({\\FeIX}--{\\sc xvi}).  The acceleration of the outer wind regions is performed by iron-group ions of lower excitation in combination with C and O. Consequently, the wind structure shows two acceleration regions, one close to the hydrostatic wind base in the optically thick part of the atmosphere, and another farther out in the wind. In addition to the radiative acceleration, the `Iron Bump' opacities are responsible for an intense heating of deep atmospheric layers.  We find that the observed narrow {\\OVI}-emissions in the optical spectra of WC\\,stars originate from this region.  By their dependence on the clumping factor we gain important information about the location where the density inhomogeneities in WR-winds start to develop. ",
        "introduction": "\\label{sec:intro}   The modeling of the atmospheres of Wolf-Rayet (WR) stars with their strong stellar winds underwent a strong development during the last two decades. Starting with pure-helium atmospheres \\citep{hil1:87,hil2:87,ham2:87}, the main goal was to augment the models by more complex atoms like C, N, and O \\citep[e.g.][]{hil1:89,koe1:95}, until a major step towards a realistic determination of stellar parameters and emergent flux distributions was performed by the inclusion of iron-group line-blanketing \\citep{hil1:98,hil1:99,gra1:02} and clumping \\citep[see][]{ham1:98}.  For the inclusion of line-blanketing major revisions of the model codes were necessary. In line-blanketed models the whole spectral range is covered by millions of overlapping lines. Due to the frequency coupling in expanding atmospheres, no sampling or re-ordering techniques can be applied \\citep[cf.\\ the technique of opacity distribution functions, e.g.][]{car1:79}, i.e., the complete radiation field must be calculated on a fine frequency grid, with each line transition at its proper place.  As a by-product of these numerically expensive calculations also the radiative acceleration on the wind material is obtained. This is of special interest for WR\\,stars, because it is not clear whether their strong winds can be driven by radiation pressure alone. In contrast to the thin winds of OB\\,stars, WR\\,winds show wind momenta $\\dot{M}\\vinf$ far above the single scattering limit $L/c$. This means that, on average, each photon leaving the star must be scattered several times to accelerate the wind to the observed terminal velocities.  The first results from line-blanketed atmosphere calculations have shown that the acceleration of the outer part of a WR\\,wind is easily explained by radiative driving, whereas the models failed completely in the deeper layers where the mass-loss is initiated \\citep{gra1:00,her1:01,gra1:02,hil1:03}. On the other hand, \\citet{nug1:02} and \\citet{lam1:02} demonstrated by a critical-point analysis that the {\\em mass-loss rates} of WR\\,stars can probably be maintained, if the critical point of the equation of motion is located in the optically thick part of the atmosphere at a temperature close to 160\\,kK.  In this case the mass-loss is initiated by the so-called `Hot Iron Bump' -- an increase of the Rosseland mean opacity, which is mainly due to the excitation of the Fe M-shell ions ({\\FeIX}--{\\sc xvi}).  Due to their method of analysis, however, it remained unclear whether the assumed conditions are actually met in WR\\,atmospheres, and if the wind can be further accelerated beyond the optically thick regime.  Due to their high wind densities and their hard radiation field, WR\\,stars develop extended atmospheres in extreme non-LTE that require sophisticated modeling techniques.  Their complex ionization structure is strongly affected by the interplay between different elements \\citep[see][]{gra1:02}. It is therefore necessary to employ a full non-LTE treatment and a radiation transport that correctly accounts for line-overlaps.  Moreover, WR\\,winds show high velocities at large optical depths, i.e., complete atmosphere models including an accurate solution of the temperature structure are required.  The numerical cost of our present models is therefore much higher than for already existent stellar wind models \\citep[e.g.\\ ][]{cas1:75,abb1:85,pau1:86,luc1:93,sch1:94,pau1:94,vin1:99,pau1:01,krt1:04} which are mostly focusing on the thin winds of OB\\,stars.  An additional important aspect for the driving of WR\\,winds, which is not included in any of the previous wind models, is the role of density inhomogeneities.  For WR\\,stars, these are directly observed as time-dependent variations of the line profiles \\citep{lep1:99,lep2:99,lep1:00,koe2:01}.  They are included in our models in the form of a clumping factor $D$, which lowers the empirically derived mass-loss rates by a factor $\\sqrt{D}$ \\citep[see][]{ham1:98}, i.e., it reduces the wind momentum problem for WR\\,stars.  Typically, values in the range of $D$\\,=\\,4--10 are obtained by a fit of the electron scattering wings of strong emission lines.  An effect of similar importance is the influence of clumping on the {\\em radiative acceleration} of the wind material.  This effect was discussed for the first time by \\citet{sch1:95} and \\citet{sch1:97} and is also detected in our models \\citep{gra2:03}. The clumping factor therefore plays a key role for the dynamics of WR\\,winds.  For the modeling of OB\\,star winds, two main approaches are presently followed. { In the first one, {\\em all} atomic populations are calculated in non-LTE (i.e., the statistical equilibrium equations are exactly solved) and the radiation field is treated in Sobolev approximation \\citep{pau1:01,krt1:04}. In the second, a Monte-Carlo technique is employed for the radiation transfer which particularly accounts for multiple scattering effects due to line-overlaps. The atomic populations of the iron-group elements, however, are treated in a modified nebular approximation \\citep{vin1:99}.}  As mentioned above, both methods are probably not suitable for the modeling of WR\\,winds, because the spectrum formation in WR\\,stars is dominated by non-LTE effects {\\em and} line-overlaps.  Despite of the approximations concerning the atomic populations, all spectral lines are treated as pure scattering opacities in the Monte-Carlo approach, i.e., they do not allow for photon leakage from one part of the spectrum to another \\citep[see][]{owo2:99,sim1:04}. This situation may be improved by the new technique of \\citet{luc1:02,luc1:03}. The Sobolev approach, on the other hand, introduces an artificial distinction between `line radiation' and `background continuum radiation' that causes problems in case of line overlaps.  A correction of the continuum radiation field is therefore necessary \\citep[see][]{pau1:94}.  In addition, the {\\em first order} Sobolev approximation breaks down in the acceleration region near the transsonic wind base -- even for thin stellar winds \\citep{sel1:93,owo1:99}.  The hydrodynamic structure of O\\,star winds was obtained for the first time by \\citet[][{}\\,later on CAK]{cas1:75}, who employed a parametrization of the radiative acceleration by means of the so-called force multiplier function. This method has been significantly improved by \\citet{pau1:86} and describes successfully the observed properties of OB\\,star winds { \\citep[except for dense winds where multiple scattering effects become important, see] []{lam1:93,pul1:96,vin1:00}. However, due to the parametrization by only two parameters, the line force is subject to an averaging process, i.e., the hydrodynamic equations are not strictly solved.}  Again, for the highly structured winds of WR\\,stars, this approach is probably not sufficient.  { The Monte-Carlo models by \\citet{vin1:99} are based on an even less accurate method where a consistent {\\em global} energy budget is provided for a {\\em prescribed} velocity field.}  \\citet{krt1:04} have recently improved this situation by a full local solution of the equation of motion {\\em and} the energy equation, i.e., their models provide a consistent velocity and temperature structure. {Nevertheless, due to the neglect of multiple scatterings, these models are limited to weaker winds.}  In the present work we overcome the problems explained above by combining full non-LTE atmosphere models with the equations of hydrodynamics.  By iteration between the non-LTE radiation transport and the hydrodynamics we derive $\\dot{M}$, $\\velo(r)$, $T(r)$, {\\em and} the non-LTE population numbers consistently with the radiation field in the co-moving frame (CMF), i.e., we solve the radiation transport, the energy equation, and the equation of motion simultaneously.  Moreover, clumping is accounted for. The resultant models describe the conditions in WR\\,atmospheres correctly, and provide synthetic spectra, i.e., they allow for a direct comparison with observational material. Nevertheless, simplifying assumptions are necessary also in our models.  These concern especially the assumption of a constant Doppler broadening velocity throughout the atmosphere and the omission of opacities, partly due to the neglect of trace elements like Ne, Ar, S, or P and partly due to incompleteness of the available data { (for ions above ionization stage {\\sc ix}, no fine-structure data, and no atomic data for elements with $Z > 26$ are available)}.  As already mentioned above, the critical-point analysis by \\citet{nug1:02} revealed that the mass-loss of early-type WR\\,stars is probably driven by the `Iron Opacity Bump' which is mainly due to line opacities of the Fe M-shell ions ({\\FeIX}--{\\sc xvi}). By inclusion of these ions in our models we are able to check whether the assumptions of \\citeauthor{nug1:02} (critical point at high optical depth and high temperature) are correct, and most importantly, if the high wind acceleration can be maintained in the outer part of the wind.  { In the Sect.\\,\\ref{sec:atmos} of the present work we review our standard atmosphere models, followed by a series of test calculations to clarify the role of the `Iron Bump' opacities and the clumping factor (Sect.\\,\\ref{sec:testcal}).  Then we describe the combination of our atmosphere models with the hydrodynamic equations (Sect.\\,\\ref{sec:hydromod}), and present a self-consistent hydrodynamic atmosphere model that reproduces all observed characteristics of an early-type WC\\,star.  Finally, we summarize our conclusions in Sect.\\,\\ref{sec:conclusions}.} ",
        "conclusions": "\\label{sec:conclusions}  In the present work we have constructed the first hydrodynamically consistent model for a WR-type (WC\\,5) stellar wind, driven by radiation pressure.  For this purpose we have combined our non-LTE atmosphere models with the hydrodynamic equations. On this way we obtain the first models for radiatively driven winds that incorporate a full non-LTE radiation transfer in the CMF, i.e., they allow for a realistic treatment of all important radiative effects throughout the whole atmosphere, including large optical depths.  Moreover, the models account for the important effects of clumping.  By inclusion of the Fe M-shell ions (\\FeIX--\\FeXVI) in our models, the radiative acceleration in deep atmospheric layers is considerably enhanced.  We show that, for suitable stellar parameters, the radiative force on these opacities is able to drive an optically thick wind. In particular, very high stellar temperatures ($T_\\star\\approx 140\\,\\mathrm{kK}$) are needed to maintain the wind acceleration above the sonic point. In accordance with \\citet{nug1:02}, the sonic point in our hydrodynamic model is located at large optical depth {\\em below} the `Hot Iron Bump' at a very high temperature of $T_{\\rm s} = 199\\,{\\rm kK}$.  The resultant wind structure shows two acceleration regions.  A steep velocity increase in deep atmospheric layers due to the Fe M-shell ions, and a shallow increase in the outer part of the wind due to iron-group ions of lower excitation in combination with C and O. The latter is in accordance with the shallow velocity laws deduced by \\citet{lep1:99} from line-profile variations.  Despite of the fact that the wind is optically thick at the sonic point, we find strong deviations from the standard CAK theory also in the outer wind regions. Here especially the reaction of the radiative force on changes of $\\velo'$ shows an abnormal behavior, which is presumably due to complex radiative processes like line emission cascades after recombination, or extreme line overlaps. Moreover, we find a strong dependence of the radiative force on wind clumping, an effect which is neglected in previous models for radiatively driven winds.  The emergent flux, as calculated from our hydrodynamic model, clearly shows the spectral characteristics of an early-type WC\\,star, i.e., the wind density and the terminal wind velocity, which are both obtained self-consistently from the hydrodynamics, are of the correct order of magnitude.  Moreover, we are able to reproduce the observed narrow {\\OVI} emissions in the optical spectrum of WC\\,stars. These features are excited by a combination of heating due to the Fe M-shell ions and the high stellar temperature, which is necessary for our wind modeling.  The {\\OVI} lines also react very sensitively on the detailed clumping structure, i.e., they help to trace deep atmospheric layers where the clumping of WC\\,star winds sets in.  When our results are compared to stellar structure calculations, the question about the location of the `surface' of a WR\\,star gains more and more in importance. Stellar structure models for WR\\,stars show an outer convection zone, which largely determines the extension of the stellar envelope. The convection however, is caused by the same increase of the mean opacity that causes the driving of the WR\\,wind in our models. A temperature around 200\\,kK, corresponding to our sonic-point temperature, is only reached deep inside the stellar envelope at the {\\em bottom} of the outer convection zone.  For a full understanding of the structure and the mass-loss of WR\\,stars models are therefore needed, that combine stellar structure and atmosphere calculations and allow for convection under dynamic conditions."
    },
    "0410/astro-ph0410374_arXiv.txt": {
        "abstract": "Several authors recently explored the Black-Hole mass ($M_{BH}$) vs. stellar velocity dispersion ($\\sigma_*$) relationship for Narrow Line Seyfert 1 galaxies (NLS1s). Their results are more or less in agreement and seem to indicate that NLS1s fill the region below the fit obtained by \\citet{tret02}, showing a range of $\\sigma_*$ similar to that of Seyfert 1 galaxies, and a lower $M_{BH}$. Until now the [O\\,{\\sc iii}] width has been used in place of the stellar velocity dispersion, but some indications begin to arise against the effectiveness of the gaseous kinematics in representing the bulge potential at least in NLS1s. \\citet{bz04b} stressed the urgency to produce true $\\sigma_*$ measurements. Here we present new stellar velocity dispersions obtained through direct measurements of the Ca\\,{\\sc ii} absorption triplet ($\\sim 8550$ \\AA) in the nuclei of 8 NLS1 galaxies. The resulting $\\sigma_*$ values and a comparison with $\\sigma_{\\rm [O~III]}$ confirm our suspects (see \\citealt{bo04}) that [O\\,{\\sc iii}] typically overestimates the stellar velocity dispersion and demonstrate that NLS1s follow the $M_{BH}-\\sigma_*$ relation as Seyfert 1, quasars and non-active galaxies. ",
        "introduction": "Recently great interest was addressed to the study of the correlation between the mass of nuclear supermassive Black-Holes ($M_{BH}$) and the stellar velocity dispersion ($\\sigma_*$) of their hosting bulges. This relation has been obtained for nearby non-active galaxies by several authors \\citep{gebet00a, fermer00, tret02, mahu03}, supporting the theory that the growth of the BH is bound to the galaxy formation. According to \\citet{dima03} this correlation, not set in primordial structures but fully established only at low redshift, can be understood if there is a simple linear relation between the total gas mass in galaxies and their $M_{BH}$.\\\\ \\citet{gebet00b} included in their sample seven Active Galactic Nuclei (AGNs) and found that also these objects followed a similar $M_{BH}-\\sigma_*$ correlation of non-active galaxies. Other authors have then confirmed the validity of this result for samples of Seyfert 1 galaxies (S1s) and quasars \\citep{ne00, fe01, wan02, bo03, sh03}. On the contrary, no consensus has been reached until now about the case of the Narrow-Line Seyfert 1 galaxies (NLS1s). These AGNs have spectroscopic properties slightly different from those of the classical S1s, as for example, narrower permitted emission lines in the optical wavelength domain and steeper power-law X-rays continua, which suggest that NLS1s are hosting BHs with smaller masses accreting at high rates, close to the Eddington limit. This hypotesis would imply that NLS1s have also less massive bulges than classical S1s, of course assuming that their $M_{BH}$ are correlated to the physical properties of their hosting bulges. Indeed, \\citet{walu01}, \\citet{wan02} and \\citet{wan04} found that there is no clear difference in the $M_{BH}-\\sigma_*$ relation between narrow-line, broad-line AGNs and non-active nearby galaxies, while opposite results were obtained by \\citet{matet01}, \\citet{bz04a}, \\citet{gm04}, and \\citet{bo04}. The main problem is that $\\sigma_*$ in AGNs is available for few S1s and only 3 NLS1s: Mrk~110 ($90\\pm 7$ \\kms) by \\citet{fe01}, NGC~4051 ($88\\pm 13$ \\kms) by \\citet{nel95}, and Mrk~766 ($106 \\pm 40$ \\kms) by \\citet{JB00}. In place of $\\sigma_*$, the [O\\,{\\sc iii}] emission linewidth was extensively used as a representation of the bulge velocity dispersion.  Here we present new direct determinations of  $\\sigma_*$ in the nuclei of 8 NLS1 galaxies, obtained using the Ca\\,{\\sc ii} absorption triplet (8498, 8542, 8662 \\AA). The targets were extracted from \\citet{bo04}, except for Mrk~766. In Section 2 we present the spectroscopic observations and the data reduction; in Section 3 we give a detailed description of the methods we used to measure stellar and gaseous kinematics, and we present the results. Finally, in Section 4 we discuss the implications of our results on the validity of the FWHM([O\\,{\\sc iii}])/2.35=$\\sigma_*$ and $M_{BH}-\\sigma_*$ relations for NLS1 galaxies.    \\begin{figure*} \\includegraphics[angle=0,width=80mm]{fig1a.eps} \\includegraphics[angle=0,width=80mm]{fig1b.eps} \\caption{The rest-frame spectra of the 8 NLS1 galaxies showing the Ca\\,{\\sc ii} triplet absorption lines. The template stars (HR 294 and HD 11559) are also plotted for comparison.} \\label{fig1} \\end{figure*} ",
        "conclusions": "Several authors tested the M$_{BH}-\\sigma_*$ correlation validity for AGNs assuming that the gaseous component could reasonably trace the nuclear gravitational potential. In particular, the [O\\,{\\sc iii}] width was used as surrogate of the stellar velocity dispersion, essentially for a practical reason: the [O\\,{\\sc iii}] emission line is usually bright in AGNs and therefore easily measurable even in case of distant objects. On the contrary stellar absorptions require longer integration time and/or large telescopes to reach the S/N ratio necessary to give a reliable estimate of the stellar velocity dispersion. The FWHM([O\\,{\\sc iii}]) to $\\sigma_*$ conversion was proposed by \\citet{nel96}, who observed a large sample of AGNs in two spectral ranges, one including H$\\beta$, [O\\,{\\sc iii}] emission lines and Mg b triplet absorption, and the other containing the Ca\\,{\\sc ii} triplet absorption. The data of 66 objects were plotted showing a moderate correlation with a significant scatter. More recently \\citet{JB00} obtained similar results on a sample of Seyfert 1 and 2 galaxies, demonstrating that the majority of the galaxies form a cloud of points around the 1:1 line of the gas vs. stellar velocity dispersion plot, with very weak or even no correlation.  None the less, the use of the [O\\,{\\sc iii}] width in place of direct $\\sigma_*$ measurements led \\citet{bz04a}, \\citet{gm04}, and \\citet{bo04} to demonstrate that Seyfert 1 galaxies agree with the M$_{BH}-\\sigma_*$ relation given by \\citet{tret02}, even if with a large scatter of $\\sigma_*$ (from 50 to more than 500 \\kms), while NLS1 galaxies do not. Moreover, NLS1s appear sistematically arranged below the Tremaine et al. fit. Interestingly, this result was already included in \\citet{bo03}, who carried out an analysis of the $M_{BH}$--[O\\,{\\sc iii}] $\\lambda$5007 width relation for 107 low--redshift radio--quiet QSOs and Seyfert 1 galaxies and found that this correlation is real, but the scatter is large, about 0.2 dex. Really, it is easy to verify that the narrow line AGNs (FWHM(H$\\beta)<2000$ \\kms) of his sample are mostly located below the Tremaine et al. fit. Therefore, authors were induced to believe that NLS1s do not follow the BH-bulge correlation as Seyefert 1, quasars and normal galaxies. But, very recently, \\citet{bz04b} obtained a good consistency in BH mass estimates of NLS1 galaxies by using H$\\beta$ and soft X-ray luminosity, while no agreement was found with the [O\\,{\\sc iii}] linewidth. Their result casts doubt on the reliability of [O\\,{\\sc iii}] as indicator of stellar velocity dispersion, at least for NLS1s. And, indeed it should be mentioned that the FWHM([O\\,{\\sc iii}])/2.35=$\\sigma_*$ relation was obtained by \\citet{nel96} for a sample of Seyfert 1 and 2 galaxies, but it was never verified for NLS1s.  The situation clearly required new direct $\\sigma_*$ measurements. We obtained these values for 8 NLS1s observing their Ca\\,{\\sc ii} triplet. The choice of this spectral range was induced by the fact that NLS1s are very often characterized by large Fe\\,{\\sc ii} emission multiplets (4400--4700 \\AA\\ and 5150--5350 \\AA) in correspondence of the absorption lines commonly used to estimate $\\sigma_*$, like Mg I $\\lambda$5175 and Fe I $\\lambda$5269. Unfortunately, as shown by \\citet{pers88}, NLS1s with strong Fe\\,{\\sc ii} are likely to have Ca\\,{\\sc ii} in emission rather than in absorption.  By comparing stellar and gaseous kinematics we found that the FWHM([O\\,{\\sc iii}]) /2.35=$\\sigma_*$ relation does not seem to hold for NLS1s, which are arranged mostly at values of $\\sigma_{\\rm [O~III]}$ larger than expected. Contrary to the recent findings, the direct $\\sigma_*$ measurements span a narrower range typically lower than 100 \\kms, and such range is in agreement with the stellar velocity dispersions expected by the $M_{BH}-\\sigma_*$ relation given by \\citet{tret02} in case of $M_{BH}\\sim10^6-10^7~M_{\\sun}$. Indeed, with our $\\sigma_*$ determinations, NLS1s are arranged around this relationship, suggesting that in NLS1 galaxies BH and bulge correlate, and lower mass BHs correspond to lower stellar velocity dispersions, or in other words lower mass bulges. We stress that this result is in agreement with our previous work \\citep{bo04}, where we showed that NLS1s are mostly confined to the lower ranges of the M$_{BH}$-L$_{bulge}(B)$ plane, and are well correlated to Seyfert 1 galaxies and quasars. The fact that the stellar velocity dispersion in place of the gaseous one yields a tighter relation when correlated with BH mass, is essentially caused by the star kinematics which is more representative of the bulge gravitational potential than the gas in AGNs. Even if the gas motion is mainly controlled by the mass of the bulge, several observations in the past showed that the base of the line profiles emitted by the Narrow Line Region (NLR) in many active galaxies is hardly reproduced simply assuming a pure virial motion in the potential of the host galaxy \\citep[][ and references therein]{veil91}. Other important factors or processes can be involved, and these additional processes can be at the origin of the larger scatter observed when the FWHM([O III]) is used as surrogate of $\\sigma_*$. The interaction bewteen the gas and the ejected radio plasma observed in many Seyfert galaxies may strongly influence the gas kinematics and contribute with outflowing motions of the NLR clouds \\citep{nel96}. Moreover, supersonic winds generated by the active nucleus are believed to accelerate the smaller clouds of the NLR located closer to the central engine and sometimes dominate their virial velocity \\citep{smith93}. This effect could be more important in NLS1 galaxies, since strong nuclear winds are expected in case of high accretion rates \\citep{lh00}. In addition, it was pointed out by \\citet{nel96} that a slight tendency exists for barred and/or disturbed active galaxies to have broader [O\\,{\\sc iii}] emission lines. In conclusion, we obtained results with important implications, which strongly encourage additional observations to enlarge the sample of the observed galaxies."
    },
    "0410/astro-ph0410360_arXiv.txt": {
        "abstract": "We present a $2.5\\sigma$ detection of the Integrated Sachs-Wolfe (ISW) effect and discuss the constraints it places on cosmological parameters. We cross-correlate microwave temperature maps from the \\WMAP satellite with a 4000 deg$^{2}$ luminous red galaxy (LRG) overdensity map measured by the Sloan Digital Sky Survey. These galaxies have accurate photometric redshifts ($\\Delta z \\sim 0.03$) and an approximately volume limited redshift distribution from $z\\sim 0.2$ to $z\\sim 0.6$ well suited to detecting the ISW effect. Accurate photometric redshifts allow us to perform a reliable auto-correlation analysis of the LRGs, eliminating the uncertainty in the galaxy bias, and combined with cross correlation signal, constrains cosmological parameters -- in particular, the matter density. We use a minimum variance power spectrum estimator that optimally weights the data according to expected theoretical templates. We find a $2.5 \\sigma$ signal in the Ka, Q, V, and W \\WMAP bands, after combining the information from multipoles $2 \\le l < 400$. This  is consistent with the expected amplitude of the ISW effect, but requires a lower matter density than is usually assumed: the amplitude, parametrized by the galaxy bias assuming $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ and $\\sigma_8=0.9$, is $b_{g} = 4.05 \\pm 1.54$ for V band, with similar results for the other bands. This should be compared to $b_g = 1.82 \\pm 0.02$ from the auto-correlation analysis. These data provide only a weak confirmation ($2.5 \\sigma$) of dark energy, but provide a significant upper limit: $\\Omega_{\\Lambda}=0.80_{-0.06}^{+0.03} (1\\sigma) _{-0.19}^{+0.05} (2\\sigma)$, assuming a cosmology with $\\Omega_{M}+\\Omega_{\\Lambda}=1$, $\\Omega_{b} = 0.05$, and $\\sigma_{8}=0.9$, and $w=-1$. The weak cross-correlation signal rules out low matter density/high dark energy density universes and, in combination with other data, strongly constrains models with $w<-1.3$. We provide a simple prescription to incorporate these constraints into cosmological parameter estimation methods for $(\\Omega_{M}, \\sigma_{8},w)$. We find no evidence for a systematic contamination of ISW signal, either from Galactic or extragalactic sources, but we do detect some large statistical fluctuations on smaller scales that could affect analyses without the template weighting. ",
        "introduction": "\\label{sec:intro}  The cosmic microwave background (CMB) observed by the \\WMAP satellite \\cite{2003ApJ...583....1B} has been a font of information for cosmology. The positions of the acoustic oscillations in the angular power spectrum of temperature fluctuations imprinted at the last scattering surface measures the curvature of the universe to unprecedented accuracy. In addition, a careful modelling of the power spectrum combined with other cosmological probes have allowed extremely precise measurements of the parameters of the $\\Lambda$CDM model. Subsequent measurements of the power spectrum by \\WMAP, as well as future experiments such as the \\PLANCK satellite will further constrain the space of cosmological models, as well as test our understanding of the physics underlying recombination.  CMB temperature fluctuations are not just sourced by density fluctuations at the last scattering surface; anisotropies also arise due to the interaction of photons with hot electrons in galaxies and clusters (the Sunyaev-Zeldovich \\cite[SZ,][]{1980ARA&A..18..537S} and kinetic-SZ \\cite{1980MNRAS.190..413S} effects), as well as with the gravitational potentials along their propagation path (the Integrated Sachs-Wolfe \\cite[ISW,][]{1967ApJ...147...73S} effect and gravitational lensing).  The ISW effect results from the red- (or blue-)shifting of CMB photons as they propagate through gravitational potential wells. If these potentials did not evolve, then the blueshift gained falling into a potential well would exactly cancel the redshift emerging from the well; evolving potentials spoil this cancellation. This is significant only on the largest scales, both because the power in the fluctuations of the potential is largest on large  scales, and because integrating along the line of sight cancels out the temperature fluctuations (a photon is as likely to be redshifted as blueshifted).  Although the above effects are too small to unambigously detect in current CMB data (for eg. \\cite{2004astro.ph..4545H}), one can attempt to isolate them by cross-correlating the temperature maps with suitable tracer populations \\cite{2003ApJ...597L..89F,2004MNRAS.350L..37F, 2003astro.ph..7335S, 2004PhRvD..69h3524A, 2004Natur.427...45B, 2004ApJ...608...10N,2004astro.ph..4348B, 2004astro.ph..6004H}. This paper is a continuation of these efforts, attempting to detect the ISW effect, but also goes one step further in making the ISW signal useful for cosmological parameter analysis. We combine the cross-correlation analysis with a galaxy auto-correlation measurement to remove the uncertainty due to galaxy bias. The auto-correlation analysis is only reliable when the galaxy population has well determined redshifts: we argue below that this is indeed the case for the sample used here. This paper does not focus on the auto-correlation analysis, so we perform a restricted analysis where we assume the shape of power spectrum and only constrain the amplitude. This is sufficient for the current purposes, as the errors are dominated by the statistical uncertainties in the ISW effect. A detailed auto-correlation analysis will be presented in a future paper.  A galaxy catalog used for detecting the ISW effect must satisfy several criteria. Since the ISW signal is only detectable on the largest scales, the catalog must cover as large an area as possible. In addition, a large number of galaxies is necessary to remove noise from random Poisson fluctuations. Finally, to make theoretical interpretation possible, these galaxies must be drawn from a uniform population with a well characterized redshift distribution.  Photometrically selected luminous red galaxies (LRGs) \\cite{2004astro.ph..7594P} from the Sloan Digital Sky Survey (SDSS)  \\cite{2000AJ....120.1579Y} are a promising candidate. These galaxies are amongst the most luminous in the Universe and so probe cosmologically interesting volumes, including the transition from matter to dark energy domination that sources the ISW signal. These galaxies are old stellar populations with very uniform spectral energy distributions that make uniform selection and accurate photometric redshifts possible. Finally, the deep and wide-angle imaging of the sky by the SDSS allows these galaxies to be selected over a large area with high densities.  Our goal is to extract the maximal amount of information on the ISW effect from the data. To do so, we employ quadratic estimator methods, which combine the data in an optimal way in the presence of noise and incomplete sky coverage \\cite{1997PhRvD..55.5895T,1998ApJ...506...64S}. In addition, we weight the information on different scales according to the expected scale dependence of the signal: the ISW effect is expected to dominate on large scales and be absent on small scales, so we weight these different scales accordingly. We apply these techniques to the latest SDSS data reductions, ensuring the largest available SDSS sky coverage to date.  Detecting an ISW effect would provide evidence of either a cosmological constant, quintessence or curvature independent of supernovae observations of the luminosity distance \\cite{2004ApJ...607..665R}. We investigate the cosmological implications of an ISW detection within the assumption of a flat universe, and in particular, on its ability to constrain the matter density. The rapid increase of the ISW signal with decreasing matter density allows us to place strong lower bounds on the matter density and suggests an alternative probe of the properties of the dark energy.  The paper is organized as follows : Sec.~\\ref{sec:theory} introduces the ISW effect, developing the formalism to compute the predicted ISW signal. Sec.~\\ref{sec:data} describes the \\WMAP and SDSS data used in this paper. Sec.~\\ref{sec:estimate} and Sec.~\\ref{sec:results} then discuss the cross correlation of these data, and the results obtained. Systematic effects are discussed in Sec.~\\ref{sec:systematics}, while cosmological implications are considered in Sec.~\\ref{sec:cosmology}. We conclude in Sec.~\\ref{sec:discussion}. Unless otherwise specified, we assume a $\\Lambda$CDM cosmology with $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda} =0.7$, $H_{0} = 100 h~\\rm{km/s/Mpc}$ and $\\sigma_{8}=0.9$. Where necessary, we use $h=0.7$. ",
        "conclusions": ""
    },
    "0410/astro-ph0410683_arXiv.txt": {
        "abstract": "Detailed 4.8- and 8.6-GHz radio images of the entire Large Magellanic Cloud with half-power beamwidths of 33$\\arcsec$ at 4.8 GHz and 20$\\arcsec$ at 8.6 GHz have been obtained using the Australia Telescope Compact Array. A total of 7085 mosaic positions were used to cover an area of 6$^{\\circ}$ on a side.  Full polarimetric observations were made.  These images have sufficient spatial resolution ($\\sim$8 and 5 pc, respectively) and sensitivity (3 $\\sigma$ of 1 mJy beam$^{-1}$) to identify most of the individual SNRs and H II regions and also, in combination with available data from the Parkes 64-m telescope, the structure of the smooth emission in that galaxy.  In addition, limited data using the sixth antenna at 4.5 to 6-km baselines are available to distinguish bright point sources ($<$ 3 and 2 arcsec, respectively) and to help estimate sizes of individual sources smaller than the resolution of the full survey.  The resultant database will be valuable for statistical studies and comparisons with x-ray, optical and infrared surveys of the LMC with similar resolution. ",
        "introduction": "The Large Magellanic Cloud (LMC) is an excellent laboratory for investigation of the processes of star birth and death, and particularly the interactions of the stellar population and various components of the interstellar medium (ISM).  To measure many properties of the sources, such as luminosity and physical size, we must know their distances. This is  hard to do for objects in the Milky Way because most of the measured quantites are themselves dependent on the distance and there is also obscuring material within the Milky Way itself which can blank our view and distort optical measurements.  In the LMC  all sources of emission are at essentially the same known distance of 50~kpc.  Radio continuum observations can be used to measure both thermal emission from H~II regions and polarized synchrotron emission from supernova remnants (SNRs).  The thermal emission has a nearly flat radio spectrum whereas the synchrotron emission, from relativistic electrons accelerated in shocks and the magnetic fields of SNRs, has a power-law spectrum with brighter emission at lower frequencies. Three important sources of stellar energy in the ISM are UV radiation, fast stellar winds, and supernovae.  Thus the radio observations supply a vital component in studies of these energy sources, by tracing the contributions from the photoionized and shock-ionized material and (through polarimetry) showing the orientation of magnetic field lines, even deep within highly obscured regions like molecular clouds.  In combination with optical line emission \\citep{s99a} and X-ray images \\citep{s99b} from the new generation surveys at those wavelengths, high-resolution radio continuum imaging will help to construct a complete physical picture of stellar feedback in the Magellanic Clouds.  We have completed 4.8- and 8.6-GHz observations of the LMC to produce images $6^{\\circ}$ on a side, including all the bright features in that galaxy.  They have resolutions of 33\\arcsec ~at 4.8~GHz (20\\arcsec ~at 8.6~GHz) or about 8 pc (5~pc) at the 50~kpc distance of the LMC.  With these moderate-resolution data in hand, one can choose smaller areas for followup imaging, and reuse these relatively short-spacing data.  The noise level of the system provides 3-$\\sigma$ measurements at 1~mJy beam$^{-1}$ and if the source fills the 33$\\arcsec$ beam at 4.8 GHz, this value corresponds to a surface brightness at of 4$\\times$10$^{-22}$ W m$^{-2}$ Hz$^{-1}$ sr$^{-1}$ or for a mean SNR spectral index of $-0.5$, a 1 GHz surface brightness of 1$\\times$10$^{-21}$ W m$^{-2}$ Hz$^{-1}$sr$^{-1}$. While this limit is ten times the brightness of the faintest SNR in the Milky Way (G182.4+4.3, \\cite*{k98}) it would include 90\\% of all the Milky Way SNRs cataloged by Green (http://www.mrao.cam.ac.uk/surveys/snrs/). ",
        "conclusions": ""
    },
    "0410/astro-ph0410156_arXiv.txt": {
        "abstract": "We report monitoring observations of 20 high-luminosity AGN, 12 of which are radio-quiet quasars (RQQs).  Intra-night optical variability (INOV) was detected for 13 of the 20 objects, including 5 RQQs.  The variations are distinctly stronger and more frequent for blazars than for the other AGN classes. By combining these data with results obtained earlier in our program, we have formed an enlarged sample consisting of 9 BL Lacs, 19 RQQs and 11 lobe-dominated radio-loud quasars. The moderate level of rapid optical variability found for both RQQs and radio lobe-dominated quasars argues against a direct link between INOV and radio-loudness. We supplemented the present observations of 3 BL Lacs with additional data from the literature. In this extended sample of 12 well observed BL Lacs, stronger INOV is found for the {\\it EGRET} detected subset. ",
        "introduction": "In a series of papers since 2003, we have reported results of a program to search for intra-night optical variability (INOV; see Wagner \\& Witzel 1995 for a review), often called microvariability, in a sample of 26 optically luminous active galactic nuclei (AGN), using the 104-cm Zeiss telescope of the Aryabhatta Research Institute of Observational Sciences (ARIES), Naini Tal, India (Gopal-Krishna et al.\\ 2003, GK03; Stalin et al.\\ 2004a, St04a; Sagar et al.\\ 2004, Sa04; Stalin et al.\\ 2004b, St04b; see also Stalin 2003). These objects belong to the four major classes of luminous AGN, namely, radio-quiet QSOs (RQQs), radio-loud lobe-dominated quasars (LDQs), radio-loud core-dominated quasars (CDQs), and BL Lac objects (BLs). The sample selection was such that the four classes are reasonably well matched in the $z$--$M_B$ plane, with $z$ ranging from 0.17 to 2.2 and $M_B$ ranging from $-24.3$ to $-30.0$ (taking $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0 = 0$).     \\begin{table} \\caption{The sample of 20 optically luminous AGN monitored in the present work} \\begin{tabular}{lllllrr} \\hline &       &        &         &          &              &       \\\\ IAU Name    &  Type & ~~$B$    & ~$M_B$  & ~~ $z$   & $P_{opt}$$^\\dag$ & log $R^{\\ast}$$^{\\ddag}$\\\\ &       &        &         &          &    (\\%)~~          &       \\\\ &       &        &         &          &              &       \\\\ \\hline &       &        &         &          &              &       \\\\ 0003$+$158  &  LDQ  & 16.51  & $-$25.7 &  0.450   &  0.65  &     2.6  \\\\ 0025$+$307  &  LDQ  & 15.79  & $-$26.7 &  0.500   &  ---   &     1.8  \\\\ 0043$+$039  &  RQQ  & 16.00  & $-$26.0 &  0.385   &  0.27  &  $-$0.7  \\\\ 0806$+$315  &  BL   & 15.70  & $-$25.0 &  0.220   &  ---   &     1.7  \\\\ 0824$+$098  &  RQQ  & 15.50  & $-$25.6 &  0.260   &  ---   &     0.5 \\\\ 0832$+$251  &  RQQ  & 16.10  & $-$25.5 &  0.331   &  ---   &     0.1 \\\\ 0846$+$513  &  CDQ  & 16.28  & $-$29.4 &  1.860   &  ---   &     2.2  \\\\ 0850$+$440  &  RQQ  & 16.40  & $-$26.1 &  0.513   &  ---   & $<-$0.5 \\\\ 0931$+$437  &  RQQ  & 16.47  & $-$25.8 &  0.456   &  ---   &     0.1 \\\\ 0935$+$416  &  RQQ  & 16.31  & $-$29.6 &  1.966   &  ---   & $<-$0.7 \\\\ 0945$+$438  &  RQQ  & 16.28  & $-$24.5 &  0.226   &  ---   & $<-$0.1  \\\\ 1029$+$329  &  RQQ  & 16.00  & $-$26.7 &  0.560   &  ---   & $<-$0.6  \\\\ 1418$+$546  &  BL   & 16.17  & $-$23.7 &  0.152   &  7.5   &     3.1  \\\\ 1422$+$424  &  RQQ  & 16.42  & $-$25.1 &  0.316   &  ---   & $<-$0.4  \\\\ 1425$+$267  &  LDQ  & 15.78  & $-$26.0 &  0.366   &  1.9   &     2.0  \\\\ 1444$+$407  &  RQQ  & 15.45  & $-$25.7 &  0.267   &  0.4   & $-$1.1  \\\\ 1522$+$101  &  RQQ  & 16.20  & $-$28.4 &  1.324   &  0.3   & $<-$0.7  \\\\ 1553$+$113  &  BL   & 15.00  & $-$26.8 &  0.360   &  ---   &     2.2  \\\\ 1631$+$395  &  LDQ  & 16.48  & $-$27.8 &  1.023   &  1.1   &     1.6  \\\\ 1750$+$507  &  RQQ  & 15.80  & $-$25.6 &  0.300   &  ---   &     0.7  \\\\ &       &        &         &          &        &       \\\\  \\hline \\end{tabular}  \\hspace*{-0.2cm} $^\\dag$References for optical polarization: Wills et al.\\ 1992;  ~~Berriman et al.\\ 1990; --- for no data available  \\hspace*{-0.2cm} $^\\ddag$$R^{\\ast}$ is the K-corrected ratio of the 5 GHz radio to 2500 \\AA -band  ~~optical flux densities (Stocke et al. 1992); references for   \\hspace*{-0.0cm}~~radio fluxes: V{\\'e}ron-Cetty \\& V{\\'e}ron 2001; NVSS (Condon et al.\\ \\hspace*{-0.0cm}~~1998); FIRST (Becker et al.\\ 1995; Bauer et al.\\ 2000)  \\end{table} \\begin{table} \\caption{Positions and magnitudes of the comparison stars} \\begin{tabular}{llrrll} \\hline &       &             &               &             &               \\\\ IAU Name  &  Star &  RA(2000)   & DEC(2000)     & ~~$B$         &  ~~$R$          \\\\ &       &             &               & (mag)       & (mag)         \\\\ \\hline &       &             &               &             &               \\\\ 0003$+$158&  S1   & 00 06 08.42 & $+$16 09 54.4 & 16.31       & 15.32         \\\\ &  S2   & 00 06 06.20 & $+$16 10 46.3 & 17.17       & 15.71         \\\\ &  S3   & 00 06 05.97 & $+$16 12 15.6 & 16.89       & 15.49         \\\\ 0025$+$307&  S1   & 00 28 25.59 & $+$31 03 19.0 & 15.57       & 14.11         \\\\ &  S2   & 00 28 15.86 & $+$31 03 09.8 & 15.57       & 13.89         \\\\ 0043$+$039&  S1   & 00 45 39.87 & $+$04 10 02.0 & 16.95       & 15.50         \\\\ &  S2   & 00 45 44.87 & $+$04 10 57.9 & 17.34       & 16.09         \\\\ &  S3   & 00 45 44.16 & $+$04 13 26.0 & 17.76       & 15.21         \\\\ 0806$+$315&  S1   & 08 09 06.08 & $+$31 22 19.3 & 16.52       & 15.09         \\\\ &  S2   & 08 09 18.58 & $+$31 22 20.7 & 16.54       & 15.11         \\\\ &  S3   & 08 09 14.89 & $+$31 20 18.7 & 17.65       & 15.92         \\\\ 0824$+$098&  S1   & 08 27 39.18 & $+$09 41 13.5 & 16.60       & 15.03         \\\\ &  S2   & 08 27 44.30 & $+$09 45 05.6 & 16.28       & 15.44         \\\\ 0832$+$251&  S1   & 08 35 26.47 & $+$24 57 12.2 & 18.86       & 16.62         \\\\ &  S3   & 08 35 47.24 & $+$24 57 19.0 & 16.56       & 15.72         \\\\ 0846$+$513&  S1   & 08 50 14.07 & $+$51 06 21.9 & 17.27       & 16.11         \\\\ &  S3   & 08 50 19.88 & $+$51 09 00.0 & 17.34       & 18.67         \\\\ 0850$+$440&  S1   & 08 53 28.75 & $+$43 46 22.8 & 18.37       & 16.08         \\\\ &  S2   & 08 53 48.92 & $+$43 48 28.1 & 18.03       & 16.39         \\\\ &  S3   & 08 53 39.97 & $+$43 46 15.4 & 18.72       & 16.46         \\\\ 0931$+$437&  S1   & 09 34 46.90 & $+$43 32 05.9 & 15.75       & 14.42         \\\\ &  S2   & 09 35 01.19 & $+$43 27 43.4 & 15.72       & 15.24         \\\\ 0935$+$416&  S1   & 09 38 40.37 & $+$41 26 11.3 & 16.11       & 15.32         \\\\ &  S3   & 09 39 02.53 & $+$41 30 37.9 & 16.27       & 15.47         \\\\ 0945$+$438&  S1   & 09 49 28.88 & $+$43 37 54.4 & 15.30       & 14.35         \\\\ &  S3   & 09 49 06.74 & $+$43 29 08.2 & 17.18       & 15.27         \\\\ &  S4   & 09 48 58.30 & $+$43 55 11.8 & 17.28       & 16.14         \\\\ 1029$+$329&  S1   & 10 32 10.68 & $+$32 36 08.1 & 16.37       & 15.02         \\\\ &  S2   & 10 32 07.49 & $+$32 37 28.2 & 17.35       & 15.33         \\\\ 1418$+$546&  S1   & 14 20 02.31 & $+$54 25 25.3 & 16.28       & 15.58         \\\\ &  S2   & 14 19 46.29 & $+$54 26 43.4 & 16.11       & 15.51         \\\\ &  S3   & 14 19 39.75 & $+$54 21 56.1 & 16.74       & 15.44         \\\\ 1422$+$424&  S1   & 14 25 03.56 & $+$42 14 41.8 & 16.27       & 15.67         \\\\ &  S2   & 14 25 09.20 & $+$42 17 21.8 & 18.14       & 16.81         \\\\ &  S3   & 14 25 11.09 & $+$42 17 51.2 & 15.97       & 15.39         \\\\ 1425$+$267&  S1   & 14 27 47.53 & $+$26 35 14.9 & 15.21       & 13.65         \\\\ &  S2   & 14 27 30.07 & $+$26 36 05.5 & 15.59       & 14.01         \\\\ 1444$+$407&  S1   & 14 16 54.96 & $+$40 36 51.9 & 15.68       & 14.10         \\\\ &  S2   & 14 46 55.62 & $+$40 36 16.6 & 17.04       & 14.99         \\\\ 1522$+$101&  S1   & 15 24 03.25 & $+$09 58 15.2 & 16.99       & 15.03         \\\\ &  S2   & 15 24 07.32 & $+$10 01 02.9 & 17.12       & 15.56         \\\\ 1553$+$113&  S1   & 15 55 35.71 & $+$11 09 33.2 & 16.11       & 15.11         \\\\ &  S3   & 15 55 51.81 & $+$11 12 28.7 & 16.55       & 15.45         \\\\ 1631$+$395&  S1   & 16 33 01.57 & $+$39 20 49.4 & 17.20       & 15.90         \\\\ &  S3   & 16 32 54.19 & $+$39 21 19.8 & 17.91       & 16.48         \\\\ 1750$+$507&  S1   & 17 51 07.39 & $+$50 45 03.5 & 20.11       & 19.55         \\\\ &  S2   & 17 51 06.32 & $+$50 44 33.9 & 16.38       & 14.81         \\\\ &  S3   & 17 51 37.59 & $+$50 43 56.5 & 15.70       & 14.80         \\\\ &       &             &               &             &               \\\\ \\hline \\end{tabular}  \\end{table}    The observations under this program typically achieved convincing detectability of INOV at a level of 0.01--0.02 mag and spanned a total of 113 nights (720 hours) between October 1998 and May 2002. This work provided the first positive detection of INOV for RQQs, though modest evidence for such variations had been obtained earlier (e.g.\\ Gopal-Krishna et al.\\ 1995, 2000; Sagar et al.\\ 1996; Jang \\& Miller 1997). It was, moreover, found that except for BLs and high optical polarization CDQs (HP-CDQs), the amplitude of detected INOV is small ($\\le$ 3 per cent) and so is the INOV duty cycle ($\\sim$10--20 per cent), {\\it irrespective of the radio loudness}. Further, for the BLs and HP-CDQs, for which a strong INOV was frequently observed, no correlation was found between the amplitude of INOV and long-term optical variability (St04b).  We argued that these results are consistent with the hypothesis that even radio-quiet QSOs possess relativistic jets emitting optical emission on sub-parsec scales, but that we are observing them at moderately large angles to the jet direction so that any variations are neither amplified in magnitude nor compressed in time-scale as they are in BL Lacs (GK03; St04a).  Wills (1996) also argued that RQQs do indeed possess jets, but that they propagate through denser gas close to the host galaxies' planes and are thus quickly snuffed out. For BLs we found no correlations between apparent brightness levels and INOV properties (St04b).  This is in accord with a recent study which indicates that  microvariability of a blazar may be correlated with the presence of longer-term flux changes, rather than its apparent brightness level (Howard et al.\\ 2004).  In this paper, we present the results of our optical monitoring for another 20 AGN belonging to all of the above mentioned four classes of luminous AGN. We then combine the present data for 3 BLs with similar high quality light-curves for another 9 BLs taken from literature, to arrive at a representative sample of 51 intra-night optical light-curves for BL Lacs.  This sample allows us to make a comparative study of the INOV properties of BL Lacs detected with the {\\it EGRET} instrument on the {\\it Compton Gamma-Ray Observatory} (Hartman et al.\\ 1999), or otherwise found to emit high energy $\\gamma$-rays, and their counterparts that were not detected by {\\it EGRET} (Section 4).  Our conclusions are summarized in Section 5. ",
        "conclusions": "We have presented new observations of intra-night optical monitoring for 20 powerful AGN, including 3 BL Lacs, 5 radio-loud quasars (RLQs) and 12 radio-quiet quasars (RQQs). INOV is detected in all three classes of AGN, consistent with the results reported in our earlier papers (GK03; St04a,b; Sa04). By combining the present data with the observations reported in our earlier papers (GK03; St04a; Sa04), we could assemble a larger AGN sample consisting of 19 RQQs, 9 BL Lacs and 16 RLQs (after excluding the high optical polarization quasar 1216$-$010 from the radio core-dominated RLQs in our sample).  The INOV duty cycles (DCs) derived for this sample are: 63\\% for BL Lacs, 18\\% for RLQs and 22\\% for RQQs. Thus, the INOV duty cycles for both RQQs and RLQs (5 of which are CDQs) are similar, and much smaller than that for BL Lacs. This supports our earlier result that the mere presence of a powerful radio synchrotron jet does not lead to an enhanced INOV (GK03; St04a). The observed similarity in the INOV of RQQs and non-blazar RLQs, both in terms of DC and $\\psi$, further suggests that the RQQs also eject relativistic jets.  Their jets are, however, probably quenched while crossing the innermost micro-arcsecond scale, possibly through heavy inverse Compton losses in the vicinity of the central engine (GK03). A similar conclusion has also been reported recently from radio variability studies of RLQs and RQQs (Barvainis et al.\\ 2004).  Further, we have formed an enlarged sample of BL Lac objects with intra-night monitoring duration $>$ 4 hours, by combining the 3 LCs reported here with 48 taken from the literature (Sect.\\ 4). The duty cycle of INOV for this entire sample of 51 LCs of BL Lacs is found to be 68\\%.  Dividing this sample of 51 LCs by the criterion of detection of $\\gamma-$rays (Table 4), we find that the $\\gamma-$ray detected BLs show somewhat stronger INOV, the formal confidence being 0.962 using K-S test (Fig.\\ 3). It is tempting to speculate about the possible origin of this difference. The synchrotron self-Compton (SSC) model for the origin of $\\gamma-$rays posits that the $\\gamma-$rays are produced by inverse Compton (IC) scattering of the synchrotron photons themselves off the relativistic jet electrons (e.g.\\ Maraschi, Ghisellini \\& Celloti 1992; Bloom \\& Marscher 1996).  The external Compton (EC) models invoke IC scattering of photons originating outside the jet, typically from the accretion disk around the central black hole (e.g.\\ Dermer, Schlickeiser \\& Mastichiadis 1992), or disk photons reprocessed by matter above the disk but outside the jet (e.g. Sikora, Begelman \\& Rees 1994; Blandford \\& Levinson 1995). A variant of the EC model, the ``mirror model'', utilizes jet photons reflected or reprocessed by clouds external to the jet (Ghisellini \\& Madau 1996).  In the IC scenario involving external seed photons for $\\gamma$--ray loud blazars, one expects the emission cone to be particularly sharp (Dermer 1995), raising the likelihood of detecting stronger and more rapid INOV (GK03). Hence, our current preliminary results provide additional support to the EC model. Nonetheless, more extensive, multi-band intra-night monitoring observations of blazars are clearly needed. For the non-blazar AGN, the amplitude of INOV continues to be found to be small ($<$ 3\\%), emphasizing the need for even more sensitive monitoring programs."
    },
    "0410/astro-ph0410010_arXiv.txt": {
        "abstract": "Simultaneous spectral differential imaging is a high contrast technique by which subtraction of simultaneous images reduces noise from atmospheric speckles and optical aberrations. Small non-common wave front errors between channels can seriously degrade its performance. We present a new concept, a multicolor detector assembly (MCDA), which can eliminate this problem. The device consists of an infrared detector and a microlens array onto the flat side of which a checkerboard pattern of narrow-band micro-filters is deposited, each micro-filter coinciding with a microlens. Practical considerations for successful implementation of the technique are mentioned. Numerical simulations predict a noise attenuation of 10$^{-3}$ at 0.5$^{\\prime \\prime}$~for a 10$^5$ seconds integration on a $m_H$~=~5 star of Strehl ratio 0.9 taken with an 8-m telescope. This reaches a contrast of 10$^{-7}$ at an angular distance of 0.5$^{\\prime \\prime}$~from the center of the star image. ",
        "introduction": "Direct imaging and spectroscopy of exoplanets are crucial to constrain their masses, study their atmospheres and, ultimately, find telltale signatures of life. The task is dauntingly difficult. The planet image is hugely fainter than the background from the brilliant stellar image. The most obvious problem is photon noise in the bright star point spread function (PSF). A remedy is coronagraphy \\citep{lyot1933,roddier1997,rouan2000,baudoz2000,abe2001,kasdin2003}, which attenuates the coherent diffraction pattern of the on-axis PSF. At high Strehl ratios the diffraction pattern dominates much of the PSF and coronagraphs can achieve strong attenuation \\citep{kilston2002}. With the relatively modest Strehl ratios achieved with current adaptive optics (AO) systems on large ground based telescopes, the efficiency of coronagraphy is limited by the presence of a rapidly changing halo of speckles generated by atmospheric phase distortion. This increases the local PSF variance by orders of magnitude above that caused by photon noise \\citep{racine1999}. In addition, aberrations in the optical train, unfiltered by the AO system, produce speckles of much longer lifetime, hence much brighter in a long-exposure image, than atmospheric speckles. This noise can be dominant for integrations longer than a few seconds \\citep{marois2003a,marois2004}. Quasi-static speckles also limit the high-contrast performance of the Hubble Space Telescope \\citep{schneider2003}, in the form of the so-called ``breathing'' problem. All this underscores the importance of PSF calibration for high contrast imaging, which is the topic of this paper.  Simultaneous spectral differential imaging \\citep{smith1987,racine1999,marois2000a,sparks2002,biller2004} is such a calibration technique. Images are acquired simultaneously at adjacent wavelengths in a spectral range where the planetary and stellar spectra differ appreciably. Judicious image subtraction removes the stellar image and conserves that of any companion. A useful spectral feature for this purpose is the sharp methane absorption band head at 1.6~$\\mu$m \\citep{rosenthal1996} found in relatively cold atmospheres \\citep{burgasser2002} such as those of cold T type brown dwarfs and Jovian planets.  The TRIDENT camera \\citep{marois2000b,marois2003a,marois2003b,marois2003c,marois2004}, a 3-channel (1.580, 1.625 and 1.680~$\\mu$m, 1\\% bandwidth) differential imager, has been extensively used with the AO system of the 3.6-m Canada-France-Hawaii Telescope \\citep{rigaut1998}. From the data, one can determine the level of noise attenuation achieved. The noise attenuation factor $\\Delta$N/N is defined as the median over an annulus of the pixel-by-pixel ratio of the absolute value of the intensity in the difference image to its absolute value in the original image after subtraction of an azimuthally averaged profile. The noise structure of the PSF at one wavelength is dominated by static aberrations and is well approximated with $\\sim$~130~nm of static phase screen following a power-law distribution P~$\\propto \\nu^{\\alpha}$ with $\\alpha$~=~$-2.7$. This level of aberration is consistent with that measured (110~nm RMS) for the static component of the PUEO AO system \\citep{rigaut1998}. Fig.~\\ref{fig1} illustrates typical noise attenuations obtained with TRIDENT. At subarcsecond separations the noise residuals in a difference image are $\\sim 10$ times larger than the sum of the photon and read noises. Results from numerical simulations indicate that $\\sim$~$\\lambda$/20 of non-common path aberrations can explain TRIDENT limited performance. A factor of $\\sim 4$ improvement is obtained by subtraction of the residuals of a reference star observed shortly after and under the same conditions as the science target. This indicates that differential aberrations between channels contribute a large fraction of the residual noise in a difference of simultaneous images. It also appears that some of the noise after subtraction of the reference star image is caused by slight changes in optical aberrations between the target and the reference star observations.  In this paper, in light of the non-common path limitations, a new differential imager concept is introduced that should improve detection limits by a large factor. Simulations are made to estimate the performances and the expected detection limits with 8-m class telescopes and high order AO systems. ",
        "conclusions": "Experience with the TRIDENT camera has shown that small non-common path aberrations in multiple optical channel instruments seriously limit PSF noise attenuation hence faint companion detection. The MCDA concept described here can overcome this problem. Numerical simulations predict that, with an MCDA, 2.5~M$_{\\rm{Jup}}$ exoplanets around 0.5 Gyr M0 stars should be detectable in 10$^{5}$ seconds with an 8-m telescope and a 4k actuator AO system.  Planet finders based on MCDAs would enable detection and characterization of exoplanets with present large telescopes and high order AO systems. The technique is relatively simple to implement and would complement coronagraph projects, or telescopes like HST, where it could calibrate evolving noise structure. An MCDA combined with a coronagraph would be a powerful tool for detecting ozone in the mid-IR \\citep{desmarais2002}, water in the near-IR, molecular oxygen in the red and, eventually, the chlorophyll ``red edge'' with the Terrestrial Planet Finder. An optimized MCDA-type polarimeter could study reflected star light from exoplanets \\citep{saar2003}, or disks and zodiacal dust. Instead of micro-filters, micro-polarizers would be used to produce simultaneous polarized images with a single optical channel."
    },
    "0410/gr-qc0410057_arXiv.txt": {
        "abstract": "Gravitational wave bursters are sources which emit repeatedly bursts of gravitational waves, and have been recently suggested as potentially interesting candidates for gravitational wave (GW) detectors. Mechanisms that could give rise to a GW burster can be found for instance in highly magnetized neutron stars (the ``magnetars'' which explain the phenomenon of soft gamma repeaters), in accreting neutron stars and in hybrid stars with a quark core. We point out that these sources  have very distinctive experimental signatures. In particular, as already observed in the $\\g$-ray bursts from soft gamma repeaters, the energy spectrum of the events is a power-law, $dN\\sim E^{-\\g}dE$ with $\\g\\simeq 1.6$, and they have a distribution of waiting times (the times between one outburst and the next) significantly different from the distribution of uncorrelated events. We discuss possible detection strategies that could be used to search for these events in existing gravitational wave detectors. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/hep-th0410082_arXiv.txt": {
        "abstract": "It is possible that superstrings, as well as other one-dimensional branes, could have been produced in the early universe and then expanded to cosmic size today.   I discuss the conditions under which this will occur, and the signatures of these strings.  Such cosmic superstrings could be the brightest objects visible in gravitational wave astronomy, and might be distinguishable from gauge theory cosmic strings by their network properties.  Presented at the Mitchell Conference on Superstring Cosmology, GR17, Cosmo 04, and the 2004 APS-DPF Meeting. ",
        "introduction": "Cosmic strings might also arise as gauge theory solitons, and for some time these were a candidate for the source of the gravitational perturbations that produced the galaxies.  It is useful to give a brief summary here to set a context for the superstring discussion, but for a more complete discussion and references see the reviews~\\cite{VilShell,Hindmarsh:1994re}; for a recent overview see ref.~\\cite{Kibblenew}.  In any field theory with a broken $U(1)$ symmetry, there will be classical solutions that are extended in one dimension~\\cite{Abrikosov:1956sx, Nielsen:1973cs}.  These are topological solitons: as one traverses a circle around the string core, the Higgs field winds around the manifold of vacua, which is also a circle.  The broken $U(1)$ can be either a global or a gauge symmetry.  In the global case the Goldstone boson field has a $1/r$ gradient at long distance, giving a logarithmic potential between strings.  In the gauge case this gradient is pure gauge and the physical fields all fall off exponentially.  Most of the discussion of cosmic strings has focused on gauge strings, though one should keep the global case in mind for later reference.  These solutions exist whenever there is a broken $U(1)$ symmetry, and whenever a $U(1)$ symmetry {\\it becomes} broken during the evolution of the universe a network of strings must actually form.  This is the Kibble argument~\\cite{Kibble:1976sj}: the Higgs field starts at zero, and then rolls down to one of the vacua.  By causality, it cannot roll in the same direction everywhere in the universe, as it cannot be correlated on distances greater than the horizon scale (in practice, the correlation length is usually less than this, being the microphysical correlation length of the field theory).  Thus it chooses random directions in different places, and inevitably there will be some trapped winding, so that strings are left over at the end.  This is indeed what simulations show.  A fraction of the string is in the form of finite sized loops, and a fraction is in the form of infinite strings; the latter enter and leave the boundaries of the simulation volume no matter how large this is taken to be.  These populations are cleanly separated because the distribution of lengths of finite loops falls rapidly for long loops, giving a convergent integral for the total amount of string in loops.  Presumably the existence of the infinite strings is implied by the causality argument. They are important because they begin to stretch with the expansion of the universe, while the small loops quickly decay away.  In the subsequent evolution of the network, the important processes in addition to the expansion of the universe are collisions of strings and the emission of gravitational radiation.\\footnote{For reasons of space I am focusing on strings whose only long-distance interactions are gravitational.  Global strings would also emit Goldstone bosons.  There are also superconducting strings~\\cite{Witten:1984eb}, which have strong couplings to gauge fields.  These are perhaps less likely to arise in the superstring case, for reasons that I will explain.}  When two strings collide they can either pass through each other or they can reconnect (intercommute), as in figure~1. \\begin{figure}[t] \\begin{center} \\includegraphics[height=.13\\textheight]{intercommute.eps} \\caption{When two strings of the same type collide, they either reconnect, with probability $P$, or pass through each other, with probability $1-P$.  For classical solitons the process is deterministic, and $P=1$ for the velocities relevant to the string network.} \\end{center} \\end{figure} In an adiabatic collision, they will always reconnect, $P=1$, because reconnection allows the flux in the string core to take an energetically favorable shortcut.  Simulations show that this remains true at the moderately relativistic velocities that are present in string networks~\\cite{Shellard:1987bv, Matzner, Moriarty:1988fx}.  During the radiation and matter dominated eras, the horizon distance grows linearly in $t$ while the comoving distance grows more slowly, as $t^{1/2}$ or $t^{2/3}$ respectively.  Thus as time goes on we see more of the universe, and if no decay processes were operating we would see more strings.  The relevant decay process is the combined effect of reconnection and gravitational radiation.  Reconnection turns long straight strings into long kinked strings and then allows loops of string to break off, and these loops decay by gravitational radiation.  Simulations show that this occurs at the maximum rate allowed by causality, so that the string network always looks the same as viewed at the horizon scale,  with a few dozen strings spanning the horizon volume and a gas of loops of various sizes.  This is known as the {\\it scaling solution}.  The scaling solution is an attractor: if we start with too much string, the higher collision rate will reduce it, while if we start with too little then there will be few collisions until the amount of string per horizon volume approaches the scaling form.  It is possible that the strings have a short-distance structure that does not scale with the horizon size.  Since we are considering strings whose only interactions are gravitational, all of their effects are controlled by the dimensionless product of Newton's constant and the string tension, $G\\mu$.  This is the string tension in Planck units, and it sets the size of the typical metric perturbation produced by a string.  For example, the geometry around a long straight string is conic, with a deficit angle $8\\pi G\\mu$.  The string network produces inhomogeneities proportional to $G\\mu$, and because of the scaling property of the string network these are scale invariant.  A value $G\\mu \\sim 10^{-5.5}$ would give scale-invariant perturbations of the right magnitude to produce the galaxies and CMB fluctuations.\\footnote{I am going to quote values of $G\\mu$ to the nearest half order of magnitude.  This is less precise than most numbers given in the literature, but it is all I will need and it  roughly reflects the uncertainties in the understanding of the string network.  To give more precise numbers would require a more detailed discussion, and so the reader should consult the references.  Incidentally, some of the bounds given here are stronger than those that I was aware of when this talk was given.} This was a viable theory for some time~\\cite{Zeldovich:1980gh, Vilenkin:1981iu}, but it is now excluded. For example, the CMB power spectrum is wrong: the actual spectrum shows a pattern of peaks and dips, whereas the spectrum from strings would be smooth.  There is a simple reason for this.  The fluctuations produced in inflation have a definite phase.  This phase is maintained from the end of inflation until the perturbations go nonlinear, and is imprinted as oscillations of the power spectrum.  Strings, on the other hand, each keep their own time, there is no common phase.  Fitting cosmological constant plus cold dark matter {\\it plus} strings to the CMB power spectrum gives an upper limit $G\\mu< 10^{-6}$~\\cite{Pogosian:2003mz, Pogosian:2004ny}.  Beyond the power spectrum, strings will produce nongaussianities in the CMB.  Recent limits are somewhat stronger than those from the power spectrum, around $G\\mu< 10^{-6.5}$~\\cite{Jeong:2004ut,Lo}.   Another limit on $G\\mu$ comes from pulsar timing.  Because the energy in the strings eventually goes into gravitational waves, strings produce a large stochastic gravitational wave background.  The classic reference on this subject~\\cite{Kaspi:1994hp} quotes a limit on the energy density in stochastic gravitational waves, per logarithmic frequency range, as $\\Omega_{\\rm GW} < 1.2 \\times 10^{-7}$, using $h^2 = 0.5$ for the Hubble parameter.  In the frequency range of interest one can estimate the stochastic background from a network of strings as $\\Omega_{\\rm GW} = 0.04 G\\mu$~\\cite{VilShell}, meaning that $G\\mu < 10^{-5.5}$.  Different analyses of {\\it the same} data, with different statistical methods, give a (controversial) bound a factor of 6 stronger~\\cite{Thorsett:1996dr} and another a factor of 1.6 weaker~\\cite{McHugh:1996hd}.  The limits from pulsars should increase rapidly with greater observation time, and recent work~\\cite{Lommen:2002je} quotes a bound 30 times stronger than~\\cite{Kaspi:1994hp}, meaning that $G\\mu < 10^{-7}$.  However, this paper again obtains much weaker limits using other methods and so this should not be regarded as a bound until there is agreement on the analysis.\\footnote{I would like to thank E. Flanagan and H. Tye for communications on this point.} The bounds here are from gravitational waves with wavelengths comparable to the size of the emitting string loop, and do not include a potentially substantial enhancement due to high-frequency cusps; see section~4.  Thus far we have quoted upper bounds, but there are possible detections of strings via gravitational lensing.  A long string will produce a pair of images symmetric about an axis, very different from lensing by a point mass.  Such an event has been reported recently~\\cite{Sazhin:2003cp, Sazhin:2004fv}.  The separation of around two arc-seconds corresponds to  $G\\mu$ equal to $4 \\times 10^{-7}$ times a geometric factor that is at least 1. This appears to exceed the upper limit from pulsars, unless there exist fewer strings than in the usual network simulations.  I have heard varying opinions on how seriously to take such an observation, as similar pairs in the past have turned out to coincidental.  There is further discussion in the review~\\cite{Kibblenew}, which also discusses a possible time-dependent lens, as from an oscillating loop, with $G\\mu \\sim 10^{-7.5}$~\\cite{Schild:2004uv}.  Network simulations show that the total energy density in the string network is of order $60 G\\mu$ times the matter density (during the matter-dominated era), so the upper bound on $G\\mu$ implies that the strings contribute only a small fraction to the total energy density.  In particular, the strings are not the dark matter, if they have the usual network properties. ",
        "conclusions": "As we have seen, each of the four conditions that we discussed at the beginning is independently model-dependent.\\footnote{I think that I am presenting this in an overly pessimistic way.  When I hear talks on new physics (especially in the context of cosmology!) there is much less of a sense of an a priori measure of what is likely.  I am quite sure that many of the ideas that are prominently discussed have an a priori probability very much less than cosmic superstrings. (I had started to include a list, but decided that it would be distracting and inflammatory.  It was rather long.)} However, quite a number of things have worked out surprisingly well: the production of strings in brane inflation, the possible stabilization of the strings as a side effect of other properties of the models (in particular, of the stabilization of the vacuum itself), the possibility to see strings over many interesting orders of magnitude of tension, and the existence of properties that distinguish different kinds of string so that after the strings are discovered we can do a lot of science with them.  In any case, searching for cosmic strings is a tiny marginal cost on top of experiments that will already be done, and it is great that string theorists will have a stake in these experiments over the coming decade or more."
    },
    "0410/astro-ph0410540_arXiv.txt": {
        "abstract": "$\\eta$ Carinae was observed by FUSE through the LWRS (30\\arcsec x30\\arcsec) and HIRS (1.25\\arcsec x20\\arcsec) apertures in March and April 2004. There are significant differences between the two spectra. About half of the LWRS flux appears to be due to two B-type stars near the edge of the LWRS aperture, 14\\arcsec from $\\eta$ Car. The HIRS spectrum (LiF1 channel) therefore reveals the intrinsic FUV spectrum of $\\eta$ Car without this stellar contamination.  The HIRS spectrum contains strong interstellar H$_2$ having high rotational excitation (up to $J=8$).  Most of the atomic species with prominent ISM features (C II, Fe II, Ar I, P II, etc) also have strong blue-shifted absorption to $v\\sim -580$ km/s that is associated with expanding debris from the 1840 eruption. ",
        "introduction": "The luminous blue variable (LBV) $\\eta$ Carinae (HD 93308), a highly evolved, massive and luminous star, shows strong evidence for a binary companion with a 5.5 year periodicity in X-rays, optical and UV (Daminelli 1996, van Genderen et al. 2003, Whitelock et al. 2004, Corcoran et al. 2001). Most of its luminosity is emitted in the infra-red, but there is flux and variability at all wavelengths from hard X-rays to radio. The observed double-lobed structure consists of ejecta from the massive outburst in 1840. The orbital parameters of $\\eta$ Car are not well constrained, although extensive modelling is on its way (Corcoran et al. 2001; Pittard  \\& Corcoran 2002; Corcoran et al. 2005, in preparation).    \\begin{figure} \\plotone{iping2_fig1.ps} \\caption{The LiF1a spectrum of $\\eta$ Car from 990 \\AA\\ to 1070 \\AA.  The thick line of the HIRS aperture spectrum; the thin line is LWRS.  The HIRS flux was scaled up by 1.67X to correct for the point source throughput (60\\%). The H$_2$ Lyman and Werner ground state transitions are shown, respectively, by short- and long-dashed lines and are identified by the rotational branch and quantum number (e.g. R3) above each line and by the vibrational quantum numbers (e.g. 0-8) on the left-hand side of each panel.  The principal atomic transitions and the R0 lines of HD are also identified.  All transitions are marked at their laboratory wavelengths.} \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410589_arXiv.txt": {
        "abstract": "The X-ray Background has been resolved in the $0.5$--$5\\kev$ band and found to consist mostly of both unabsorbed and absorbed AGN with column densities $<10^{23}\\pcmsq$. This contrasts with the local AGN population where the column density range extends to Compton-thick objects and beyond ($>1.5\\times 10^{24}\\pcmsq$). Stacking analysis of the integrated emission of sources detected by XMM-Newton in the Lockman Hole, and by Chandra in the CDF-N and S reveals that the resolved fraction of the X-ray Background drops above $6\\kev$ and is about $50\\%$ above $8\\kev$. The missing flux has the spectrum of highly absorbed AGN, making it likely that the range of column density at redshift one is similar to that locally, and that many AGN are as yet undetected in well-studied fields. ",
        "introduction": "Deep images with Chandra and XMM-Newton have shown that the X-ray Background (XRB) is composed, at least in the $2$--$5\\kev$ band, of the emission of many active galactic nuclei (AGN) \\cite{giacconi02}\\cite{mushotzky00}\\cite{alexander03}\\cite{barger02}. This was anticipated by earlier work with ROSAT \\cite{hasinger98} which resolved the $0.5$--$2\\kev$ XRB in a similar manner. The XRB therefore reveals the integrated X-ray emission of AGN and is an important tool for studying the accretion history of the Universe and the evolution of AGN.  The background spectrum in the $2$--$10\\kev$ band is flatter than typical AGN and requires that most are absorbed \\cite{setti89}\\cite{madau94}\\cite{comastri95}. This ties in with what we know of AGN locally where absorbed Seyfert II galaxies outnumber Seyfert Is. What remains to be determined is whether powerful objects such as quasars have a significant obscured fraction, how the obscuration evolves with time in an object and how obscuration evolves with redshift within the population. ",
        "conclusions": "Studies of the XRB continue to reveal the evolution and properties of AGN. Most of the XRB is resolved up to an energy of 5~keV and is due to AGN -- Seyferts and quasars -- with absorbing column densities $<2\\times 10^{23}\\pcmsq$. The identification of large samples of such objects shows that the lower luminosity ones (Seyfert-like) evolve rapidly to $z\\sim 0.8$ while quasars evolve more slowly to $z\\sim 2$.  Assuming that the Seyfert-like ones have lower black hole masses, those black holes continued to grow at later stages of the Universe than the more massive ones associated with quasars.  The resolved fraction of the XRB drops to $60$--$70\\%$ above 6~keV and $50\\%$ above 8~keV. The missing fraction between 6--8~keV is not due to the known faint sources in the 2--5~keV band since it would have been detected by the stacking analysis which corresponds to an exposure time of 0.5~Gs or more.  The missing fraction is most simply explained as due to the high column density ($N_{\\rm H}>2\\times 10^{23}\\pcmsq$) extension to the observed Seyfert-like objects evolving rapidly to $z\\sim 1$. Such a population was implicit in XRB synthesis models such as that of \\cite{gandhi03} or \\cite{gilli01}.  Presumably there is also a Compton-thick population as well, which would make less of an impact on the XRB since Compton down-scattering then depletes the emergent X-ray luminosity. The total black hole density due to Compton-thin sources is then about $\\rho_{\\rm h}\\approx 5\\times 10^5 \\eta^{-1}_{-1}\\msunpmpccub$, where $\\eta=0.1\\eta_{-1}$. This is consistent with recent estimates of the local value of $\\rho_{\\rm h}$ \\cite{marconi04}\\cite{shankar04} for $\\eta=0.1$.  The issue of highly obscured and Compton-thick quasars remains unresolved \\cite{comastri04}. Powerful radio galaxies represent one, relatively small, population. Some powerful, highly obscured, radio-quiet quasars have been found \\cite{crawford02}\\cite{norman02}\\cite{stern02}\\cite{wilman03}, with \\cite{gandhi04} estimating a surface density of at least 10~deg$^{-2}$, The uncertainties are currently too large to assess their general importance. Future work combining X-ray and Spitzer data will be instructive."
    },
    "0410/astro-ph0410406_arXiv.txt": {
        "abstract": "{ We use a new technique to extract the spectrum of a supernova from that of the contaminating background of its host galaxy, and apply it to the specific case of high-redshift Type Ia supernova (SN Ia) spectroscopy. The algorithm is based on a two-channel iterative technique employing the Richardson-Lucy restoration method and is implemented in the IRAF code \\emph{specinholucy}. We run the code both on simulated (SN Ia at $z=0.5$ embedded in a bright host galaxy) and observed (SNe Ia at various phases up to $z=0.236$) data taken with VLT+FORS1 and show the advantages of using such a deconvolution technique in comparison with less elaborate methods. This paper is motivated by the need for optimal supernova spectroscopic data reduction in order to make meaningful comparisons between the low and high-redshift SN Ia samples. This may reveal subtle evolutionary and systematic effects that could depend on redshift and bias the cosmological results derived from comparisons of local and high-$z$ SNe Ia in recent years. We describe the various aspects of the extraction in some detail as guidelines for the first-time user and present an optimal observing strategy for successful implementation of this method in future high-$z$ SN Ia spectroscopic follow-up programmes. ",
        "introduction": "The initial claim made by two independent teams -- the High-Z Supernova Search Team \\citep{schmidt98} and the Supernova Cosmology Project \\citep{P99} -- that the apparent dimming of Type Ia supernovae (SN Ia) at redshifts of $z \\approx 0.5$ implies a present accelerating expansion of the universe \\citep{R98,P99}, and its subsequent confirmation with improved precision \\citep{tonry03,knop03,barris04}, have prompted an increased interest in these astrophysical events. On the theoretical side, a lot of effort and computing time has been devoted to multi-dimensional modelling of the explosion in order to provide physical input parameters for spectral synthesis calculations (see \\citealt{wolfi00} for a review). On the observational side several teams are currently using SNe Ia both to probe the decelerating expansion at higher ($z \\gtrsim 1$) redshifts -- the Hubble Higher-$z$ Supernova Search \\citep{riesshz1} -- as well as measure second order effects (the equation-of-state parameter $\\omega=p/\\rho c^2$) at intermediate ($0.2 \\lesssim z \\lesssim 0.8$) redshifts -- such as ESSENCE \\citep{miknaitis05,matheson05} and the CFHT SuperNova Legacy Survey, or SNLS \\citep{pain02}.  However, the physical parameters governing the SN Ia explosion mechanism are today not fully understood and the empirical parametrisation of the entire SN Ia class -- the so-called ``Phillips relation'' \\citep[see][]{phillips93} -- has no physical basis. The aforementioned observational programmes will detect several hundreds of high-$z$ SNe Ia, at which stage the statistical errors in the sample will have reached the systematic error floor. The cosmological effect detected via observations of Type Ia supernovae is indeed a subtle one (see \\citealt{bruno01} for a review) and the field is limited by the systematic errors involved in such measurements and their possible evolution with redshift.  Spectroscopy is an ideal way to probe these potential evolutionary effects through systematic comparison of high-$z$ SN Ia spectra with local templates. This requires the extraction of clean SN Ia spectra devoid of background contamination, mainly by the host galaxy and sky lines (for ground-based observations). For high signal-to-noise (hereafter S/N) cases, when the supernova is bright with respect to its immediate underlying background, standard extraction software will in general work well. However, when the supernova is faint compared to the background, e.g. at late phases, heavily superposed on its host galaxy or simply at a high redshift, a more elaborate technique is required to ensure that we are not extracting non-SN flux.  The purpose of this paper is to test such a method that has been implemented in the IRAF\\footnote{IRAF is distributed by the National Optical Astronomy Observatories, operated by the Association of Universities for Research in Astronomy, Inc., under contract to the National Science Foundation of the United States.} code \\emph{specinholucy} \\citep{LW03}. It is based on a two-channel restoration algorithm that restores a point spread function (PSF)-like component in a 2D image and an underlying extended background \\emph{separately}. It is of wide astronomical use and has already been successfully applied to the restoration of point-source spectra in highly inhomogeneous backgrounds \\citep[see][]{LW03}, such as is often the case for high-$z$ SNe Ia embedded in their host galaxy. The fact that the entire 2D spectrum of the background (with or without the inclusion of the restored point source spectrum) is restored in such an algorithm means that we have a firm hold on potential systematic errors associated with the restoration. This is done by simply comparing the residuals in the restored 2D spectrum with the statistical noise limit of the input 2D spectrum (the square root of the number of photoelectrons). Should these residuals fall below this limit then we can consider the restoration to be successful. This is a clear advantage over traditional spectral extraction routines where no secure hold on systematic errors is possible. Furthermore, the very nature of the algorithm presented in this paper implies that all non PSF-like components in the input 2D spectrum are automatically allocated to the background channel. We are making no \\emph{a priori} assumptions on the nature of the contaminating extended background component and only rely on the point-like nature of the supernova's spatial profile.  In section \\ref{algo} we present the algorithm and its application to the case of supernovae embedded in their host galaxy. We dedicate sections \\ref{ssfissue} and \\ref{spatialres} to the specific issues of the PSF spectrum -- or slit spread function (SSF) -- used to restore the point source and of the spatial resolution kernel used to discriminate between the point source and the underlying extended background, respectively, as they are of crucial importance to the proper functioning of \\emph{specinholucy}. In sections \\ref{testsim} and \\ref{real} we test the algorithm on simulated and observed data, respectively. Section \\ref{comparison} serves the purpose of comparing our technique to alternative spectral extraction methods, and we further discuss the advantages of using our approach alongside an optimal observing strategy for its successful implementation in section \\ref{conclusion}.  The need for comparisons between the high-$z$ and the local SN Ia spectra in order to make quantitative statements on their potential differences requires the use of such an elaborate two-channel deconvolution technique. The code \\emph{specinholucy} enables a truly quantitative analysis of high-$z$ SN Ia spectra and opens the path for meaningful detections of potential evolutionary signatures in their spectra. ",
        "conclusions": "}  We have presented the advantages of using a more elaborate spectral extraction algorithm for obtaining supernova spectra devoid of background contamination, and hence enable a truely quantitative analysis of high-$z$ Type Ia supernova spectra. This will allow us in the future to set constraints on potential evolutionary effects that could affect SNe Ia as function of redshift. Not only are we able to extract the PSF-like components of an input spectrum, one can also restore the entire 2D spectrum and hence get a hold on systematic errors linked with the restoration process. We thus have a quantitative way of judging whether or not the extracted SN spectrum is optimal for subsequent analysis and whether we are not mistakenly ``seeing'' SN-like features in a pure galaxy spectrum, as is sometimes the case in low S/N high-$z$ SN Ia spectra.  However, there are some caveats mainly associated with the complexity of \\emph{specinholucy}. Errors associated with the use of an inappropriate SSF for restoring the point-source channel or a mis-evaluation of the width of the spatial resolution kernel necessary for the point-source/extended background discrimination will strongly affect the restoration (see section \\ref{ssfissue}). Through restoration of the 2D background spectrum one is however able to produce diagnostic plots to critically evaluate the error in each case and repeat the restoration with different settings (see Fig. \\ref{skernel}).  Note that a valuable by-product of this novel technique is the acquisition of a host galaxy spectrum with no contamination from the supernova flux! Galaxies at $z \\sim 0.5$ are at most a few arcseconds in size and one can essentially obtain integrated spectra of high-$z$ SN host galaxies by centering a slit on both the SN and its host. This in turn enables the study of properties of high-$z$ SN Ia host galaxies and compare them with integrated spectra of local SN Ia hosts \\citep{sullivan03,williams03,gallagher04}.  Thus, \\emph{specinholucy} provides the high-$z$ supernova community with a reliable tool to obtain clean SN Ia spectra and in turn to derive accurate physical quantities associated with the ejecta. Two of the authors (St\\'ephane Blondin and Bruno Leibundgut) are part of the ESSENCE project, and Gr\\'egory Sainton is part of the SNLS collaboration. Both teams recognise the need for a quantitative analysis of their results in order to constrain potential systematic effects that could distort the cosmological results derived from observations of high-$z$ SNe Ia. We have shown \\emph{specinholucy} to be the ideal tool for this purpose. In fact, the VLT data taken for the ESSENCE project will be shown in their restored form via \\emph{specinholucy} in \\citealt{matheson05}."
    },
    "0410/nucl-th0410074_arXiv.txt": {
        "abstract": "{ Core collapse supernovae are gigantic explosions of massive stars that radiate 99\\% of their energy in neutrinos.  This provides a unique opportunity for large scale parity or charge conjugation violation.  Parity violation in a strong magnetic field could lead to an asymmetry in the neutrino radiation and recoil of the newly formed neutron star.  Charge conjugation violation in the neutrino-nucleon interaction reduces the ratio of neutrons to protons in the neutrino driven wind above the neutron star.  This is a problem for r-process nucleosynthesis in this wind.  On earth, parity violation is an excellent probe of neutrons because the weak charge of a neutron is much larger than that of a proton.  The Parity Radius Experiment (PREX) at Jefferson Laboratory aims to precisely measure the neutron radius of $^{208}$Pb with parity violating elastic electron scattering.  This has many implications for astrophysics, including the structure of neutron stars, and for atomic parity nonconservation experiments.}  \\PACS{ {25.30.Bf}, {11.30.Er}, {26.60+c}, and {97.60.Bw}} % ",
        "introduction": "\\label{intro} Can there be large scale parity violation in astrophysics?  For example, the 1960s science fiction movie {\\it Journey to the Far Side of the Sun} visits a parity double of earth.  In the real world, core collapse supernovae provide a unique opportunity for macroscopic parity violation.  These gigantic explosions of massive stars are dominated by neutrinos that transport 99\\% of the energy.  This is because no other known particle can diffuse through the very dense matter of the collapsed star.  In Section \\ref{supernovae} we discuss how the weak interactions of these neutrinos may lead to large scale parity or charge conjugation violation.  On earth, parity violation can be used to probe neutrons because the weak charge of a neutron is much larger than that of a proton.  In Section \\ref{PREX} we describe the Parity Radius Experiment (PREX) at Jefferson Laboratory that aims to use parity violating electron scattering to measure the neutron radius in $^{208}$Pb.  This measurement has many implications for astrophysics including properties of neutron stars, nuclear structure, and atomic parity nonconservation. ",
        "conclusions": "\\label{conclusions}  Core collapse supernovae are dominated by neutrinos.  This provides a unique opportunity for large scale parity or charge conjugation violation.  Parity violation in a strong magnetic field could lead to an asymmetry in the neutrino radiation and recoil of the newly formed neutron star.  Charge conjugation violation in the neutrino driven wind above a neutron star reduces the ratio of neutrons to protons.  This is a large problem for r-process nucleosynthesis in the wind.  On earth, parity violation probes neutrons because the weak charge of a neutron is much larger than that of a proton.  The Parity Radius Experiment (PREX) at Jefferson Laboratory aims to measure the neutron radius of $^{208}$Pb to 1\\% by using parity violating elastic electron scattering.  Because it is an electroweak reaction, the theoretical interpretation of this measurement is very clean.  The neutron radius of Pb has many implications for atomic parity nonconservation experiments, neutron stars, and the structure of neutron rich nuclei."
    },
    "0410/astro-ph0410177_arXiv.txt": {
        "abstract": "We present the results of a detailed analysis of multi-wavelength observations of a very impulsive solar flare 1B/M6.7, which occurred on 10 March, 2001 in NOAA AR 9368 (N27 W42). The observations show that the flare is very impulsive with very hard spectrum in HXR that reveal non-thermal emission was most dominant. On the other hand this flare also produced type II radio burst and coronal mass ejections (CME), which are not general characteristics for impulsive flares. In H$\\alpha$ we observed the bright mass ejecta (BME) followed by drak mass ejecta (DME). Based on the consistence of the onset times and direction of BME and CME, we conclude that these two phenomena are closely associated. It is inferred that the energy build-up took place due to photospheric reconnection between emerging positive parasitic polarity and predominant negative polarity, which resulted as a consequence of flux cancellation. The shear increased to $>$80$^o$ due to further emergence of positive parasitic polarity causing strongly enhanced cancellation of flux. It appears that such enhanced magnetic flux cancellation in a strongly sheared region triggered the impulsive flare. ",
        "introduction": "Impulsive flares produce a wide range of emission from $\\gamma$-rays, X-ray, EUV, visible, microwave (MW) and longwave radio emission (Miller et al., 1997). The X-ray emission is comprised of thermal, superhot and non-thermal components. Based on the classical nonthermal thick-target model it is proposed that the electrons accelerated to $>$ 100 keV in the corona spiral downwards, and creating MW (Brown, 1971; Lin and Hudson, 1976; Kundu and White, 2001). When they reach the footpoints of a coronal loop, they produce hard X-ray (HXR) and drive evaporation, which fills the loop with hot plasma emitting in soft X-rays (SXR).  For better understanding of the various aspect of flare phenomena at different height of solar atmosphere, it is important to resolve the  different structures involved and temporal evolution during the flare. Therefore simultaneous  observations of chromospheric (optical) coronal (X-ray and MW) with high temporal and spatial resolutions can provide wealth of information for flare diagnostic.  In this paper we present analysis of 10 March 2001 impulsive flare at different wavelengths, viz., H$\\alpha$, HXR, SXR and radio wave to understand the various physical processes taking place at different height in the above flare.  \\vspace*{-0.6cm} ",
        "conclusions": "NOAA AR 9368 appeared on 02 March 2001 as most simple $\\alpha$ type magnetic class region, which slowly became magnetically complex as a function of time. Several MMF of positive polarity were observed near following spots group  of south dominating polarity and causing cancellation of magnetic field. It appears that magnetic field cancellation at following spot side as well as increasing shear was the basic physical process for energy build-up. The high shear, and hence the enhanced magnetic stress erupted the flare, which was associated with the CME through a shock produced by it. The fast passage of CME induced the reconnection, which in turn was responsible for the impulsive energy release as observed in the flare and associated HXR and MW emission. It seems that compact and confined impulsive flare associated with CME require continual appearance of opposite polarity emerging fluxes (MMF) and increasing shear for impulsive energy release."
    },
    "0410/gr-qc0410130_arXiv.txt": {
        "abstract": "\\noindent We present a brief review of brane-world models --- models in which our observable Universe with its standard matter fields is assumed as localized on a domain wall (three-brane) in a higher dimensional surrounding (bulk) spacetime. Models of this type arise naturally in M-theory and have been intensively studied during the last years. We pay particular attention to the covariant projection approach, the Cardassian scenario, to induced gravity models, self-tuning models and the Ekpyrotic scenario. A brief  discussion is given of their basic properties and their connection with conventional FRW cosmology. \\\\  \\noindent {\\bf PACS numbers:} 04.50.+h, 11.25.Mj, 98.80.Jk ",
        "introduction": "}   In the brane-world picture\\index{brane world}, it is assumed that ordinary matter (consisting of the fields and particles of the Standard Model) is trapped to a three-dimensional submanifold  (a three-brane) in a higher dimensional (bulk) spacetime\\index{bulk spacetime}. This world-brane is identified with the currently observable Universe. Such kind of scenarios are inspired by string theory/M-theory. One of the first examples of this type was the Ho$\\check{\\mbox{r}}$ava-Witten setup \\cite{HW1} of $11-$dimensional M-theory compactified on an $S_1/Z_2$ orbifold\\index{orbifold} (i.e. a circle folded on itself across a diameter). In this setup, $E_8$ gauge fields are confined to two $(1+9)-$dimensional planes/branes located at the two fixed points of the $S_1/Z_2$ orbifold. In the low-energy limit, the $E_8$ symmetry is down-broken to the $SU(3)\\times SU(2)\\times U(1)$ group of the Standard Model. Simultaneously, the $9-$brane is endowed with the structure of a product manifold $M_3\\times M_6$ consisting of a noncompact real 3D manifold $M_3$ and a compact real 6D manifold $M_6$ with the complex structure of a Calabi-Yau threefold $M_6\\sim CY_3$ and  a characteristic compactification scale $r_{CY_3}$. When the distance between the planes is much larger than the compactification scale, $R >> r_{CY_3}$,  then the $CY_3$ degrees of freedom can be integrated out (see e.g. \\cite{Lukas1}) and one arrives at an effective 5-D heterotic M-theory with two $(1+3)-$dimensional planes. \\begin{figure}[htb] \\centerline{\\hbox{\\includegraphics[width=0.5\\textwidth]{fig1.eps}}} \\caption{Schematic of the dimensional reduction of a heterotic $S^1/\\ZZ_2$ orbifold setup.} \\end{figure} One of these planes describes our brane/Universe and the other one corresponds to a hidden brane (see Fig. 1). In the formal limit $R\\to \\infty$, one obtains a model with a single brane.  The described scenario is one of the possible scenarios for an embedding of branes into a higher-dimensional bulk manifold. Because up to now only certain aspects of the underlying fundamental higher dimensional theory are understood, the model building process is still going on and an uncovering of new aspects of brane nucleation/embedding mechanisms can be expected also for the future.  The presently known brane world models can be distinguished by the number of branes involved, their co-dimensions, as well as the topological structure of the branes and the higher dimensional bulk spacetime. One of the basic tests that any brane-world model has to pass, is the test on its compatibility with observable cosmological data and their implications. From the large number of currently present scenarios we select in the present mini-review only a few ones --- to illustrate some of the possible deviations from usual FRW cosmology. The aim of the brief discussion is twofold. Firstly, for each of the selected models we briefly review the modifications of the Friedmann-like cosmological equations and demonstrate the regimes where the corresponding physics coincides with conventional FRW cosmology; and, secondly, we briefly indicate the type of observable deviations from it. ",
        "conclusions": "}    We are living in very interesting time because new observational data directly indicate that our knowledge about the Universe and its evolution is in great extent limited. In this situation new ideas and new theoretical models are required to explain the observational data. In their turn, these new models will predict new observable phenomena which can be used as test tools to single out those scenarios which are most close to nature.  One possible direction to modify the currently accepted standard theory of cosmology are brane-world models --- M-theory inspired setups in which our observable Universe is interpreted as a 3D submanifold embedded in a higher dimensional bulk spacetime. The aim of the present mini-review was to give a very brief description of some of the basic features of these models --- illustrating them with the help of a few concrete toy model setups.    \\vspace*{1ex}  \\mbox{} \\\\ {\\bf Acknowledgments}\\\\ We thank the High Energy, Cosmology and Astroparticle Physics Section of the ICTP (Trieste) for their warm hospitality during the preparation of this mini-review as well as for financial support from the Associate Scheme (A.Z.) and from HECAP (U.G.). Additionally, U.G. acknowledges support from DFG grant KON/1806/2004/GU/522."
    },
    "0410/astro-ph0410341_arXiv.txt": {
        "abstract": "We present a catalog of 450 high signal-to-noise microlensing events observed by the MACHO collaboration between 1993 and 1999. The events are distributed throughout our fields and, as expected, they show clear concentration toward the Galactic center. No optical depth is given for this sample since no blending efficiency calculation has been performed, and we find evidence for substantial blending. In a companion paper we give optical depths for the sub-sample of events on clump giant source stars, where blending is not a significant effect.  Several events with sources that may belong to the Sagittarius dwarf galaxy are identified. For these events even relatively low dispersion spectra could suffice to classify these events as either consistent with Sagittarius membership or as non-Sagittarius sources. Several unusual events, such as microlensing of periodic variable source stars, binary lens events, and an event showing extended source effects are identified.  We also identify a number of contaminating background events as cataclysmic variable stars. ",
        "introduction": "The structure and composition of our Galaxy is one of the outstanding problems in contemporary astrophysics.  While inventories of bright stars have been made, it is known that the bulk of the material in our Galaxy is dark.  In addition, the number and mass distribution of stellar remnants such as white dwarfs, neutron stars and black holes is quite uncertain, as is the number of faint stars, brown dwarfs and extra-solar planets.  Gravitational microlensing was suggested as a probe to detect compact objects including dark matter (\\pac\\ 1986, Griest 1991a, Nemiroff 1991) and was observationally discovered in 1993 (Alcock et al.\\ 1993; Aubourg et al.\\ 1993; Udalski et al.\\ 1993).  The line-of-sight towards the Large Magellanic Cloud (LMC) is best for dark matter detection, but as a probe of faint objects from planets to black holes, the line-of-sight towards the Galactic bulge is superior (Griest et al.\\ 1991b; \\pac\\ 1991). The high density of stars in the disk and bulge means that the vast majority of events detected by microlensing experiments are in this direction.   The amount of lensing matter between a source and the observer is typically described using the optical depth to microlensing, $\\tau$, defined as the probability that a given source star will be magnified by any lens by more than a factor of 1.34. Early predictions (Griest et al.\\ 1991b; \\pac\\ 1991) of the optical depth towards the Galactic center included only disk lenses and found values near $\\tau = 0.5\\ten{-6}$. The early detection rate (Udalski et al.\\ 1993, 1994a) seemed higher than this, and further calculations (Kiraga \\& \\pac~1994) added bulge stars to bring the prediction up to $0.85 \\ten{-6}$.  The first measurements were substantially higher than this, $\\tau \\geq 3.3 \\pm 1.2 \\ten{-6}$ (Udalski et al.\\ 1994b) based upon 9 events and $\\tau = 3.9^{+1.8}_{-1.2}$ (Alcock et al.\\ 1997) based upon an efficiency calculation and 13 clump-giant events taken from their 45 candidates. Many additional calculations ensued, including additional effects, especially non-axisymmetric components such as a bar (e.g. Zhao, Spergel \\& Rich 1995; Metcalf 1995; Zhao \\& Mao 1996; Bissantz et al.\\ 1997; Gyuk 1999; Nair \\& Miralda-Escud\\'e 1999; Binney, Bissantz \\& Gerhard 2000; Sevenster \\& Kalnajs 2001; Evans \\& Belokurov 2002; Han \\& Gould 2003). Values in the range $0.8 \\ten{-6}$ to $2 \\ten{-6}$ were predicted for various models, and values as large $\\tau=4\\ten{-6}$ were found to be inconsistent with almost any model.  More recent measurements have all used efficiency calculations and have found values of $\\tau = 3.2 \\pm 0.5 \\ten{-6}$ from 99 events in 8 fields using difference imaging (Alcock et al.\\ 2000b), $\\tau = 2.0 \\pm 0.4 \\ten{-6}$ from around 50 clump-giant events in a preliminary version of the companion paper (Popowski et al.\\ 2001a), $\\tau = 0.94 \\pm 0.29 \\ten{-6}$ from 16 clump-giant events (Afonso et al.\\ 2003), and $\\tau = 3.36^{+1.11}_{-0.81}$ from 28 events (Sumi et al.\\ 2003).  In releases similar in spirit to our catalog here Udalski et al.\\ (2000) presented a catalog of 214 microlensing events from the 3 seasons of the OGLE-II bulge observation, and Wo\\'zniak et al.\\ (2001) presented a catalog of 520 events, mainly from difference imaging.  In this work we present our complete catalog of high signal-to-noise microlensing events that were found with point spread function fitting photometry. In the companion paper (Popowski et al.\\ 2004), we make an accurate determination of the bulge optical depth using the 62 clump giant events (60 unique) listed here and find $\\tau = 2.17^{+0.47}_{-0.38} \\ten{-6}$ at $(l,b) = (1 \\hbox{$.\\!\\!^\\circ$} 50, -2 \\hbox{$.\\!\\!^\\circ$} 68)$. We do not calculate an optical depth for our entire sample of microlensing events since a complete blending efficiency calculation has not been performed, and we caution against using the entire sample of events for such purposes.  Initially envisioned as a probe of dark matter, microlensing has evolved into a more general astronomical tool, useful for several distinct purposes. For example, since the duration of the microlensing event is proportional to the square root of the lens mass, microlensing is sensitive to compact objects in the $10^{-7}\\msun$ to $10^1\\msun$ range, independent of the object's luminosity, so it facilitates inventories of brown dwarfs, white dwarfs, and black holes.  However, the lens mass measurement is degenerate with the lens distance and speed, severely limiting the accuracy of the mass distribution measurement.  Our catalog includes several long duration events that may be massive black holes and several short duration events that may be brown dwarfs.  In order to break the mass/distance/speed degeneracy several techniques have been applied to rare classes of events such as those with binary lenses, binary sources, large annual parallaxes, etc.  Our catalog lists events which may be members of exotic microlensing classes. Finally microlensing has emerged as a powerful method of detecting or constraining the existence of extra-solar planets orbiting the lens (Mao \\& \\pac~1991; Gould \\& Loeb 1992; Griest \\& Safizadeh 1998; Rhie et al.\\ 2000; Gaudi et al.\\ 2002). A key for these searches is careful follow-up on microlensing events, almost all of which are towards the Galactic bulge.  Our catalog can be used to determine the frequencies of detectable lensing in various directions towards the bulge.  For comparison with other works we note that our definition of microlensing event duration, $\\that$, is the Einstein ring diameter crossing time, twice the more commonly used Einstein ring radius crossing time. ",
        "conclusions": "In conclusion, light curves and parameters of 528 microlensing events found by the Macho Project (1993-1999) are presented.  Included are 5 events on variable stars, 17 binary events, 24 potential binary events, and 1 extended source event.  Also included is a representative sample of 36 contaminant events, consisting of 16 cataclysmic variables, and 20 duplicate events.  In addition we select 37 (34 unique) events that are potentially lensed Sagittarius sources.  The sample of over 500 events presented here is effected significantly by blending and should not be used for quantitative studies.  We present light curves for all 564 microlensing and non-microlensing events.  Data and figures will be  available at {\\tt http://wwwmacho.mcmaster.ca} upon acceptance of this paper."
    },
    "0410/astro-ph0410388_arXiv.txt": {
        "abstract": "When the gravitational lensing potential can be approximated by that of a circularly symmetric system affected by weak perturbations, it is found that the shape of the resulting (tangential) caustics is entirely specified by the local azimuthal behaviour of the affecting perturbations. This provides a common mathematical groundwork for understanding problems such as the close-wide ($d\\leftrightarrow d^{-1}$) separation degeneracy of binary lens microlensing lightcurves and the shear-ellipticity degeneracy of quadruple image lens modelling. ",
        "introduction": "The recent announcement by \\citet{Bo04} of the most convincing planetary microlensing event \\citep{MP91} to date is one of the greatest success stories of the nearly two-decade old promise of microlensing \\citep{Pa86}. However, despite such successes, the study of microlensing is still hampered by various hurdles, one of the most persistent being the problem of degenerate lightcurves \\citep[e.g.,][]{Do99a,Af00} -- that is, several distinct physical systems can be used to model observed lightcurves.  For the purpose of formal discussion, the degeneracy of microlensing lightcurves may be categorized into two separate but related phenomena. The first, which may be termed as `external' degeneracy, is caused by imperfect observations yielding highly correlated measurements of the parameters that control the observed lightcurves. For instance, in most highly blended events, without extremely good coverage of certain wing portions of the lightcurve, one only expects to measure the ratio of the time-scale to the peak magnification (or the product of the time-scale and the source flux) but not the parameters individually \\citep[c.f.,][]{Go96}. On the other hand, the second type of degeneracy, which will henceforth be referred to as `internal' degeneracy, is due to the fact that the number of independent lightcurve-controlling parameters is smaller than the number of parameters that are used to describe the underlying lens systems. In other words, even with perfectly specified lightcurves (consequently, perfectly determined lightcurve parameters), the underlying lens system may not be uniquely recovered. This internal degeneracy of microlensing lightcurves can be further divided into two broad classes; the `intrinsic' degeneracy among magnification structures from different systems, and different lens systems tracing identical or intrinsically degenerate paths in the given magnification structure (hereafter the `extrinsic' degeneracy). The most widely acknowledged degeneracy of microlensing lightcurves, that is, the `time-scale degeneracy' relating to the mass of the lens, the distances to the source and the lens, and the relative motions between the observer, the source, and the lens, is an example of the extrinsic degeneracy. Other examples of the extrinsic degeneracy are the discrete parallax degeneracies \\citep*[e.g.,][]{SMP03,Go04}. However, it is the intrinsic degeneracy that is of the most theoretical interest as it can be studied through the analysis of the lens equation as a whole without specified reference to the path of the system through the magnification structure or the observational conditions. Furthermore, understanding of the intrinsic degeneracy can greatly facilitate the identification of various extrinsic and external degeneracies. The most obvious example of an intrinsic degeneracy is the azimuthal symmetry as well as the radial similarity of the lensing magnification of the point mass lens. The well-known result of the `inverse-square-root' behaviour of magnification `inside' fold caustics \\citep[e.g.,][]{GP02} is a less obvious example of a \\emph{local} intrinsic degeneracy.  It is notable that the intrinsic degeneracy is intimately related to the symmetry of the lens equation. In general, if (the potential of) the system possesses a certain symmetry, one can expect that there should exist an intrinsic degeneracy of magnification related to it. Note that the magnification is basically the second derivative of the potential. This also means that the study of the intrinsic degeneracy beyond the incidental case studies can be greatly systematized by recognizing the associated symmetry structure of the lens system. However, beyond the obvious example of the circularly symmetric lens, studies of degeneracies rooted in symmetries have been minimal, at best. This paper is one of the first attempts to understand certain intrinsic degeneracies of microlensing by searching for a symmetry in the lens system. It is found that this approach proves itself by providing a unified scheme of understanding the problem of a certain well-known microlensing degeneracy and a seemingly remote problem of lens modelling of a quadruple images.  The primary focus of this manuscript is a point mass lens (and a more general circularly symmetric lens in later sections) that is under the external influence of certain perturbations. It should be noted not only that the symmetry in such systems can be analysed with relative ease, but also that they can be applied to various realistic systems. In Section~\\ref{sec_scl}, the basic properties of the point mass lens are recapitulated, and in Section~\\ref{sec_per}, a perturbative approach to gravitational lensing is developed. In particular, in section~\\ref{ssec_ccc}, the perturbative approach is used to derive an approximate expression for caustics and critical curves. Following this, in Section~\\ref{sec_isp}, it is shown that there exists a certain set of perturbations that yields invariant caustics, which may be seen as the main finding of this monograph. This finding is applied to interpret one of the well-known microlensing degeneracy problems in the next Section~\\ref{sec_egs}. The restrictions imposed on the studied lens system are subsequently relaxed in Section~\\ref{sec_dis}. For example, the discussion is extended from a point mass lens under perturbations to a general circularly symmetric lens system under perturbations in Section~\\ref{ssec_csl}, and from external perturbations to perturbations associated with the system itself in Section~\\ref{ssec_scp}. These two generalizations lay the basis for a study of the shear and ellipticity in lens modelling later in the same section.  In this paper, the lens equation formulated in terms of complex numbers \\citep{BK75,Wi90,WM95} is extensively used. In Appendix~\\ref{asec_caf}, most of the mathematical terminologies regarding complex analysis found in this manuscript are summarized, and in Appendix~\\ref{asec_ceq}, the basic gravitational lens theory is redeveloped using complex-number notation. The treatments given there are minimal and deliberately casual. More in-depth treatments of complex analysis including rigorous definitions and proofs can be found in any standard text of complex analysis \\citep[e.g.,][]{Ah79} or of mathematical methods \\citep[e.g.,][]{AW00}. Interested readers may also consult references regarding potential theory in a plane. ",
        "conclusions": "\\label{sec_dis}  \\subsection{What about the critical curves?} \\label{ssec_wcc}  \\begin{figure} \\plot{ca.eps} \\caption{ Geometrical determination of the azimuthal positions and magnifications of images with a linear approximation of the caustics. For this example, the two cusps lying along the horizontal direction are positive so that the tangent to a caustic point on the bottom half points upwards and vice versa. For a given shape the caustics, if one can draw a tangent to the caustics through a given source position (represented by a dot), there exists an image towards the direction (represented by an arrow) parallel to that tangent from the image plane centre. Note that the image plane centre is not necessarily coincident with the source plane centre nor with the caustics centre. The magnification of the image is proportional to the inverse of the distance between the source and the tangent point. }\\label{ims} \\end{figure}  So far, the discussion on the LCIP pair has been focused on the identical (up to linear approximation) caustics and the degenerate (total) magnification. A natural question to follow is whether they are indistinguishable (at least up to the linear order) in every sense. The answer is actually negative as one can easily discover the difference in the critical curves and the individual image positions. For the monomial perturbation of $f(z)=az^n$, one finds, from equation~(\\ref{lcr}), the corresponding critical curves to be; \\begin{equation} \\label{pcr} z_c(\\phi)=\\E^{\\I\\phi}\\left\\{ 1-\\frac{\\epsilon}{2}n|a|\\cos[(n+1)\\phi+\\phi_a]\\right\\}, \\end{equation} where $\\phi_a$ is the argument of $a$, that is, $a=|a|\\E^{\\I\\phi_a}$. It is straightforward to show that equation~(\\ref{pcr}) is \\emph{not} invariant under the LCIP transformation of $n\\leftrightarrow-(n+2)$ and $a\\leftrightarrow-\\ab$ (equivalently, $\\phi_a\\leftrightarrow \\upi-\\phi_a$), but rather the amplitude of the cosine linear deviation term changes to $\\epsilon(n+2)|a|/2$. In addition, while the azimuthal positions of the images for the LCIP pair are the same since equation~(\\ref{peq}) is identical for them, one finds that the radial positional deviations given by \\begin{equation} \\label{rad} \\delta r=\\frac{\\E^{-\\I\\phi}}{2}\\left[ \\zeta+\\epsilon\\f\\!(\\E^{-\\I\\phi})\\right]\\ (\\in\\mathbf R) \\end{equation} cannot be the same for the two different perturbations. However, it should be noted that the leading order descriptions of the critical curves and the image radial positions are dominated not by the first order but by the zeroth order terms, and they are still identical up to their leading zeroth order. More importantly, despite this difference, the actual observational distinction between the LCIP pair would be rather difficult since the critical curve is physically unobservable and two dominant observable characteristics of the individual images -- the magnification and the azimuthal position -- are indeed identical up to the linear order of the perturbations.  To find the magnification and the azimuthal position of the individual images, one must solve equation~(\\ref{peq}), which is a transcendental equation for $\\phi$, and thus it may not be possible to find algebraic solutions although it is straightforward to find its roots graphically. Incidentally, it is also possible to find them in a purely geometrical manner provided that one can draw its caustics (eq.~\\ref{cau}). From equation~(\\ref{dzp}), one finds that $(d\\zeta_c/d\\phi)\\parallel \\E^{\\I\\phi}$, that is, the parameter $\\phi$ also determines the tangent direction to the caustics. This, together with the fact that equation~(\\ref{ljc}) is real when $\\phi$ is the solution of equation~(\\ref{peq}), also implies that, if one can draw a tangent to the caustic through the given source position, the angle of that tangent line $\\phi$ is a root of equation~(\\ref{peq}). Hence, there exists an image with its azimuthal position with respect to the image plane centre being $\\phi$ and magnification given by $(2l)^{-1}$ where $l$ is the distance between the source and the tangent point. Examples of this procedure for the perturbation given by $f(z)=az$ or its LCIP $f(z)=-\\ab z^{-2}$ is found in Fig.~\\ref{ims}. For definiteness, $a$ is assumed to be real (horizontal direction) in this example, and therefore, the two cusps lying along the horizontal direction are positive. Finally, note that not all images are actually solutions of equation~(\\ref{peq}). However, missing images, if any, are typically higher order solutions that do not significantly contribute to the total magnification provided that the linear approximation is valid and dominant and that the source lies sufficiently close to the centre (and therefore the caustics).  \\subsection{What is so special about the point-mass lens?} \\label{ssec_csl} The discussion until now has been limited to the case for which the unperturbed lens system is described by the point-mass lens equation~(\\ref{le0}). However, the perturbation approach outlined in Section~\\ref{sec_per} can be applied to any lens system whose solutions are known. In fact, the procedure in Section~\\ref{sec_per} can be generalized to the case for which the unperturbed lens system is given by any circularly symmetric lens system that can form an Einstein ring. Suppose that the system is described by the lens equation~(\\ref{ple}) with the unperturbed part given by \\begin{equation} \\zeta_s=z\\left(1-\\frac{1}{r}\\frac{d\\psi}{dr}\\right) =z\\left[1-\\frac{4G}{D_\\mathrm{rel}c^2}\\frac{M(<r)}{r^2}\\right], \\end{equation} where $r=|z|$ is the radial coordinate of image, $\\psi(r)$ is a circularly symmetric lensing potential, and $M(<r)$ is the mass enclosed within a given angular radius $r$. If the equation \\begin{equation} \\label{ein} \\frac{4G}{D_\\mathrm{rel}c^2}M(<r)=r^2 \\end{equation} possesses a positive root $r=r_0$, one can find the linear approximation for the tangential caustics of the system by series-expanding the lens equation at $z_0=r_0\\E^{\\I\\phi}$ and following procedure similar to Section~\\ref{ssec_ccc}. In particular, one finds \\begin{equation} \\zeta=(2-\\beta_0)\\,\\delta r \\E^{\\I\\phi}-\\epsilon\\f\\!(r_0\\E^{-\\I\\phi}), \\end{equation} \\begin{equation} F\\!(\\phi)=\\zeta+\\epsilon\\f\\!(r_0\\E^{-\\I\\phi})- \\E^{2\\I\\phi}[\\bar\\zeta+\\epsilon f\\!(r_0\\E^{\\I\\phi})]=0, \\end{equation} \\begin{equation} \\J=(2-\\beta_0)\\,\\frac{\\zeta-\\zeta_c(\\phi)}{z_0}, \\end{equation} \\begin{equation} \\frac{\\zeta_c(\\phi)}{r_0}=-\\frac{\\epsilon}{2} \\left[\\E^{3\\I\\phi}f^\\prime\\!(r_0\\E^{\\I\\phi})+ \\E^{-\\I\\phi}\\f^\\prime\\!(r_0\\E^{-\\I\\phi})+ 2\\frac{\\f\\!(r_0\\E^{-\\I\\phi})}{r_0}\\right], \\end{equation} and \\begin{equation} \\frac{z_c(\\phi)}{r_0}= \\E^{\\I\\phi}-\\frac{\\epsilon}{2(2-\\beta_0)} \\left[\\E^{3\\I\\phi}f^\\prime\\!(r_0\\E^{\\I\\phi})+ \\E^{-\\I\\phi}\\f^\\prime\\!(r_0\\E^{-\\I\\phi})\\right], \\end{equation} in place of equations~(\\ref{lle}), (\\ref{peq}), (\\ref{ljc}), (\\ref{cau}), and (\\ref{lcr}). Here, $\\beta_0=(d\\ln M/d\\ln r)|_{r=r_0}$ is the local power index for the mass $M$ at the projected angular radius $r=r_0$. In other words, the shape of (the linear approximation of) the caustics is completely determined by the perturbation alone once they are scaled by the angular Einstein ring radius (note that eq.~\\ref{ein} actually defines $r_0$ to be the angular Einstein ring radius for the given circularly symmetric system), and the unperturbed mass distribution only affects the image magnification as a common multiplicative factor, and therefore, the magnification ratios between images are also completely independent of the unperturbed mass distribution. Moreover, it is easy to establish the same LCIP relation $f(z)=ar_0(z/r_0)^n$ and $f(z)=-\\ab r_0(z/r_0)^{-(n+2)}$, or more generally, that the pair of perturbation potentials with an inverse symmetry with respect to Einstein ring radius $r_0$ [i.e., $\\delta\\psi^t(\\bmath r)$ and $\\delta\\psi^p(\\bmath r)= \\delta\\psi^t(r_0^2\\bmath r/r^2)$] has the same expression for the linear approximation of the tangential caustics.  \\subsection{Are there LCIP pairs for convergent perturbations?} \\label{ssec_scp} As has been noted, the LCIP relation originates from a certain intrinsic symmetry of harmonic functions, that is, their azimuthal structures are invariant under the radial distance inversion. This implies that the two-fold correspondence LCIP does not exist for convergent perturbations, for which the perturbation potential is no longer harmonic. However, examination of equations~(\\ref{pty}) and (\\ref{pmp}) reveals more basic properties of the LCIP, that is, the azimuthal Fourier coefficients of the perturbation potentials at the angular Einstein ring radius are the same, which is indeed the generalization of LCIP to include convergent perturbations.  Let us think of the potential \\begin{equation} \\label{ppo} \\delta\\psi(r,\\phi)= \\sum_n\\left[\\epsilon_n(r)\\cos n\\phi+\\varrho_n(r)\\sin n\\phi\\right]= \\frac{1}{2}\\sum_n\\left[\\frac{z^n}{r^n}\\bar\\varepsilon_n(r)+ \\frac{r^n}{z^n}\\varepsilon_n(r)\\right] \\end{equation} perturbing a circularly symmetric lens system. Here, $\\varepsilon_n= \\epsilon_n+\\I\\varrho_n$. Then, the lens equation is given by \\begin{equation} \\label{gle} \\zeta=z\\left(1-\\frac{4G}{D_\\mathrm{rel}c^2}\\frac{M}{r^2}\\right)-\\delta\\alpha, \\end{equation} where \\begin{equation} \\label{dap} \\delta\\alpha=2\\upartial_{\\z}\\delta\\psi=\\sum_n \\left[\\frac{z^{n+1}}{2r}\\frac{d}{dr}\\left(\\frac{\\bar\\varepsilon_n}{r^n}\\right) +\\frac{1}{z^{n-1}}\\frac{1}{2r}\\frac{d}{dr}\\left(r^n\\varepsilon_n\\right)\\right] =\\frac{r_0}{2}\\sum_n \\left[\\frac{r_0^{n+1}}{\\z^{n+1}}\\bar\\fz_n -\\frac{\\z^{n-1}}{r_0^{n-1}}\\gz_n\\right]. \\end{equation} \\begin{equation} \\fz_n\\!(r)=\\frac{r^{2n+1}}{r_0^{n+2}} \\frac{d}{dr}\\left(\\frac{\\varepsilon_n}{r^n}\\right) \\,;\\ \\ \\ \\gz_n\\!(r)=-\\frac{r_0^{n-2}}{r^{2n-1}} \\frac{d}{dr}\\left(r^n\\varepsilon_n\\right) \\end{equation} Here, the calculation utilizes the fact that $r^2=z\\z$ and consequently that $\\upartial_{\\z} r=z(2r)^{-1}=r(2\\z)^{-1}$. That is, for any real function $F(r)$, one can relate the complex derivative and the ordinary derivatives via $\\upartial_{\\z} F=F'\\upartial_{\\z} r=zF' (2r)^{-1}=rF'(2\\z)^{-1}$. To find the expression for the linear approximation of the caustics, one series-expands equation~(\\ref{gle}) and its Jacobian determinant at $z_0=r_0\\E^{\\I\\phi}$ as in Section~\\ref{ssec_ccc}. After some tedious but none the less straightforward algebra, one can find \\begin{equation} \\label{cvj} \\J=2(1-\\kappa_{s,0})\\,\\frac{\\zeta-\\zeta_c(\\phi)}{z_0} \\end{equation} \\begin{equation} \\label{cvc} \\zeta_c(\\phi)=\\frac{r_0}{4}\\sum_n\\left[ (n+1)\\left(\\fz_{n,0}+\\gz_{n,0}\\right)\\E^{-(n-1)\\I\\phi}+ (n-1)\\overline{\\left(\\fz_{n,0}+\\gz_{n,0}\\right)}\\E^{(n+1)\\I\\phi} \\right] \\end{equation} in place of equations~(\\ref{ljc}) and (\\ref{cau}), where $\\kappa_{s,0}$ is the (local) convergence of the unperturbed circularly symmetric lens at $r=r_0$ (note that $\\beta_0= 2\\kappa_{s,0}$ at $r=r_0$), and the subscript `$_0$' indicates the value of the function at $r=r_0$ [e.g., $\\fz_{n,0}= \\fz_n(r_0)$ and $\\gz_{n,0}=\\gz_n(r_0)$]. Then, \\begin{equation} \\fz_{n,0}=\\fz_n\\!(r_0)=r_0^{n-1} \\left.\\frac{d}{dr}\\left(\\frac{\\varepsilon_n}{r^n}\\right)\\right|_{r=r_0}= \\frac{\\varepsilon^\\prime_{n,0}}{r_0}-n\\frac{\\varepsilon_{n,0}}{r_0^2} \\end{equation} \\begin{equation} \\gz_{n,0}=\\gz_n\\!(r_0)=-\\frac{1}{r_0^{n+1}} \\left.\\frac{d}{dr}\\left(r^n\\varepsilon_n\\right)\\right|_{r=r_0}= -\\frac{\\varepsilon^\\prime_{n,0}}{r_0}-n\\frac{\\varepsilon_{n,0}}{r_0^2} \\end{equation} so that \\begin{equation} \\zeta_c(\\phi)=-\\frac{r_0}{2}\\sum_nn\\left[ (n+1)\\E^{-(n-1)\\I\\phi}\\frac{\\varepsilon_{n,0}}{r_0^2}+ (n-1)\\E^{(n+1)\\I\\phi}\\frac{\\bar\\varepsilon_{n,0}}{r_0^2}\\right], \\end{equation} or \\begin{equation} \\frac{\\Re[\\zeta_c]}{r_0}=-\\sum_nn\\left[ \\frac{\\epsilon_{n,0}}{r_0^2}(n\\cos n\\phi\\cos\\phi+\\sin n\\phi\\sin\\phi)+ \\frac{\\varrho_{n,0}}{r_0^2}(n\\sin n\\phi\\cos\\phi-\\cos n\\phi\\sin\\phi)\\right] \\end{equation} \\begin{equation} \\frac{\\Im[\\zeta_c]}{r_0}=-\\sum_nn\\left[ \\frac{\\epsilon_{n,0}}{r_0^2}(n\\cos n\\phi\\sin\\phi-\\sin n\\phi\\cos\\phi)+ \\frac{\\varrho_{n,0}}{r_0^2}(n\\sin n\\phi\\sin\\phi+\\cos n\\phi\\cos\\phi)\\right]. \\end{equation} Note that the expressions do not involve the radial derivatives at all. In other words, up to the linear order, the caustics are completely specified by $\\varepsilon_{n,0}=\\epsilon_n(r_0)+ \\I\\varrho_n(r_0)$ or the azimuthal behaviour of the potential $\\delta\\psi(r_0,\\phi)$ exactly \\emph{at} $r=r_0$, and independent of its radial behaviour even in its immediate neighbourhood. In terms of the deflections due to the perturbations (instead of the perturbation potentials) given as in equation~(\\ref{dap}), this translates into the linear approximation of the caustics being a function of $\\fz_{n,0}+\\gz_{n,0}$ alone, and thus the two perturbations producing the same linear approximation of the caustics if $\\fz_{n,0}+\\gz_{n,0}$ is the same. The null-convergent LCIP pair is a special case of a pair of two perturbations given respectively by $(\\fz_n,\\gz_n)=(c_n,0)$ and $(\\fz_n,\\gz_n)=(0,c_n)$ where the $c_n$'s are (complex) constants. In terms of the potentials, they are the pair of $\\varepsilon^t_n=a_n(r/r_0)^n$ and $\\varepsilon^p_n=a_n(r_0/r)^n$ where the $a_n$'s are again some common constants.  \\subsubsection{ellipticity, external shear, and quadrupole moment} In many applications of gravitational lensing, the lensing convergence may be elliptically symmetric, that is, $\\kappa(r,\\phi)= \\kappa_a[a(r,\\phi)]$ is the function of $r$ and $\\phi$ only through the combination $a=r\\rho^{-1}(\\sin^2\\phi+\\rho^2\\cos^2\\phi)^{1/2}$, where $\\rho$ is a constant. It is straightforward to show that this yields concentric ellipses with eccentricity of $|1-\\rho^2|^{1/2}$ as iso-convergence contours. The analysis of these types of lensing systems is of great interest, but also very much more complicated than the circularly symmetric system. However, if $\\rho\\simeq1$, the problem can be greatly simplified by applying Fourier analysis. Let us consider the radial Taylor expansion of $\\kappa(r,\\phi)$ with fixed $\\phi$, that is, $\\kappa(r,\\phi)\\approx\\kappa(\\tilde r,\\phi)+(d\\kappa/dr) |_{r=\\tilde r} (r-\\tilde r)$. Suppose that the expansion was at $\\tilde r=2^{1/2}a (1+\\rho^{-2})^{-1/2}$. Then the convergence at $r=a\\rho(\\sin^2\\phi+\\rho^2\\cos^2\\phi)^{-1/2}$ is \\begin{equation} \\label{ell} \\kappa_a(a)\\approx\\kappa(\\tilde r,\\phi)\\left\\{1+ \\left.\\frac{d\\ln\\kappa}{d\\ln r}\\right|_{r=\\tilde r} \\left[\\left(1-\\frac{1-\\rho^2}{1+\\rho^2}\\cos2\\phi\\right)^{-1/2} -1\\right]\\right\\}. \\end{equation} Here, by noticing that $\\kappa_a(a)$ and $(d\\ln\\kappa/d\\ln r)|_{r= \\tilde r}$ do not involve any azimuthal dependence, equation~(\\ref{ell}) can be used to derive the quadrupole approximation \\citep[c.f.,][]{Ko87a} of the elliptically symmetric convergence \\begin{equation} \\kappa(\\tilde r,\\phi)\\approx \\kappa_a\\left[\\frac{\\tilde r}{\\rho}\\sqrt{\\frac{1+\\rho^2}{2}}\\right] \\left(1-\\frac{1}{2}\\,\\frac{1-\\rho^2}{1+\\rho^2} \\left.\\frac{d\\ln\\kappa}{d\\ln r}\\right|_{r=\\tilde r} \\cos2\\phi\\right), \\end{equation} provided that $|1-\\rho^2|\\ll1$.  Suppose that the lensing convergence is given by \\begin{equation} \\label{con} \\kappa(r,\\phi)=\\kappa_s(r) \\left\\{1+\\xi\\cos[2(\\phi-\\phi_e)]\\right\\}. \\end{equation} Then, with the canonical boundary condition, the potential can be found, using Green's function, to be \\begin{equation} \\psi(r,\\phi)=2\\left[ \\ln r\\int_0^r\\!\\D\\hat r\\,\\hat r\\kappa_s(\\hat r)+ \\int_r^\\infty\\!\\D\\hat r\\,\\hat r\\kappa_s(\\hat r)\\ln\\hat r\\right] -\\frac{1}{2}\\left\\{ \\frac{1}{r^2}\\int_0^r\\!\\D\\hat r\\,\\hat r^3\\xi\\kappa_s(\\hat r) \\cos[2(\\phi-\\phi_e)]+r^2\\int_r^\\infty\\!\\frac{\\D\\hat r}{\\hat r}\\, \\xi\\kappa_s(\\hat r)\\cos[2(\\phi-\\phi_e)]\\right\\}. \\end{equation} If one defines radial functions $\\mu_0(r)\\equiv\\int_0^r\\kappa_s \\hat r\\D\\hat r$, which is basically a scaled mass associated with the circularly symmetric part of the convergence, and $\\mu_1(r)\\equiv \\int_0^r\\xi\\kappa_s\\hat r\\D\\hat r$, then the potential (up to an addictive constant) may also be written in (assuming $\\phi_e$ is constant) \\begin{equation} \\label{pot} \\psi(r,\\phi)=2\\psi_0+ 2\\int_{\\underline{r}}^r\\!\\D\\hat r\\,\\frac{\\mu_0(\\hat r)}{\\hat r} -\\frac{1}{2}\\left[ \\frac{1}{r^2}\\int_0^r\\!\\D\\hat r\\,\\hat r^2\\mu_1^\\prime(\\hat r) +r^2\\int_r^\\infty\\!\\D\\hat r\\, \\frac{\\mu_1^\\prime(\\hat r)}{\\hat r^2}\\right]\\cos[2(\\phi-\\phi_e)], \\end{equation} where $\\underline{r}$ is some fiducial radius, $\\psi_0$ is a constant that may be formally defined by $\\psi_0\\equiv\\mu_0(\\underline{r}) \\ln\\underline{r}+\\int_{\\underline{r}}^\\infty\\mu_0'(r)\\ln r\\D r$. In addition, one can also find the deflection function \\begin{equation} \\label{def} \\alpha=2\\upartial_{\\z}\\psi=2\\frac{\\mu_0(r)}{r^2}z+ \\frac{1}{\\z^3}\\int_0^r\\!\\D\\hat r\\,\\hat r^2\\mu_1^\\prime(\\hat r)\\E^{-2\\I\\phi_e} -\\z\\int_r^\\infty\\!\\D\\hat r\\,\\frac{\\mu_1^\\prime(\\hat r)}{\\hat r^2} \\E^{2\\I\\phi_e}. \\end{equation} One can note that the potential given by equation~(\\ref{pot}) can be considered as a circularly symmetric potential perturbed by a small perturbing potential of the form of equation~(\\ref{ppo}) with $n=2$ provided that both of the integrals are small. Assuming that the circularly symmetric part of the potential in equation~(\\ref{pot}) allows to form an Einstein ring at $r=r_0$ [i.e., $2\\mu_0(r_0)= r_0^2$], the linear approximation of the tangential caustics is given by \\begin{equation} \\label{cca} \\frac{\\zeta_c(\\phi)}{r_0\\E^{\\I\\phi_e}}= \\frac{(\\fz^\\star_0+\\gz^\\star_0)}{2} \\left[3\\E^{-\\I(\\phi-\\phi_e)}+\\E^{3\\I(\\phi-\\phi_e)}\\right] =2(\\fz^\\star_0+\\gz^\\star_0) \\left[\\cos^3(\\phi-\\phi_e)-\\I\\sin^3(\\phi-\\phi_e)\\right] \\end{equation} \\begin{equation} \\fz^\\star_0=\\int_0^{r_0}\\!\\frac{\\D r}{r_0} \\left(\\frac{r}{r_0}\\right)^2 \\frac{\\mu_1^\\prime(r)}{r_0} \\,;\\ \\ \\ \\gz^\\star_0=\\int_{r_0}^\\infty\\!\\frac{\\D r}{r_0} \\left(\\frac{r_0}{r}\\right)^2 \\frac{\\mu_1^\\prime(r)}{r_0}. \\end{equation} Here, the caustic is a tetracuspoid so that there are four images for a source position within the caustic. Equation~(\\ref{cca}), from a comparison to equations~(\\ref{ccc}), (\\ref{cri}), and (\\ref{are}), also implies that the cross-section of quadruple images is $(3\\upi/2)(\\fz^\\star_0+ \\gz^\\star_0)^2$. Furthermore, the lensing behaviour of the source near the lens centre is essentially described by two \\emph{constants} $\\fz^\\star_0$ and $\\gz^\\star_0$. In particular, when the source and the centre of the lens are perfectly aligned ($\\zeta=0$), there are four images associated with the caustics, and their azimuthal coordinates are given by $\\phi=n\\upi/2+\\phi_e$ where $n$ is an integer. It is easy to verify that equation~(\\ref{cvj}) is real and that \\begin{equation} z=(r_0+\\delta r)\\E^{\\I\\phi}= r_0\\E^{\\I\\phi_e}\\left[1+(-1)^n \\frac{(\\fz^\\star_0-\\gz^\\star_0)}{2(1-\\kappa_{s,0})} \\right]\\E^{n\\I\\upi/2} \\,;\\ \\ \\ \\J=(-1)^{n+1}4(1-\\kappa_{s,0}) (\\fz^\\star_0+ \\gz^\\star_0) \\end{equation} so that the resulting images are an alternating-parity equal-magnification two-fold symmetric crucifix-form quartet with an axis ratio of $1-|(\\fz^\\star_0-\\gz^\\star_0)/(1-\\kappa_{s,0})|$. For more general source positions, once the caustic is drawn using equation~(\\ref{cca}), the azimuthal coordinates and the individual magnifications of the images can be found by following the same procedure as in Section~\\ref{ssec_wcc} using equation~(\\ref{cvj}). In addition, since the form of the perturbation here is specified, it is also possible to find the radial coordinates of the images by the equivalent of equation~(\\ref{rad}) \\begin{equation} \\label{dtr} \\delta r= \\frac{\\E^{-\\I\\phi}}{2(1-\\kappa_{s,0})}\\left\\{\\zeta+r_0\\left[ \\E^{\\I(3\\phi-2\\phi_e)}\\fz^\\star_0 -\\E^{\\I(2\\phi_e-\\phi)}\\gz^\\star_0 \\right]\\right\\}=\\frac{\\zeta \\E^{-\\I\\phi}+r_0\\left[ (\\fz^\\star_0-\\gz^\\star_0)\\cos[2(\\phi-\\phi_e)] +\\I(\\fz^\\star_0+\\gz^\\star_0)\\sin[2(\\phi-\\phi_e)] \\right\\}}{2(1-\\kappa_{s,0})}, \\end{equation} for the given radial coordinate of the images $\\phi$. Algebraically, $\\phi$ can be determined from the condition that $\\delta r$ is real, that is, $\\Im[\\zeta \\E^{-\\I\\phi}]+(\\fz^\\star_0+ \\gz^\\star_0)\\sin[2(\\phi-\\phi_e)]=0$. In other words, for a given source position, $\\phi$ is determined by $\\fz^\\star_0 +\\gz^\\star_0$, which is in accordance with the geometric construction of $\\phi$ since it is the same coefficient for the caustics (eq.~\\ref{cca}).  Note the linear approximation of the caustics (eq.~\\ref{cca}) is identical (after a proper rotation of coordinate) to those caused either by the null-convergent compact quadrupole perturbation with the quadrupole moment of $(\\fz^\\star_0+\\gz^\\star_0)r_0^4$ or by the null-convergent tidal perturbation by a constant external shear of $-(\\fz^\\star_0+\\gz^\\star_0)$ \\citep*[c.f.,][]{WMS95}. In fact, the examination of equation~(\\ref{def}) suggests that those locally (near Einstein ring region) null-convergent perturbations can be caused by the quadrupole moment of the mass distribution, either entirely enclosed within the angular Einstein ring radius ($f\\propto z^{-3}$), or lying completely outside of the angular Einstein ring radius ($f\\propto z$). In general, the lensing behaviour for the source near the lens centre when the convergence is given by equation~(\\ref{con}) with $|\\xi|\\ll1$ is basically equivalent to that of the system with a circularly symmetric convergence $\\kappa_s$ under the perturbation of $\\f=r_0 [\\bar\\fz^\\dag_0(\\z/r_0)^{-3}-\\gz^\\dag_0(\\z/r_0)]$, where $\\fz^\\dag_0=\\fz^\\star_0\\E^{2\\I\\phi_e}$ and $\\gz^\\dag_0= \\gz^\\star_0\\E^{2\\I\\phi_e}$. This can be generalized to the elliptically symmetric system that is additionally affected by external shear and/or compact quadrupole mass moment.\\footnote{In much of literature, a quadrupole lens and a lens with shear are used interchangeably without distinction. Although this practice has a certain justification since both the effects due to the constant external shear and the compact quadrupole moment are associated with $\\cos2\\phi$ terms in the azimuthal part of the potential, they are always explicitly distinguished in this paper.} For this case, $\\fz^\\dag_0= \\fz^\\star_0\\E^{2\\I\\phi_e}+\\hat Q_4$ and $\\gz^\\dag_0= \\gz^\\star_0\\E^{2\\I\\phi_e}+\\gamma_0$,\\footnote{The sign for the external shear is chosen to follow the usual convention in strong lensing literatures \\citep[e.g.,][]{HKM04}. This choice leads the lensing shear due to the external shear term to be `$-\\gamma_0$,' as defined in Appendix~\\ref{asec_ceq}.} and they are not necessarily parallel to each other any more. Here, note that $\\hat Q_4=|\\hat Q_4|\\E^{2\\I\\phi_Q}$ and $\\gamma_0=|\\gamma_0| \\E^{2\\I\\phi_\\gamma}$ are complex numbers. (Strictly speaking, they are complex number representations of certain symmetric 2$\\times$2 traceless tensors. The absolute value of the complex number corresponds to the positive eigenvalue of the tensor and its argument is the same as twice the angle between the corresponding eigenvector and the real direction.) Since equations~(\\ref{cvj}) and (\\ref{cvc}) are still valid regardless of $\\fz_n$ and $\\gz_n$ not being parallel, one finds, in the place of equation~(\\ref{cca}), -- note the factor of two difference in the definition compared to eq.~(\\ref{dap}) -- that \\begin{equation} \\zeta_c(\\phi)=\\frac{r_0\\E^{\\I\\phi_+}}{2} \\left|\\fz^\\dag_0+\\gz^\\dag_0\\right| \\left[3\\E^{-\\I(\\phi-\\phi_+)}+\\E^{3\\I(\\phi-\\phi_+)}\\right] =2r_0\\E^{\\I\\phi_+}\\left|\\fz^\\dag_0+\\gz^\\dag_0\\right| \\left[\\cos^3(\\phi-\\phi_+)-\\I\\sin^3(\\phi-\\phi_+)\\right] \\end{equation} and also that \\begin{equation} \\delta r= \\frac{\\E^{-\\I\\phi}}{2(1-\\kappa_{s,0})}\\left\\{\\zeta+r_0\\left[ \\E^{3\\I\\phi}\\bar\\fz^\\dag_0-\\E^{-\\I\\phi}\\gz^\\dag_0 \\right]\\right\\}=\\frac{\\zeta \\E^{-\\I\\phi}+r_0\\left\\{ \\left|\\fz^\\dag_0-\\gz^\\dag_0\\right|\\cos[2(\\phi-\\phi_-)] +\\I\\left|\\fz^\\dag_0+\\gz^\\dag_0\\right|\\sin[2(\\phi-\\phi_+)] \\right\\}}{2(1-\\kappa_{s,0})}, \\end{equation} in place of equation~(\\ref{dtr}). Here, $\\fz^\\dag_0+ \\gz^\\dag_0=|\\fz^\\dag_0+\\gz^\\dag_0|\\E^{2\\I\\phi_+}$ and $\\fz^\\dag_0-\\gz^\\dag_0=|\\fz^\\dag_0-\\gz^\\dag_0| \\E^{2\\I\\phi_-}$. As before, $\\delta r\\in\\mathbf R$ for $\\phi$ that is the azimuthal coordinate of the image for the given source position. For $\\zeta=0$, one finds the azimuthal coordinates of images $\\phi=n\\upi/2+\\phi_+$ and also that $2(1-\\kappa_{s,0})\\delta r=(-1)^n r_0|\\fz^\\dag_0-\\gz^\\dag_0|\\cos[2(\\phi_+-\\phi_-)]$. This implies that the resulting image quartets are identical for the set of $\\fz^\\dag_0$ and $\\gz^\\dag_0$ if $|\\fz^\\dag_0+\\gz^\\dag_0|$ and $|\\fz^\\dag_0-\\gz^\\dag_0||\\cos[2(\\phi_+-\\phi_-)]|$ are constant. For a more general source position $\\zeta\\ne0$, while one finds that the azimuthal coordinates and the magnification ratios of images are completely specified by $\\fz^\\dag_0+\\gz^\\dag_0$, the projection that controls the radial coordinates depends on each azimuthal coordinate $\\phi$ of the image, so that they are in general not common for the four images. However, if the source is sufficiently close to the centre, the projections are more or less parallel to one another, and therefore, there still exists a certain near degeneracy of the images concerning the choice of $|\\fz^\\dag_0|$, $|\\gz^\\dag_0|$, and their relative orientation in the underlying lens model.  If the convergence is given by a power-law ellipsoid \\begin{equation} \\kappa[a(r,\\phi)]=\\frac{\\beta}{2}\\left(\\frac{a_0}{a}\\right)^{2-\\beta}= \\frac{\\beta}{2}\\left(\\frac{a_0}{r}\\right)^{2-\\beta} \\left(\\rho^{-2}\\sin^2\\phi+\\cos^2\\phi\\right)^{-(2-\\beta)/2}= \\frac{\\beta}{2}\\left(\\frac{a_0}{r}\\right)^{2-\\beta} \\left(\\frac{2\\rho^2}{1+\\rho^2}\\right)^{(2-\\beta)/2} \\left(1-\\frac{1-\\rho^2}{1+\\rho^2}\\cos2\\phi\\right)^{-(2-\\beta)/2}, \\end{equation} it can be approximated by a quadrupole convergence in a form of equation~(\\ref{con}) with \\begin{equation} \\kappa_s=\\frac{\\beta}{2}\\left(\\frac{a_0}{r}\\right)^{2-\\beta} \\left(\\frac{2\\rho^2}{1+\\rho^2}\\right)^{(2-\\beta)/2}\\,;\\ \\ \\ \\xi=\\frac{1-\\rho^2}{1+\\rho^2}\\left(1-\\frac{\\beta}{2}\\right). \\end{equation} Here, $0<\\beta<2$. In addition, \\begin{equation} \\mu_0(r)=\\int_0^r\\!\\D\\hat r\\,\\hat r\\kappa_s(\\hat r)= \\frac{\\beta}{2}\\left(\\frac{2\\rho^2}{1+\\rho^2}\\right)^{(2-\\beta)/2} a_0^{2-\\beta}\\int_0^r\\!\\D\\hat r\\,\\hat r^{\\beta-1}= \\left(\\frac{2\\rho^2}{1+\\rho^2}\\right)^{(2-\\beta)/2} \\frac{a_0^2}{2}\\left(\\frac{r}{a_0}\\right)^\\beta \\end{equation} so that one finds $r_0=2^{1/2}a_0(1+\\rho^{-2})^{-1/2}$ by solving $2\\mu_0(r_0)=r_0^2$. Furthermore, $\\mu_1(r)=\\xi\\mu_0(r)$ since $\\xi$ is now constant, and $2\\kappa_{s,0}=\\beta$. Note that $2\\mu_0(r)=r_0^{2-\\beta}r^\\beta$ and $(d\\ln\\mu_0/d\\ln r)= (d\\ln\\mu_1/d\\ln r)=\\beta$. Finally, \\begin{equation} \\fz^\\star_0=\\int_0^{r_0}\\!\\frac{\\D r}{r_0} \\left(\\frac{r}{r_0}\\right)^2\\frac{\\mu_1^\\prime(r)}{r_0}= \\xi\\frac{\\beta}{2}\\int_0^1\\!\\D\\rdash\\,\\rdash^{\\beta+1}= \\frac{\\xi\\beta}{2(\\beta+2)}= \\frac{1-\\rho^2}{1+\\rho^2}\\frac{\\beta}{4}\\frac{2-\\beta}{2+\\beta}, \\end{equation} \\begin{equation} \\gz^\\star_0=\\int_{r_0}^\\infty\\!\\frac{\\D r}{r_0} \\left(\\frac{r_0}{r}\\right)^2\\frac{\\mu_1^\\prime(r)}{r_0}= \\xi\\frac{\\beta}{2}\\int_1^\\infty\\!\\D\\rdash\\,\\rdash^{\\beta-3}= -\\frac{\\xi\\beta}{2(\\beta-2)}= \\frac{1-\\rho^2}{1+\\rho^2}\\frac{\\beta}{4}, \\end{equation} and thus \\begin{equation} \\fz^\\star_0+\\gz^\\star_0= \\frac{1-\\rho^2}{1+\\rho^2}\\frac{\\beta}{2+\\beta} \\,;\\ \\ \\ \\fz^\\star_0-\\gz^\\star_0= -\\frac{1-\\rho^2}{1+\\rho^2}\\frac{\\beta^2}{2(2+\\beta)}. \\end{equation} Hence, for a singular isothermal ellipsoid (SIE) lens ($2\\kappa_{s,0}=\\beta=1$), with an ellipticity parameter $\\ee=(1-\\rho^2)/(1+\\rho^2)\\simeq 1-\\rho$, one finds that $\\fz^\\star_0+\\gz^\\star_0=\\ee/3$ and $\\fz^\\star_0-\\gz^\\star_0=-\\ee/6$, and thus, the four-image cross-section is given by $(3\\upi/2)(\\fz^\\star_0+\\gz^\\star_0)^2=(\\upi/6)\\ee^2$, and the axis ratio for a crucifix-form quartet is $1-|(\\fz^\\star_0-\\gz^\\star_0)/(1-\\kappa_{s,0})|=1-(\\ee/3)$ \\citep*{KKS97}.  If this SIE lens is subject to perturbations due to the constant external shear, one finds that \\begin{equation} \\fz_0^\\dag= \\frac{\\ee}{12}\\E^{2\\I\\phi_e}= \\frac{1}{12}\\left(\\ee\\E^{2\\I\\phi_e}+4|\\gamma_0|e^{2\\I\\phi_\\gamma}\\right) -\\frac{|\\gamma_0|}{3}\\E^{2\\I\\phi_\\gamma}\\,;\\ \\ \\ \\gz_0^\\dag= \\frac{\\ee}{4}\\E^{2\\I\\phi_e}+|\\gamma_0|\\E^{2\\I\\phi_\\gamma}= \\frac{1}{4}\\left(\\ee\\E^{2\\I\\phi_e}+4|\\gamma_0|e^{2\\I\\phi_\\gamma}\\right). \\end{equation} By comparing this to another SIE lens system that is affected by the compact quadrupole moment, $\\fz_0^\\dag= (\\tilde\\ee/12)\\E^{2\\I\\tilde\\phi_e}+|\\hat Q_4|\\E^{2\\I\\phi_Q}$ and $\\gz_0^\\dag=(\\tilde\\ee/4)\\E^{2\\I\\tilde\\phi_e}$, one can conclude that the lensing behaviours of two systems are identical up to the linear approximation provided that the external shear and the compact quadrupole moment are related by $|\\gamma_0|=3|\\hat Q_4|$ and $\\E^{2\\I(\\phi_Q-\\phi_\\gamma)}=-1$ and the two ellipticity parameters by \\begin{equation} \\tilde\\ee\\E^{2\\I(\\tilde\\phi_e-\\phi_e)}=\\ee\\left[1+ \\frac{4|\\gamma_0|}{\\ee}\\E^{2\\I(\\phi_\\gamma-\\phi_e)}\\right] \\,;\\ \\ \\ \\ee\\E^{2\\I(\\phi_e-\\tilde\\phi_e)}=\\tilde\\ee\\left[1+ \\frac{12|\\hat Q_4|}{\\tilde\\ee}\\E^{2\\I(\\phi_Q-\\tilde\\phi_e)}\\right]. \\end{equation} In other words, the lens system that can be modelled by a SIE lens with an external shear can also be modelled by a different SIE lens with a quadrupole moment, the magnitude of which is a third of that of the external shear, provided that the system's departure from the circular symmetry is small. This can be further generalized to systems that are subject to an external shear and a compact quadrupole moment at the same time to yield a certain degeneracy of the lens model regarding the ellipticity, the shear, and the quadrupole moment (as well as their relative orientations).  If one models an image quartet using a SIE lens with a constant external shear but no quadrupole moment, then $\\fz_0^\\dag+\\gz_0^\\dag=|\\fz_0^\\dag+\\gz_0^\\dag|\\E^{2\\I\\phi_+}= (\\ee/3)\\E^{2\\I\\phi_e}+|\\gamma_0|\\E^{2\\I\\phi_\\gamma}$ and $\\fz_0^\\dag-\\gz_0^\\dag=|\\fz_0^\\dag-\\gz_0^\\dag|\\E^{2\\I\\phi_-}= -(\\ee/6)\\E^{2\\I\\phi_e}-|\\gamma_0|\\E^{2\\I\\phi_\\gamma}$. Hence, the four-image cross-section is found to be \\begin{equation} \\frac{3\\upi}{2}\\left|\\fz_0^\\dag+\\gz_0^\\dag\\right|^2= \\upi\\left\\{\\frac{\\ee^2}{6}+\\frac{3|\\gamma_0|^2}{2}+ \\ee|\\gamma_0|\\cos[2(\\phi_\\gamma-\\phi_e)]\\right\\}. \\end{equation} It is easy to see that, for given $\\ee$ and $|\\gamma_0|$, the cross-section is largest [$=\\upi(\\ee+3|\\gamma_0|)^2/6$] if $\\phi_\\gamma=\\phi_e$ while it is at minimum [$=\\upi(\\ee-3|\\gamma_0|)^2/6$] when $\\phi_\\gamma-\\phi_e=\\upi/2$. In addition, one can find that \\begin{equation} \\left(\\frac{\\ee}{3}\\right)^2-|\\gamma_0|^2= |\\fz_0^\\dag+\\gz_0^\\dag|\\left\\{3|\\fz_0^\\dag+\\gz_0^\\dag| +4|\\fz_0^\\dag-\\gz_0^\\dag|\\cos[2(\\phi_+-\\phi_-)]\\right\\}. \\end{equation} That is to say, for fixed $|\\fz_0^\\dag+\\gz_0^\\dag|$ and $|\\fz_0^\\dag-\\gz_0^\\dag|\\cos[2(\\phi_+-\\phi_-)]$, which produce identical crucifix-form image quartets, there exists a degeneracy running along a hyperbolic path in ($\\ee$, $|\\gamma_0|$) space. This may be seen as a prototype for the `cancellation branches' of the `{\\sf U}' shape degenerate path found in fig.~8 of \\citet{KKS97}. In reality, it is highly unlikely that the source is perfectly aligned with the lens centre so that the actual degeneracy becomes rather more complex. (Furthermore, as the ellipticity and/or the shear get larger, the whole perturbative approach will start to break down.) Nevertheless, the basic idea that there exists a certain set of $\\fz_0^\\dag+\\gz_0^\\dag$ and $\\fz_0^\\dag-\\gz_0^\\dag$ that yields similar (or nearly degenerate) image configurations (with the proper rotation of frames) is thought to be more or less valid, and this can lead to the observed degeneracy of the shear and ellipticity combination."
    },
    "0410/astro-ph0410031_arXiv.txt": {
        "abstract": "We consider general curvature-invariant modifications of the Einstein-Hilbert action that become important only in regions of extremely low space-time curvature. We investigate the far future evolution of the universe in such models, examining the possibilities for cosmic acceleration and other ultimate destinies. The models generically possess de Sitter space as an unstable solution and exhibit an interesting set of attractor solutions which, in some cases, provide alternatives to dark energy models. ",
        "introduction": "The acceleration of the universe presents one of the greatest problems in theoretical physics today. The increasingly accurate observations of type Ia supernovae light-curves, coupled with exquisite measurements of CMB anisotropies and large scale structure data~\\cite{Riess:1998cb,Perlmutter:1998np,Tonry:2003zg,Bennett:2003bz,Netterfield:2001yq,Halverson:2001yy} have forced this issue to the forefront of those facing particle physicists, cosmologists and gravitational physicists alike.  This problem has been attacked head on, but no compelling, well-developed and well-motivated solutions have yet emerged. While much work has focused on the search for new matter sources that yield accelerating solutions to general relativity, more recently some authors have turned to the complementary approach of examining whether new gravitational physics might be responsible for cosmic acceleration.  There have been a number of different attempts~\\cite{Deffayet:2001pu,Freese:2002sq,Arkani-Hamed:2002fu,Dvali:2003rk,Nojiri:2003wx,Carroll:2003wy,Capozziello:2003tk,Arkani-Hamed:2003uy,Abdalla:2004sw} to modify gravity to yield accelerating cosmologies at late times. The path we are concerned with in this paper is the direct addition of higher order curvature invariants to the Einstein-Hilbert action. The first example of this was provided by the model of Carroll, Duvvuri, Trodden, and Turner (CDTT)~\\cite{Carroll:2003wy} (see also~\\cite{Capozziello:2003tk}). For subsequent work on various aspects and extensions of this model, see \\cite{Vollick,Dick,Nojiri_1,Dolgov,Nojiri_2,Meng1,Chiba,Meng2,Flanagan,Woodard,Ezawa,Flanagan_2,Rajaraman,Vollick_2,Nunez,Allemandi,Lue:2003ky}. In particular, the simplest model has been shown to conflict with solar system tests of gravity~\\cite{Dolgov,Chiba,Woodard}.  Our approach is purely phenomenological. The evidence for cosmic acceleration is very sound. In pure Einstein gravity, as matter dilutes away in the expanding universe, the expansion rate inevitably slows. Our question is can this be avoided within the gravitational sector of the theory, as opposed to adding new energy sources. One way is by adding a cosmological constant. In this paper we explore a wider class of modifications that share the feature of late-time accelerating behavior, thus fitting cosmological observations.  Consider a simple correction to the Einstein-Hilbert action, \\begin{equation} \\label{action} S =\\frac{\\mpl^2}{2}\\int d^4 x\\, \\sqrt{-g}\\left(R-\\frac{\\mu^4}{R}\\right) +\\int d^4 x\\, \\sqrt{-g}\\, {\\cal L}_M \\ . \\end{equation} where $\\mu$ is a new parameter with units of $[{\\rm mass}]$ and ${\\cal L}_M$ is the Lagrangian density for matter. This action gives rise to fourth-order equations of motion. In the case of an action depending purely on the Ricci scalar (and on its derivatives) it is possible to transform from the frame used in~(\\ref{action}), which we call the {\\em matter frame}, to an {\\em Einstein frame}, in which the gravitational Lagrangian takes the Einstein-Hilbert form and the additional degrees of freedom ($\\ddot H$ and $\\dot H$) are represented by a scalar field $\\phi$. The details of this can be found in~\\cite{Carroll:2003wy}. The scalar field is  minimally coupled to Einstein gravity, non-minimally coupled to matter, and has a potential given by \\begin{equation} \\label{potential} V(\\phi)=\\mu^2 \\mpl^2 \\exp\\left(-2\\sqrt{\\frac{2}{3}}\\frac{\\phi}{\\mpl} \\right)\\sqrt{\\exp \\left(\\sqrt{\\frac{2}{3}}\\frac{\\phi}{\\mpl} \\right)-1} \\ . \\end{equation}  Consider vacuum cosmological solutions. We must specify the initial values of $\\phi$ and $\\phi'$, denoted as $\\phi_i$ and ${\\phi'}_i$.  For simplicity we take $\\phi_i \\ll \\mpl$. There are three qualitatively distinct outcomes, depending on the value of ${\\phi'}_i$.  {\\em 1.  Eternal de Sitter.}  There is a critical value of ${\\phi'}_i \\equiv {\\phi'}_C$ for which $\\phi$ just reaches the maximum of the potential $V(\\phi)$ and comes to rest.  In this case the Universe asymptotically evolves to a de~Sitter solution. This solution requires tuning and is unstable, since any perturbation will induce the field to roll away from the maximum of its potential.  {\\em 2.  Power-Law Acceleration.}  For ${\\phi'}_i > {\\phi'}_C$, the field overshoots the maximum of $V(\\phi )$ and the Universe evolves to late-time power-law inflation, with observational consequences similar to dark energy with equation-of-state parameter $w_{\\rm DE}=-2/3$.  {\\em 3.  Future Singularity.}  For ${\\phi'}_i < {\\phi'}_C$, $\\phi$ does not reach the maximum of its potential and rolls back down to $\\phi =0$.  This yields a future curvature singularity.  In the more interesting case in which the Universe contains matter, it is possible to show that the three possible cosmic futures identified in the vacuum case remain in the presence of matter.  By choosing $\\mu\\sim 10^{-33}\\,$eV, the corrections to the standard cosmology only become important at the present epoch, making this theory a candidate to explain the observed acceleration of the Universe without recourse to dark energy.  Since we have no particular reason for choosing this value of $\\mu$, such a tuning is certainly not attractive. However, it is worth commenting that this small correction to the action, unlike most small corrections in physics, is destined to be important as the Universe evolves.  Clearly our choice of correction to the gravitational action can be generalized.  Terms of the form $-\\mu^{2(n+1)}/R^n$, with $n>1$, lead to similar late-time self acceleration, with behavior similar to a dark energy component with equation of state parameter \\begin{equation} \\label{geneos} w_{\\rm eff} = -1 + \\frac{2(n+2)}{3(2n+1)(n+1)} \\ . \\end{equation} Therefore, such modifications can easily accommodate current observational bounds~\\cite{Melchiorri:2002ux,Spergel:2003cb} on the equation of state parameter $-1.45< w_{\\rm DE} <-0.74$ ($95\\%$ confidence level). In the asymptotic regime, $n=1$ is ruled out at this level, while $n\\geq 2$ is allowed; even $n=1$ is permitted if we are near the top of the potential.  In this paper we seek to extend this approach. For the actions considered in this paper, a tranformation between the matter and Einstein frames does not necessarily make sense. Therefore, we would like to analyze the dynamics in the matter frame itself. A technique for this analysis is presented here.  As an example, we begin by applying this technique to the CDTT model. We will work in the matter frame with a spatially flat Robertson-Walker metric \\begin{equation} ds^2=-dt^2+a(t)^2\\,d\\vec x^2\\ , \\end{equation} with $a(t)$ being the scale factor. The time-time component of the field equations, the {\\it Friedmann equation} in this non-standard cosmology, is \\begin{equation} \\label{newfriedmann} 3H^2 - \\frac{\\mu^4}{12({\\dot H}+2H^2)^3}\\left(2H{\\ddot H} +15H^2{\\dot H}+2{\\dot H}^2+6H^4\\right) = \\frac{\\rho_M}{\\mpl^2}\\ , \\end{equation} where an overdot denotes differentiation with respect to cosmic time, and $H=\\dot a/a$.  To perform a phase-space analysis of such equations we write ${\\dot H}$ as a function of $H$. Let \\begin{eqnarray} x &=& -H(t) \\nonumber \\\\ y &=& {\\dot H}(t) \\ , \\end{eqnarray} so that \\begin{equation} \\ddot H =\\frac{d\\dot H}{dt} = \\frac{d\\dot H}{dH}\\,\\frac{dH}{dt} = - y\\,\\frac{dy}{dx}\\ . \\end{equation}  In this way we may write~(\\ref{newfriedmann}), a third order equation in $a(t)$, as a second order equation in $H(t)$ and hence as a first order equation in $y(x)$. The fact that the FRW equation for any $f(R)$ theory can be reduced to the first-order equation for the spatially flat case (second order in case of a non-zero spatial curvature) was first introduced in \\cite{Gurovich:1979xg}, \\cite{Starobinsky:1980te} .  Many cosmologically interesting solutions, including accelerating ones, are power-law solutions of the form $a(t)\\propto t^p$. In such cases ${\\dot H}=-H^2/p$ (i.e. $y=-x^2/p$). In anticipation of finding such solutions as asymptotic solutions to our equations, we define a new function $v(x)$ by \\begin{equation} v(x) = -\\frac{x^2}{y} \\ , \\end{equation} with $v\\neq 0$. Power-law solutions in the asymptotic future are then easily identified as  $v(x)\\rightarrow p= {\\rm constant}$ as $H\\rightarrow 0$.  Furthermore, if \\begin{equation} |\\dot H|=|y|\\to\\infty ,\\qquad{\\rm then}\\qquad |v|\\to 0\\quad \\textrm{if $x\\neq0$.} \\end{equation} So if $x$ is not zero, as $v\\to0$ we approach the singularity $|\\dot H|\\to\\infty$. This trick is invoked throughout this paper.  Here, let us apply it to the simple case of~(\\ref{newfriedmann}). The relevant first order equation is \\begin{equation} x\\,\\frac{dv}{dx}=2v + \\frac1{2\\mu^4} \\left[x^4\\,(36 - 216 v +432 v^2 -288 v^3) + \\mu^4\\,(2 v -15 v^2 + 6 v^3)\\right],\\label{eq_R} \\end{equation} with the resulting phase plot shown in figure~\\ref{Rmodelfig}. \\begin{figure}[ht] {\\centering \\includegraphics[width=3in]{R.eps} \\includegraphics[width=3in]{pR.eps} \\par } \\caption{Two phase portraits for the modified gravity model proposed in~\\cite{Carroll:2003wy}. The left portrait is in the coordinates $(x,v)$, for which an attractor at constant $v=p$ corresponds to a power-law solution with $a(t)\\propto t^p$. The right portrait is for the same theory in the $({\\dot H},H)$ plane, with the unstable de Sitter solution at $(0,1)$.} \\label{Rmodelfig} \\end{figure}  Since the $x$-axis on this plot is (minus) the Hubble parameter, earlier times in the universe lie to the (negative) left and late times lie closer to $x=0$ (in all but exponential or phantom evolution).  Note that the numerical solution shows the accelerating attractor $a(t)\\propto t^2$, corresponding to $v(x)\\rightarrow 2$ as $x\\rightarrow 0$, as expected from the analytic, Einstein-frame method. Indeed, for general $n$ this attractor, at $v=(2n+1)(n+1)/(n+2)$ can be obtained directly from the asymptotic form of the generalization of equation~(\\ref{eq_R}). In addition, our method pinpoints a singularity in the phase space, corresponding to a power-law evolution with exponent $p=1/2$ (that of radiation). Both features are also evident from the study of the asymptotically late time behavior of equation~(\\ref{eq_R}). In order to have a constant $v_0$ as a solution, we need \\begin{equation} 2v_0^2-5v_0+2=0\\ , \\end{equation} which has two real solutions: $v_0=1/2,2$, as expected.  The singularity at $p=1/2$ occurs because the Ricci scalar vanishes for this particular power. One might worry that this is problematic for describing the radiation-dominated phase in standard cosmology, since nucleosynthesis occurs during this epoch and provides a particularly strong constraint on deviations from the standard Friedmann equation at that time. However, as we shall see later, this is not a problem when matter sources are included explicitly, since even during radiation domination the Ricci scalar does not vanish exactly, but rather, has a small contribution from non-relativistic matter.  Nevertheless, the singularity is a new feature that was not found in the Einstein-frame analysis~\\cite{Carroll:2003wy}. This is presumably because  $R=0$ is a singular point of the conformal transformation used to reach the Einstein frame. We shall see similar singularities for some of the more general actions we consider in this paper.  Another way to visualize the solutions to our models is to use a more traditional phase portrait in the $({\\dot H},H)$ plane.  The outline of this paper is as follows. In the next section we shall introduce the general class of actions we are interested in. In section~\\ref{sec:vacuum} we analyze the vacuum equations, describing the singularity and attractor structure in detail before moving on to some simple special cases. In section~\\ref{sec:matter} we introduce matter into the equations, demonstrating briefly that the late-time attractor solutions of the system remain unchanged and setting up the formalism used in the appendix to establish stability of the system. In section~\\ref{sec:conclusions} we summarize our findings and comment on the status of these models as origins of cosmic acceleration. The paper contains two appendices. Appendix A contains definitions of a number of functions used in the body of the paper and Appendix B consists of a proof of the stability of the vacuum solutions under the addition of matter. ",
        "conclusions": "\\label{sec:conclusions} In the extreme low curvature regime, our only tests of general relativity are cosmological. The discovery of new phenomena at these scales may point to new matter sources, but alternatively may hint at hitherto undetected modifications of gravity.  The acceleration of the universe provides a particular challenge to modifications of gravity. Unlike the known perturbative corrections to the Einstein-Hilbert action arising from string theory, late time acceleration requires modifications that become important at extremely low energies, so low that only today, at the largest scales in the universe, is the resulting curvature low enough to lead to measurable deviations from general relativity.  These considerations led some of the current authors and others to consider new terms in the action for gravity that consist of inverse powers of the Ricci scalar. It is easy to show that such an approach introduces a de Sitter solution. However, this solution is unstable to a late-time accelerating power law attractor. For appropriate choices of parameters, this theory is a candidate to explain cosmic acceleration without the need for dark energy, although the simplest such theories are in conflict with solar system tests. For Lagrangians that are functions of only the Ricci scalar, there exists a map to an Einstein frame, in which the new degrees of freedom are represented by a scalar field. As a result, such modified gravity theories share many features in common with some dark energy models.  In this paper we have introduced a much more general class of modifications to the Einstein-Hilbert action, becoming relevant at extremely low curvatures. Specifically, we have considered inverse powers of arbitrary linear combinations of the curvature invariants $R^2$, $P\\equiv  R_{\\mu\\nu}\\,R^{\\mu\\nu}$ and $Q \\equiv R_{\\alpha\\beta\\gamma\\delta}\\,R^{\\alpha\\beta\\gamma\\delta}$. Such modifications are not simply equivalent to Einstein gravity plus scalar matter sources.  We have performed a general analysis of the late-time evolution of cosmological solutions to these theories.  Many of the theories exhibit late-time attractors of the form $a(t)\\propto t^p$, with $p$ some constant power. Indeed, there are often multiple such attractors. For a large class of theories there exists at least one attractor satisfying $p>1$, corresponding to cosmic acceleration.  The detailed structure of cosmological solutions to these theories turns out to be quite rich and varied, depending on the dimensionless parameters entering the particular linear combination. Two distinct types of singularities may exist, as well as the late-time power law attractors. We have identified all those theories for which the late time behavior is consistent with the observed acceleration of the universe, providing a whole new class of theories - generalized modified gravity theories - which are alternatives to dark energy.  The results we have found for the modification $f=\\mu^6/(a\\,R^2+b\\,P+c\\,Q)$ may be generalized to modifications \\begin{equation} f=\\frac{\\mu^{4n+2}}{(a\\,R^2+b\\,P+c\\,Q)^n} \\ , \\end{equation} using exactly the same calculational techniques.  In general there are power-law attractors with the following exponents, whenever they are real \\begin{equation} v^{\\rm gen}_{1,2}=\\frac{8n^2+10n+2-3\\alpha\\pm \\sqrt\\Gamma}{4(n+1)}\\ , \\end{equation} where \\begin{eqnarray} \\Gamma&=&9n^2\\alpha^2-(80n^3+116n^2+40n+4)\\,\\alpha\\nonumber\\\\ &&\\qquad{}+64n^4+160n^3+132n^2+40n+4 \\ . \\end{eqnarray} Clearly, as $n\\rightarrow \\infty$, the smaller attractor tends to 0, whereas the larger one increases linearly as $4n$.  Two special cases are \\begin{eqnarray} f &=& -\\frac{m^{4n+2}}{P^n} \\ , \\nonumber \\\\ f &=& -\\frac{M^{4n+2}}{Q^n} \\ , \\end{eqnarray} for which the power law attractors, $v_{1,2}$ are \\begin{eqnarray} v^{(P)}_{1,2}&=&\\frac{12n^2+9n+3\\pm \\sqrt{144n^2+120n^3-15n^2-30n-3}}{2(3+3n)} \\label{VPN}\\\\ v^{(Q)}_{1,2}&=&\\frac{4n^2+2n+1\\pm \\sqrt{16n^4-16n^2-10n-1}}{2n+2}\\ . \\label{VQN} \\end{eqnarray}  For $n=1$, all the above expressions agree with the values found earlier. It is interesting to note that $v^{(Q)}_{1,2}$ are imaginary for $n=1$, but are real for all $n>1$.  Of course, much remains to be done. We have not addressed solar system tests of these theories since this is a complicated analysis that is beyond the scope of our current paper. We have not focused on detailed comparisons between our models and the supernova data and it is possible that there are specific signatures of this new physics in such data. For example, the modified Friedmann equations arising from the theories presented here should directly provide information about the jerk parameter. We intend to consider such effects, not only in the present models, but in those proposed by us and other authors, in a separate paper focused on the connections of these models with observations. In this context, as with the case of the simple modifications introduced in~\\cite{Carroll:2003wy}, it may also be interesting to study more complicated functions of the curvature invariants we have considered."
    },
    "0410/astro-ph0410207_arXiv.txt": {
        "abstract": "Evolutionary population synthesis models for a wide range of metallicities, ages, star formation histories, initial mass functions, and Horizontal Branch morphologies, including blue morphologies at high metallicity, are computed. The model output comprises spectral energy distributions, colours, stellar M/L ratios, bolometric corrections, and near-infrared spectral line indices. The energetics of the post Main Sequence evolutionary phases are evaluated with the fuel consumption theorem. The impact on the models of the stellar evolutionary tracks (in particular with and without overshooting) is assessed. We find modest differences in synthetic broad-band colours as induced by the use of different tracks in our code (e.g., $\\Delta (V-K)\\sim 0.08~\\rm mag$; $\\Delta(B-V)\\sim 0.03~\\rm mag$). Noticeably, these differences are substantially smaller than the scatter among other models in the literature, even when the latter adopt the same evolutionary tracks. The models are calibrated with globular cluster data from the Milky Way for old ages, and the Magellanic Clouds plus the merger remnant galaxy NGC~7252, both for young ages of $\\sim 0.1-2\\;$Gyr, in a large wavelength range from the $U$-band to the $K$-band. Particular emphasis is put on the contribution from the Thermally-Pulsing Asymptotic Giant Branch phase. We show that this evolutionary phase is crucial for the modelling of young stellar populations by the direct comparison with observed spectral energy distributions of Magellanic Clouds clusters, which are characterised by relatively high fluxes both blueward and redward the $V$-band.  We find that the combination of the near-IR spectral indices $\\rm C_{2}$ and $\\rm H_{2}O$ can be used to determine the metallicity of $\\sim 1~\\rm Gyr$ stellar populations.  As an illustrative application, we re-analyze the spectral energy distributions of some of the high-$z$ galaxies ($2.4 \\lapprox z \\lapprox 2.9$) observed with the Spitzer Space Telescope by Yan et al.~(2004). Their high rest-frame near-IR fluxes are reproduced very well with the models including Thermally-Pulsing Asymptotic Giant Branch stars for ages in the range $\\sim 0.6-1.5~\\rm {Gyr}$, suggesting formation redshifts for these objects around $z\\sim 3-6$. ",
        "introduction": "The evolutionary population synthesis (EPS) is the technique to model the spectrophotometric properties of stellar populations, that uses the knowledge of stellar evolution. This approach was pioneered by B.~Tinsley in a series of fundamental papers, that provide the basic concepts still used in present-day computations. The models are used to determine ages, element abundances, stellar masses, stellar mass functions, etc., of those stellar populations that are not resolvable in single stars, like galaxies and extra-galactic globular clusters. Due to the ubiquitous astrophysical applications of EPS models, a rich literature has been developed so far (Bruzual~1983; Renzini \\& Buzzoni~1986; Chiosi, Bertelli and Bressan~1988; Buzzoni~1989; Charlot \\& Bruzual~1991; Bruzual \\& Charlot~1993; Worthey~1994; Vazdekis et al.~1996; Tantalo et al.~1996; Fioc \\& Rocca-Volmerange~1997; Bressan, Granato \\& Silva~1998; Maraston~1998; Leitherer et al.~1999; Brocato et al.~2000; Thomas, Maraston \\& Bender~2003; Thomas, Maraston \\& Korn~2004).  In the simplest flavour of an EPS model, called {\\it Simple Stellar Population} (hereafter SSP) by Renzini~(1981), it is assumed that all stars are coeval and share the same chemical composition. The advantage of dealing with SSPs is twofold.  First, SSPs can be compared directly with globular cluster (GC) data, since these are the ``simplest'' stellar populations in nature. This offers the advantage of {\\it calibrating} the SSPs with those GCs for which ages and element abundances are independently known, an approach introduced in the review by Renzini \\& Fusi Pecci (1988). This step is crucial to fix the parameters that are used to describe that part of the model ``input physics'' - convection, mass loss, mixing - that cannot be derived from first principles.  The calibrated models can be applied with more confidence to the study of extragalactic stellar populations. This step is taken in the models of Maraston~(1998) and in their extension presented here. Second, complex stellar systems which are made up by various stellar generations are modelled by convolving SSPs with the adopted star formation history (e.g. Tinsley~1972; Arimoto \\& Yoshii 1986; Rocca-Volmerange \\& Guiderdoni 1987; Vazdekis et al.~1996; Kodama \\& Arimoto~1997; Barbaro \\& Poggianti 1997; Bruzual \\& Charlot~2003). Therefore the deep knowledge of the building blocks of complex models is very important.  Two techniques are adopted to compute SSP models, which differ according to the integration variable adopted in the post-Main Sequence: isochrone synthesis and `fuel consumption based' algorithms. With the `isochrone synthesis' (e.g. Chiosi, Bertelli \\& Bressan~1988; Charlot \\& Bruzual~1991) the properties of a stellar population are calculated by integrating the contributions to the flux in the various passbands of all mass-bins along one isochrone, after assuming an Initial Mass Function (IMF). Usually isochrones are computed up to the end of the Early Asymptotic Giant Branch phase. Later stellar phases like the Thermally-Pulsing Asymptotic Giant Branch are added following individual recipes or are neglected.  In the `fuel consumption' approach (Renzini 1981; Renzini \\& Buzzoni 1986; Buzzoni~1989; Maraston~1998), the integration variable in Post Main Sequence is the so-called {\\it fuel}, that is the amount of hydrogen and/or helium that is consumed via nuclear burning during a given post Main Sequence phase. The fuel at a given age is computed on the stellar evolutionary track of the {\\it turnoff} mass (i.e., the mass completing the hydrogen-burning phase), thereby neglecting the dispersion of stellar masses in post main sequence. However since a mass difference of only few percent exists between the {\\it turnoff} mass and the mass at any other post-Main Sequence phase, this assumption can be made safely, as also shown by Charlot \\& Bruzual~(1991). The advantages of the fuel consumption approach are of two kinds. First, the fuel as integration variable is very stable since it is directly proportional to the contributions of the various phases to the total luminosity. This is very important in luminous, but short-lived evolutionary stages, e.g. the bright Red Giant Branch phase, where the evolutionary mass practically does not change. We note that the problem of the numerical instability on the RGB was early recognized by Tinsley \\&~Gunn (1976). Second, and more important, there are several relevant stellar phases (e.g. blue Horizontal Branch, Thermally Pulsing Asymptotic Giant Branch, very hot old stars etc.)  whose theoretical modeling is uncertain because of mass loss and for which complete stellar tracks are not available. The fuel consumption provides useful analytical relations that link the Main Sequence to the post Main Sequence evolution, by means of which one can include into the synthesis the energy contributions of these uncertain phases using, e.g. observations. The `isochrone synthesis' technique is used in all models in the literature, with the exception of the models by Buzzoni~(1989), Maraston~(1998) and those presented here, that adopt the fuel consumption theorem.  Besides the method used to compute them, evolutionary population synthesis models keep the uncertainties inherent in the stellar evolutionary tracks and in the spectral transformations. Charlot, Worthey \\& Bressan~(1996) investigated both classes of uncertainties by comparing their EPS models. The analysis is very illustrative for the three considered EPS, however does not yield information about the sole impact of the stellar tracks. Here we make a different exercise, because we can use the same EPS code and just vary the model ingredients. In this way we can isolate their impact on the final model. To quantify the model uncertainties is indeed very important since the cosmological inferences that are derived on the basis of galaxy ages and metallicities rely ultimately on the stellar population models.  Maraston (1998) presents a fuel-consumption-based code for evolutionary population synthesis, and SSP models for solar metallicity, and ages from 30 Myr to 15 Gyr. Distinct features of that work are: \\begin{enumerate} \\item the extension of the fuel consumption theorem to compute models for young and intermediate-range ages; \\item the inclusion of a well-calibrated semi-empirical Thermally-pulsing Asymptotic Giant Branch phase, and the computation of realistic colors of intermediate-age ($t\\lapprox~2~\\rm Gyr$) stellar populations \\item the modular structure of the code, that allows experiments with the EPS ingredients. \\end{enumerate} Subsequentely the code has been updated to the computations of SSP models covering a wide range in metallicities and the model output have been extended to e.g. the spectral energy distributions, spectral indices, redshift evolution, with several applications being already published (Maraston \\& Thomas 2000; Saglia et al.~2000; Maraston et al. 2001; Maraston et al.~2003; Thomas, Maraston \\& Bender ~2003; Thomas, Maraston \\& Korn 2004; Ferraro et al.~2004). This work is devoted to discuss comprehensively the overall EPS models and ingredients, in particular, the inclusion of the TP-AGB phase in the synthetic spectral energy distributions.  The paper is organized as follows. In Section~2 the properties of the evolutionary population synthesis code are recalled. The model ingredients, i.e. the fuel consumptions, the temperature distributions and the model atmospheres, for the various metallicities and ages, are described in detail in Section~3. In particular, the recipes for Horizontal Branch and Thermally Pulsing Aymptotic Giant Branch are presented in Section 3.4. The various model output are discussed in Section~4, where the comparisons with observational data and with models from the literatures are also presented. Section~5 deals with the model uncertainties, while Section~6 presents an high-redshift application of the model SEDs in which the TP-AGB has a primary importance. Finally a summary of the main results is given in Section~7. ",
        "conclusions": "Evolutionary population synthesis models with various ages, metallicities, star formation histories, Horizontal Branch morphologies, are presented. The evolutionary code is based on the fuel consumption theorem (Renzini \\& Buzzoni~1986) for the evaluation of the energetics of the post Main Sequence evolutionary phases. The code was introduced in our previous work (Maraston~1998) in which stellar population models complete in all main stellar evolutionary phases, but only for solar metallicity were computed. The present paper provides the extension of the modelling to a wider range of stellar population parameters. The most important features of the models are: i) the inclusion of the energetics and the spectra of TP-AGB stars and their calibrations with observations; ii) the allowance for various Horizontal Branch morphologies, in particular blue morphologies at high metallicities. We will come back to both points below.  In the first part of the paper we perform a comprehensive analysis of the model ingredients, focusing in particular on: i) the metallicity effects on energetics and temperatures; ii) the impact of different stellar evolutionary tracks. The most relevant results can be summarized as follows.  The energetics of stellar populations are higher at lower metallicity because of the greater hydrogen abundance. The most affected stellar phases are the main sequence and the helium burning phase. The differences become smaller at ages later than the RGB phase-transition ($t\\gapprox~1~\\rm Gyr$) when the RGB phase becomes the most important contributor to the total bolometric. This is due to the RGB fuel consumption being not affected dramatically by the chemical abundance, a result consistent with previous findings (Sweigart, Greggio \\& Renzini~1989).  The impact of stellar evolutionary tracks is evaluated by comparing energetics and temperatures of two different stellar models, namely the set of Frascati tracks by Cassisi and collaborators, that are classical tracks without overshooting, with those by the Padova school, that include the effects of overhsooting. We find that the mass- luminosity relations for MS stars are rather similar among the explored tracks, which is explained by the modest amount of overshooting included in recent computations. Instead, major differences are found in post-MS phases, which involve both the energetics and the temperatures of the Red Giant Branch. In the Padova tracks the onset and development of the RGB is delayed with respect to the Frascati tracks, an effect of overshooting. For example, a $\\sim 0.8~\\rm Gyr$ stellar population model based on the Padova tracks has $\\sim 30$ per cent less light than one based on the Frascati tracks, a result influencing the integrated output, especially the $M/L$. It should be noticed that a recent comparison of the onset age of the RGB with what is observed in MC GCs favours the earlier RGB development of the Frascati tracks (Ferraro et al.~2004).  The temperatures of the RGB, especially towards the tip are cooler in the Padova tracks, at solar metallicity and above, a feature also found in the Yale tracks of Yi et al.~(2003). The RGB temperature is the ingredient having the largest influence on the SSP output. For example, the integrated $V-K$ colours of old, metal-rich, stellar populations change by $\\Delta V-K\\sim 0.08~\\rm mag$. The RGB temperatures are connected to the mixing-length and the underlying stellar models. The calibration of these effects is difficult since the comparisons with data require the adoption of temperature-colors transformations, that are uncertain as well. In spite of this, it would be highly valuable to understand the origin of such discrepancies at least theoretically.  Relevant to evolutionary population synthesis models is the following result. The differences induced by the use of various stellar tracks in our code are substantially smaller than the scatter among stellar population models in the literature that adopt the same stellar tracks.  \\medskip In the second part of the paper we describe in detail the recipes for the stellar phases that are affected by mass-loss and therefore not described sufficiently by stellar evolutionary tracks. These are the Horizontal Branch (HB) in old stellar populations and the Thermally-Pulsing Asymptotic Giant Branch (TP-AGB) in stellar populations with ages between 0.2 and 2 Gyr. For the HB we apply mass-loss to the RGB tracks such that we obtain a good match to stellar population spectrophotometric indicators that are sensitive to the HB morphology (e.g. the Balmer lines). We therefore calibrate the amount of mass-loss with MW GCs of known HB morphology. In order to account for the observed scatter in the HB properties at a fixed metallicity, we provide the models with a couple of choices for the HB morphology. In particular, we also compute blue HB at high $Z$ and show how much blue indicators like e.g. the Balmer lines and the $U,B,V$ magnitudes are affected. The effects of BHBs at high $Z$ are quantified in a form that can be easily used also when other evolutionary population synthesis are adopted.  A result interesting for extra-galactic GC studies is that the D-4000 \\AA~break, unlike the Balmer lines is insensitive not only to age but also to the Horizontal Branch morphology at low metallicity ($\\zh \\lapprox -1$). Therefore this colour can be used as metallicity indicator in the metal-poor regime.  The energetic of the TP-AGB phase was calibrated with Magellanic Clouds GCs data in Maraston (1998). It was shown that this phase is the dominant one in stellar populations with ages between 0.3 and 2 Gyr, providing 40 per cent of the bolometric and up to 80 per cent of the luminosity in the near-IR. Here we make a step forward by including the spectra of TP-AGB stars in the synthetic spectral energy distributions (SEDs) of the stellar population models, using available empirical spectra of TP-AGB C-rich and O-rich stars. We compare the resulting model SEDs with the observed SEDs of a sample of MC GCs with data from $U$ to $K$. The models provide a very good fit to the data in the whole age range relevant to the TP-AGB, a success that is entirely due to the proper inclusion of the phase in the models. We further compute the integrated spectral indices $\\rm C_{2}, H_{2}O, CO$ in the near-IR ($1.7~\\lapprox \\lambda/\\rm \\mu m \\lapprox 2.4$), that are effective tracers of C and O stars. We find that the combination of $\\rm C_{2}$ and $\\rm H_{2}O$ can be used to determine the metallicity of $\\sim 1~\\rm Gyr$ stellar populations, independently of the star formation history.  The model SEDs presented here can be applied to the analysis of high-$z$ galaxies extending the concept of TP-AGB as age indicator for $\\sim~1~{\\rm Gyr}$ stellar populations to primeval objects in the early universe. As an illustrative application we re-analyze some high-$z$ ($2.4 \\lapprox z\\lapprox 2.9$) galaxies recently observed with the Spitzer Space Telescope (Yan et al.~2004). The distinctive features of these objects is to have strong fluxes in both the rest-frames blue and near-IR relative to the flux in the $V$-band. This is what characterizes TP-AGB dominated stellar populations. We show that the data can be explained by stellar populations with ages in the range 0.6 to 1.5 Gyr. These ages are comfortable given the age of the universe at those redshifts ($\\sim~2~\\rm Gyr$) and imply formation redshifts between 3 and 6.  \\medskip Models are available at www-astro.physics.ox.ac.uk/$\\sim\\,$maraston."
    },
    "0410/astro-ph0410561_arXiv.txt": {
        "abstract": "One of the two breakout reactions from the hot CNO (HCNO) cycle is ${}^{15}\\textrm{O}(\\alpha,\\gamma){}^{19}\\textrm{Ne}$, which at low temperatures depends strongly on the resonance strength of the 4.033 MeV state in ${}^{19}\\textrm{Ne}$. An experimental upper limit has been placed on its strength, but the lower limit on the resonance strength and thereby the astrophysical reaction rate is unconstrained experimentally. However, this breakout reaction is crucial to the thermonuclear runaway which causes type I X-ray bursts on accreting neutron stars. In this paper we exploit astronomical observations in an attempt to constrain the relevant nuclear physics and deduce a lower limit on the reaction rate. Our sensitivity study implies that if the rate were sufficiently small, accreting material would burn stably without bursts. The existence of type I X-ray bursts and superbursts consequently suggests a lower limit on the ${}^{15}\\textrm{O}(\\alpha,\\gamma){}^{19}\\textrm{Ne}$ reaction rate at low temperatures. ",
        "introduction": "\\label{sec:introduction} Type I X-ray bursts were discovered independently by \\cite{Grindlay76a} and \\cite{Belian76} and occur in matter-transfering binary systems in which a neutron star accretes hydrogen and helium from an unevolved companion star \\citep{Woosley76,Joss77,Joss78b,Taam80,Ayasli82, Lewin95, Bildsten98c}. The transferred matter is heated to $1-2\\times 10^8$ K while it slowly descends into the neutron star atmosphere as freshly infalling material continuously piles on top of it. At these temperatures hydrogen burns stably into helium at a constant and temperature-independent rate given by the half lives of two nuclei participating in the $\\beta$-limited HCNO cycle \\citep{Wallace81}. \\begin{equation}\\label{eq:hcno} \\frac{dX_H}{dt}=-\\frac{\\lambda_{{}^{14}O}\\lambda_{{}^{15}O}}{\\lambda_{{}^{14}O}+\\lambda_{{}^{15}O}} X_{HCNO}\\,, \\end{equation} where the decay constants $\\lambda=\\ln 2/T_{1/2}$ are given by the half lives of ${}^{14}\\textrm{O}$ and ${}^{15}\\textrm{O}$, and $X_{HCNO}$ is the total mass fraction of these two nuclei. The triple-$\\alpha$ reaction concurrently increases the mass fraction of HCNO material while depleting helium \\citep{Bildsten98c}. The mass fraction of HCNO material is important because it regulates the rate at which hydrogen burns to helium, determining the initial composition and the characteristics of the burst \\citep{Fujimoto81,Fushiki87,Cumming03a,Narayan03}.  The mass fraction of HCNO material can decrease through processing into heavier isotopes via the ${}^{15}\\textrm{O}(\\alpha,\\gamma)$${}^{19}\\textrm{Ne}$ and $^{18}$Ne($\\alpha,p)^{21}$Na breakout reactions \\citep{Wallace81,Wiescher99}. This can happen either during the burst, when peak temperatures exceed $10^9$ K, or during the quiescent, non-bursting phase. At temperatures characteristic of the quiescent phase ($T \\lesssim 2\\times 10^8$ K) only the ${}^{15}\\textrm{O}(\\alpha,\\gamma)$${}^{19}\\textrm{Ne}$ reaction allows significant leakage out of the HCNO cycle \\citep{Hahn96}. In this case the leak of HCNO material into heavier isotopes depends on the ratio between the $\\alpha$ capture rate of ${}^{15}\\textrm{O}$ and its decay rate \\begin{equation}\\label{eq:leak} R=\\frac{N_A<\\sigma v>\\rho X_\\alpha/A_\\alpha}{\\lambda_{{}^{15}O}}\\,, \\end{equation} where $<\\sigma v>$ is the thermally averaged $^{15}$O($\\alpha,\\gamma)$ reaction rate per particle pair, $N_A$ is the Avogadro constant, $\\rho$ is the stellar mass density, $X_\\alpha$ is the ${}^4\\textrm{He}$ mass fraction and $A_\\alpha$ its atomic mass. %  What happens therefore depends on the relatively unknown ${}^{15}\\textrm{O}(\\alpha,\\gamma){}^{19}\\textrm{Ne}$ reaction rate: If the rate is comparable to the \\cite{Caughlan88} rate, then the leak will be important in reducing the HCNO mass fraction of Eq.~[\\ref{eq:hcno}] and therefore in decreasing the rate at which hydrogen burns into helium as matter sinks. With long burst recurrence times of thousands of seconds or even days \\citep{Strohmayer03}, even small reaction rates can be significant, allowing hydrogen to survive to great depths. In addition, a rate of this magnitude would deplete ${}^{15}\\textrm{O}$ above the ignition point during the burst runaway, ensuring that little or no ${}^{15}\\textrm{O}$ survives for the next burst and requiring the accretion of a fresh layer of material.  However, the $\\alpha$ width of the dominant 4.033 MeV state in ${}^{19}\\textrm{Ne}$ is not well known, so the rate could be lower. If the rate is lower, there will not be any significant leak out of the HCNO cycle, whence hydrogen would burn to helium at a faster rate, depleting hydrogen from the ignition zone. In addition, a lower rate would prevent or weaken the breakout via ${}^{15}\\textrm{O}(\\alpha,\\gamma)$${}^{19}\\textrm{Ne}$ ensuring less energy generation during the initial stages of the burst \\citep{Taam79,Ayasli82,Hanawa84,Wiescher99}. The survival of a significant fraction of ${}^{15}\\textrm{O}$ in the upper layers would increase the hydrogen burning prior to subsequent bursts, thus altering the composition at the time of the runaway.  In this paper we show that a sufficiently small $^{15}$O($\\alpha,\\gamma)$ rate can even prevent subsequent runaways, causing hydrogen and helium to burn stably. Stable burning creates copious amounts of ${}^{12}\\textrm{C}$ and other $A\\lesssim 20$ nuclei which have implications for superbursts. Astronomical observations may lead to better constraints on this rate and motivate future measurements of the ${}^{15}\\textrm{O}(\\alpha,\\gamma){}^{19}\\textrm{Ne}$ reaction rate. Such measurements would help refine current X-ray burst and superburst models. ",
        "conclusions": "\\label{sec:conclusion} We have demonstrated the significance of the ${}^{15}\\textrm{O}(\\alpha,\\gamma){}^{19}\\textrm{Ne}$ reaction to X-ray burst behavior and suggested a lower limit on the rate which could and should be pursued experimentally.  Long intermittent periods of stable burning, which may last for months, have been observed from burst sources; a lower value of the ${}^{15}\\textrm{O}(\\alpha,\\gamma){}^{19}\\textrm{Ne}$ reaction rate will help stabilize stable burning. The reaction rate determines the value of the critical accretion rate.  In addition we have shown a way to generate the necessary carbon for the superburst theory to be valid. We have also shown that several different kinds of $A\\lesssim 20$ nuclei are generated; applying this to superburst calculations requires proper knowledge of the fusion rates between these nuclei."
    },
    "0410/astro-ph0410082_arXiv.txt": {
        "abstract": "New constraints on the large-scale wind structures responsible for discrete absorption components are obtained from far-ultraviolet time-series observations of O supergiants in the Large Magellanic Cloud. ",
        "introduction": "Discrete absorption components (DACs) are optical depth enhancements in the absorption troughs of P~Cygni profiles, which (a) occur very commonly in wind profiles of OB-type stars; (b) accelerate slowly compared with the expected wind flow; and (c) often recur on time scales related to stellar rotation. Since the radial optical depth at any position $v \\equiv v(r)$ in an expanding wind is given by the Sobolev approximation \\begin{displaymath} \\tau_{\\rm rad}(v)\\propto ~ q_i ~\\rho~ \\left( {{\\rm d}v \\over {\\rm d}r } \\right)^{-1}, \\end{displaymath} such enhancements can in general be produced by {\\em increasing} the local ion fraction, $q_i$\\,; {\\em increasing} the local wind density, $\\rho$\\,; {\\em flattening} the local velocity gradient, ${\\rm d}v / {\\rm d}r$\\,; or by some combination of these factors, which are inter-related. The far-ultraviolet region of the spectrum provides access to many different resonance lines, and consequently time-series observations with {\\it FUSE\\/} provide the diagnostic capability required to disentangle these factors. ",
        "conclusions": ""
    },
    "0410/astro-ph0410427_arXiv.txt": {
        "abstract": "Dwarf nova outbursts often show coherent ($Q\\sim10^4-10^6$) sinusoidal oscillations with the largest pulsed fraction in the extreme ultraviolet. Called dwarf nova oscillations (DNOs), they have periods of $P\\approx3-40\\textrm{ s}$ and scale with luminosity as $P\\propto L^{-\\beta}$ with $\\beta\\approx0.1-0.2$. We propose that DNOs may be produced by nonradial oscillations in a thin hydrostatic layer of freshly accreted material, the ``spreading layer'' (SL), at the white dwarf (WD) equator. This would naturally explain a number of key properties of DNOs, including their frequency range, sinusoidal nature, sensitivity to accretion rate, and why they are only seen during outburst. In support of this hypothesis we construct a simple model that treats the SL as a cavity containing shallow surface waves, each with the same radial structure, but split into three different modes denoted by their azimuthal wavenumber, $m$. The $m=0$ latitudinally propagating mode best matches the periods and scalings associated with most DNOs, and DNOs with periods shorter than the WD Keplerian period are explained by the $m=-1$ prograde mode. We also predict a third set of oscillations, produced by the $m=1$ retrograde mode, and show its expected dependence on accretion rate. ",
        "introduction": "Rapid oscillations are often seen in the optical, soft X-ray, and extreme ultraviolet (EUV) flux of cataclysmic variables (CVs) during dwarf nova (DN) outbursts. These oscillations have periods in the range of $P\\approx3-40\\textrm{ s}$, show a monotonic relationship with EUV luminosity (Mauche 1996) and thus accretion rate, $\\dot{M}$, and are fairly coherent ($Q=|dP/dt|^{-1}\\sim10^4-10^6$). They are called dwarf nova oscillations (DNOs; Patterson 1981) in contrast to the much less coherent ($Q\\sim5-20$) quasi-periodic oscillations (QPOs) also seen from CVs and the highly coherent oscillations ($Q\\sim10^{12}$) that reflect the spin of a magnetic primary (e.g., DQ Her). Since their discovery (Warner \\& Robinson 1972), DNOs have been seen in $\\sim50$ systems (Warner 2004). The DNO periods roughly track the surface Keplerian period of the accreting white dwarf (WD) [$P_{\\rm K}= 2\\pi/(GM/R^3)^{1/2}$; see Fig. 19 of Patterson 1981, or Knigge et al. 1998], so that it is generally believed that the oscillations are created near or on the WD surface.  A number of explanations have been proposed for DNOs, but none have quantitatively explained the majority of their key features. The characteristic period of the DNOs, $P\\sim P_{\\rm K}$, implies a pulsational, rotational, or even perhaps inner accretion disk origin. Since the period changes on the same timescale as accretion, it is improbable that the entire WD takes part in creating DNOs, ruling out global {\\it g}-modes or the WD spin modulated by hot spots. This led Paczy\\'{n}ski (1978) to suggest that if accretion torques were causing the period drifts then only a small amount of material is involved in making DNOs ($\\lesssim10^{24}\\textrm{ g}$). Papaloizou \\& Pringle (1978) studied rotationally modified global {\\it g}-modes and {\\it r}-modes on WDs. Since truly global modes are not sensitive to the instantaneous accretion rate, they proposed that DNOs are high radial-order modes that have large amplitudes near the WD surface. Their work explains many of the properties of DNOs, but some questions still remain, including: (1) what physical mechanism favors modes that are strongly concentrated in the outer layers of the star, (2) why the radiating area associated with DNOs is such a small fraction of the total WD surface area ($\\sim10^{-4}$ to $10^{-2}$ as inferred from the EUV emission of SS Cyg during DN outbursts; C\\'{o}rdova et al. 1980; Mauche \\& Robinson 2001, hereafter MR01; Mauche 2004), and (3) why the modes are only excited during outbursts. These failings led others to suggest alternatives, including magnetic accretion onto a slipping belt (Warner \\& Woudt 2002) and non-axisymmetric bulges on the inner part of accretion disks (Popham 1999).  Inogamov \\& Sunyaev (1999) presented a new way of studying disk accretion close to an accreting neutron star (NS), which Piro \\& Bildsten (2004; hereafter PB04) extended to the case of WDs. Instead of using a boundary layer model, they follow the latitudinal flow of accreted material over the stellar surface, starting at the equator and streaming up toward the pole, using a spreading layer (SL) model. In the case of WDs the covering fraction (which is equivalent to the spreading angle) is found to be small, $f\\sim10^{-3}$ to $10^{-1}$ (for $\\dot{M}\\approx10^{16}-10^{18}\\textrm{ g s}^{-1}$), so that the size of the freshly accreted layer is set by the thickness of the accretion disk.  The material in the SL is much hotter in temperature and lower in density than the underlying WD. This contrast allows waves in the SL to travel freely, unencumbered by the material below. We propose that DNOs are shallow surface waves in the layer of recently accreted material confined to the WD equator. We argue that the $m=0$ mode provides the period that is most often identified as a DNO, and we show that the $m=-1$ prograde mode explains the occurrence of some of the higher frequency DNOs, including the extra periods seen from SS Cyg and VW Hyi (MR01; Woudt \\& Warner 2002). We also discuss the $m=1$ retrograde mode and whether it also corresponds to an observed frequency.  In \\S 2 we describe a simple model to estimate the shallow surface wave periods and show how they scale with $\\dot{M}$ and the WD mass and radius. We then compare these periods to observed DNOs from CVs in \\S 3. We conclude in \\S 4 with a summary of the DNO properties successfully understood using our model along with a discussion of problems that are still left unresolved. ",
        "conclusions": "We propose that DNOs in outbursting CVs are nonradial oscillations in a hot layer of freshly accreted material near the WD equator. A large number of DNO properties are then simply understood: (1) the highly sinusoidal nature of the oscillations is consistent with nonradial oscillations, (2) the periods can change on the timescale of accretion because there is little mass in the layer ($\\lesssim10^{21}\\textrm{ g}$; PB04), (3) the periods vary inversely with $\\dot{M}$ because they have the temperature scaling of shallow surface waves, (4) the covering fraction is naturally small for the SL, (5) the DNOs are only seen during DN outbursts because this is the only time when an optically thick layer of material can build up at the equator, and (6) the largest pulsed amplitude is in the EUV, consistent with the SL temperature.  In support of our hypothesis, we presented a simple, phenomenologically motivated model that quantitatively explains many of the DNO's features. We compared our model with the overall population of DNOs, both as a function of WD mass and accretion rate. In each case, it is necessary to consider both the $m=0$ mode and the $m=-1$ (prograde) mode to explain the range of periods observed and the different $P$-$\\dot{M}$ scalings. The majority of DNOs are consistent with the period of the $m=0$ mode, and this may be related to its latitudinal propagation. The speed of shallow surface waves is $\\approx gh$ so that they slow down as the scale height decreases. This steers an initially azimuthally propagating mode to instead travel perpendicular to the SL edge, similar to ocean waves at the beach, and may explain why $m=0$ modes are favored over $m\\neq0$ modes. In the 1996 DN outburst of SS Cyg (MR01), the $m=-1$ mode is only seen near the outburst peak. This may indicate that a wide enough region without differential rotation is only present when there is sufficient spreading at high $\\dot{M}$.  There are still many difficulties that must be addressed about this idea for explaining DNOs. From the SL model (PB04) we borrowed the concept of hot material in hydrostatic balance covering a small fractional area of the WD, but we ignored important details of these calculations, such as differential rotation and the change of scale height with latitude. The data do not require such additional complications, so we refrain from including them for now. In a more sophisticated model, it would probably still be true that $k\\propto1/(fR)$, but an eigenvalue calculation would determine the constant of proportionality. A full calculation may also help explain the variety of power laws, $\\beta\\equiv-d\\log P/d\\log L$, for various CVs, since our model only predicts one unique power law index for the $m=0$ mode. The WD spin (which has so far been neglected) may be an important differentiating variable as well.  Another important property of WDs that would affect the modes is a magnetic field. A strong field inhibits shearing between the SL and WD and modifies the frequency of shallow surface waves. It is therefore interesting that no intermediate polars (IPs) have shown DNOs or lpDNOs. Even LS Peg and V795 Her, neither of which are IPs, but both showing polarization modulations (Rodr\\'{i}quez-Gil et al. 2001, 2002) indicative of a reasonably strong magnetic field, are without DNOs or lpDNOs. Further studies should also work toward an understanding of the excitation mechanism for the modes, which would help explain the high coherence typical of DNOs. Material deposited at the WD equator spreads fairly quickly, $10-100\\textrm{ s}$ (PB04), so that this material must be moving through the oscillating region on timescales of order the mode period. In light of the many strengths of the SL mode explanation of DNOs we do not abandon it because of this problem, but this is definitely a weakness that puts limits any proposed excitation mechanism.  Our explanation of DNOs raises interesting questions about the relationship between oscillations originating from accreting compact objects. Mauche (2002b) showed that there is a strong correlation in the high to low oscillation frequency ratio of WDs, NSs, and black holes (BHs). Using this picture, DNOs are associated with the kilohertz QPOs of low mass X-ray binaries. Suggestively, the Fourier frequency resolved spectroscopy of NSs (Gilfanov, Revnivtsev \\& Molkov 2003) imply that both the normal branch oscillations and the kilohertz QPOs are created in the NS boundary layer, strengthening the possibility that our SL mode model may also apply in this case. For a typical NS mass and radius our model results in frequencies in the range expected for kHz QPOs, but it does not explain their complicated $P$-$L$ relation (the ``parallel tracks;'' van der Klis 2000). Interestingly, DNOs may also show the parallel tracks phenomenon, as seen in three observations of SS Cyg (Mauche 2002a), once again supporting the correlation. On the other hand, continuing the analogy to WD and BH oscillations is problematic because in the case of BHs there is no surface for nonradial oscillations.  We thank Phil Arras and Christopher Mauche for many helpful discussions. We also thank the referee for asking the right questions, which inspired a significantly revised manuscript. This work was supported by the National Science Foundation under grants PHY99-07949 and AST02-05956, and by the Joint Institute for Nuclear Astrophysics through NSF grant PHY02-16783."
    },
    "0410/astro-ph0410611_arXiv.txt": {
        "abstract": "We present a first deep colour-magnitude diagram of the putative  central region (0.5$\\arcdeg \\times$ 0.5$\\arcdeg$) of the Canis Major stellar over-density $(l,b)=(240,-8)$ found by Martin et al. (2004), which has been proposed as the remnant of a dwarf  satellite accreted onto the Milky Way on a near-equatorial orbit. We find a narrow (in apparent magnitude) main-sequence, extending 6 magnitudes below the turn-off to our limiting magnitude of B$\\sim$ 24.5\\,mag. This main sequence has very high constrast ($>$3) with respect to the thin/thick disk/halo background; its narrowness at brighter magnitudes clearly implies the presence of a distinct and possibly still bound stellar system. We derived the line-of-sight size (r$_{1/2}$) of this system based on the B-band width of the lower main sequence, obtaining  0.94 $\\pm$ 0.18 (random) $\\pm$ 0.18 (systematic)\\,kpc. That size matches a model prediction for the main body of the parent galaxy of the Monoceros tidal stream. The high density contrast and limited spatial extent in the radial direction are very hard to reconcile with the alternative explanation put forward to explain the Canis Major stellar-overdensity: a flared or warped Galactic disk viewed in projection (Momamy et al. 2004). We also derived a central surface brightness of $\\mu_{V,0}= 23.3 \\pm 0.1$\\,mag arcsec$^{-2}$  and an absolute magnitude of  $M_{V}=-14.5 \\pm 0.1$\\,mag. These values place the Canis Major object in the category of dwarf galaxy in the the $L_{V}$--size and  $M_V -\\mu_{\\rm V}$ planes for such objects.  However, like the Sagittarius dwarf, it is an outlier in  the $[\\rm Fe/H] -M_V$ plane in the sense that it is too metal rich for its estimated absolute magnitude. This suggests that the main mechanism driving its  recent and current star formation history  (possibly tidal stripping) is different to that of isolated dwarfs. ",
        "introduction": "The Milky Way offers a unique laboratory for testing the hierarchical galaxy formation scenario through direct evidence of  past  merging and tidal disruption events, which result in extensive stellar streams or large scale substructures in the Galactic halo or even in the disk (Navarro 2004, and references therein). Resolving such tidal streams into stars and measuring  the phase-space coordinates for these stars,  provides a  fossil dynamical accretion record of unparalleled accuracy. These can be directly compared with N-body simulations in order to reconstruct their dynamical history and their impact on the Milky Way at the present and recent epochs.    Recently, the Sloan Digitized Sky Survey (SDSS) team reported the discovery of a coherent giant stellar structure  at low galactic latitudes (Newberg et al. 2002; Yanny et al. 2003), which appears to form a ring around the Milky Way. Since then, there has been a tremendous amount of follow-up observational effort probing its structure and kinematics in order to understand its origin (see Majewski 2004 and references therein). N-body simulations (Martin et al. 2004a; Pe\\~narrubia et al. 2004) shows that this ring structure can be naturally attributed to  the tidal stream of a dwarf galaxy  (named the Monoceros tidal stream, or the Galactic Anticenter tidal stream). If this stellar overdensity ring is a tidal tail feature, it must have had a ``parent'' galaxy, which may or may not be completely disrupted by now. The existence and location of such a parent galaxy is still controversial; the best candidate is the Canis Major (CMa) dwarf galaxy candidate, a strong, ellipsoidal overdensity of red giant stars, discovered in the namesake constellation by Martin et al. (2004a) from an analysis of the 2MASS survey. Bellazzini et al. (2004) presented a color-magnitude diagram (CMD) in the surroundings of the candidate CMa dwarf, concluding that the system is situated at 8$\\pm$ 1 kpc from the Sun and that it is composed of a metal-rich, intermediate-age stellar population.  As CMa is only $\\sim 8$ degrees from the Galactic plane, Momany et al. (2004) suggested an alternative interpretation, namely that this over-density  is  the signature of the Galactic warp in this direction of the sky, and not a distinct satellite. However, Martin et al. (2004b) reported a narrow distribution of radial velocities for the center of this structure,  and argued on this basis for the interpretation of CMa as an  accreted dwarf satellite remnant, in an orbit near the Galactic plane.  To help discriminate between these two hypothesis -- accreted dwarf galaxy, or stellar warp -- , we present here deep broad-band photometry of the CMa center. As our results confirm and greatly strengthen the case that the stellar over-density towards CMa is part of a distinct dwarf galaxy, we refer to it as 'CMa dwarf' throughout. ",
        "conclusions": "From our data we have established that the putative center of CMa is a very high contrast density feature that as has very narrow radial extent, $r_{1/2}/R_{G.C.} < 0.1$. These results strongly support the interpretation of a distinct, possibly still bound stellar system whose  properties are consistent with  those expected for the remnant of a partially disrupted dwarf satellite. In turn, the high density contrast and the limited line-of-sight extent of the CMa would be very difficult to be reconciled (if at all) with a flared or warped  Galactic disk viewed in projection (Momamy et al. 2004). Although the obtained properties for this satellite (Table 2) in this study are subject to considerable uncertainties (stellar population, Galactic contamination, radial profile), they are in agreement with those of the known Milky Way dSph satellites (Mateo 1998) and with the known size-luminosity relation followed by Milky Way dwarf spheroidals.  At a distance of 8 kpc, CMa is the closest dwarf galaxy known. It may not only be the parent of the Galactic low-latitude stellar stream, but also a unique laboratory for testing galaxy evolution theories. In particular, our CMD shows the availability  for first time of thousands of MS stars of different ages in a magnitude range suited to high resolution spectroscopy-based abundance studies using  4-meter class telescopes. This will allow one to study the SFH of a dwarf galaxy with unprecedented  spectral resolution, and is therefore a top candidate for chemo-dynamical modelling of dwarf galaxy evolution in the Galaxy."
    },
    "0410/hep-ph0410330_arXiv.txt": {
        "abstract": "s{We review recent developments for the description of nonequilibrium quantum fields, continuing hep-ph/0302210.\\cite{SEWM02}} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410119_arXiv.txt": {
        "abstract": "{Extended CCD monitoring of \\object{RR Gem} revealed that it is a Blazhko type RRab star with the shortest Blazhko period ($7\\fd23$) and smallest modulation amplitude ($\\Delta M_{max}<0.1$ mag) currently known. The short period of the modulation cycle enabled us to obtain complete phase coverage of the pulsation at each phase of the modulation. This is the first multicolour observation of a Blazhko star which is extended enough to define accurate mean magnitudes and colours of the variable at different Blazhko phases. Small, but real, changes in the intensity mean colours at different Blazhko phases have been detected. The Fourier analysis of the light curves shows that, in spite of the mmag and smaller order of the amplitudes, the triplet structure is noticeable up to about the 14th harmonic. The modulation is concentrated to a very narrow, 0.2 phase range of the pulsation, centred on the supposed onset of the H emission during rising light. These observational results raise further complications for theoretical explanation of the long known but poorly understood Blazhko phenomenon. ",
        "introduction": "Most RR~Lyrae stars repeat their light curves with remarkable regularity. About 30\\% of the known galactic RR~Lyrae stars, however, display cyclic modulation in the shape and amplitude of the light curve over tens to thousands of pulsational cycles \\citep{smith}. The phenomenon is often called the Blazhko effect since \\citet{blazhko} was the first to recognize that the light maximum of an RR~Lyrae star (namely \\object{RW~Dra}) showed phase modulation on a long time-scale (around 42 days).  In several well-studied RR~Lyrae stars, the amplitude and the period of the Blazhko effect has been observed to be variable on a time-scale of years \\citep{szeidl,smith99,lacluyze}. In this respect \\object{RR~Geminorum} provides a very interesting case. \\citet{detre70} remarked that his unpublished photographic observations from the 1930s demonstrated the Blazhko effect in \\object{RR~Gem}, but (in his words) it `ceased to exist since about 1940.' That the  Blazhko effect might be temporary in nature may be crucial in understanding the physical background of the phenomenon. Therefore we decided to begin an intensive photometric investigation of the star.  \\object{RR~Gem} ($\\alpha=07^h 21^m 33\\fs53, \\delta=+30\\degr 52\\arcmin 59\\farcs5$, J2000) was discovered by L. P. Ceraski \\citep{ceraski} on photographic plates taken by Blazhko. The first visual observations of the star was made by \\citet{graff} and he was able to determine the pulsation period (0\\fd397). Since that time a great number of visual observations have been published, but there are only a few published photoelectric observations available (e.g., \\citealp{fitch}; \\citealp{liu89}).  The absolute parameters of \\object{RR~Gem} were determined by \\citet{liu90} from a Baade-Wesselink analysis. It belongs to the metal rich ([Fe/H]$=-0.14$~dex as compiled in \\citet{jk96}) disk population group of RR~Lyrae stars \\citep{layden}.  Most of the papers dealing with Blazhko stars list the possible theoretical explanations of the phenomena and also the different observational arguments (see e.g. \\citealp{lacluyze,macho,chadid}). Neither of the two groups of possible models  a) the magnetic oblique rotator (\\citealp{shiba,cousens})  b) resonant coupling of nonradial mode(s) to the dominant radial mode \\citep{nowak1}  \\noindent{can fully explain the large variety of observational results of Blazhko behaviour. It is important to note, however, that both of the models connect the Blazhko period to the rotation of the star.}  In this first part of a series of two papers on \\object{RR~Gem}, we examine the photometric behaviour of the star exclusively using recent CCD observations. A reexamination of all the previous observations in order to reveal the long term behaviour of the modulation will be the subject of the second paper.  \\begin{figure}[t] \\centering \\includegraphics[width=8.8cm]{rrgfold.eps} \\caption{Folded $V$ light curve of \\object{RR~Gem}. The scatter of the light curve is caused by Blazhko modulation and not by observational inaccuracies. The amplitude of the modulation is very small, it does not reach 0.1~mag even in maximum light, and the shape of the modulation points to a very regular behaviour. In this, and in the other figures, phase 0.5 is set to the middle of the ascending branch defined as the phase where the $V$ flux is equal to its time averaged value.} \\label{rrgfold} \\end{figure} ",
        "conclusions": "The extended CCD observations of \\object{RR~Gem} exposed some previously unknown or unexamined properties of the Blazhko phenomenon. The new results concerning the properties of the modulation of \\object{RR~Gem} are summarized in the next items.  \\begin{figure}[h] \\centering \\includegraphics[width=8.8cm]{amp.eps} \\caption{Intensity and magnitude mean $\\overline{V}$ magnitude and $\\overline{B}-\\overline{V}$ and $\\overline{V}-\\overline{I}$ colours of \\object{RR~Gem} in the different Blazhko phases. The intensity mean colours indicate  colour variations during the Blazhko cycle, namely the star is slightly cooler in the smaller amplitude phase than during the highest amplitude period. } \\label{amp} \\end{figure}  \\begin{itemize}  \\item {First detection of a very low amplitude Blazhko modulation with symmetrically placed triplets in the Fourier spectrum.}  \\item {The modulation is concentrated in a 0.2 phase interval of the pulsation, centred on the mid of the ascending branch.}  \\item {In contrast to the exponential decrease in the amplitudes of the harmonic components of the radial pulsation, the decrease in the amplitudes of modulating components is close to linear.}  \\item {In contrast to theoretical expectations the amplitude ratios and phase differences of the $\\pm$ modulation components are not the same for the different order modulation component pairs and in the different colours.}  \\item {According to the intensity mean colours and magnitudes in the different phases of the Blazhko cycle, there is very small if any mean $V$ brightness variation along the Blazhko period, while the mean colours indicate small but definite changes in the sense that the star is 30-40~K cooler in its small amplitude phase.}  \\item {The light curve cannot be completely fitted with symmetrically placed Fourier frequency components alone. Residual scatter of the light curve is still concentrated in the ascending branch, indicating some irregular behaviour of the modulation.}  \\end{itemize}  Any model which relates the Blazhko periodicity to the rotation of the star predicts a similar percentage of small amplitude modulation as large amplitude ones. Because of the incidental angle of view of the rotational axes, the observed amplitudes of the modulation has to vary uniformly from zero up to some tenths of a magnitude. On the contrary, no small amplitude Blazhko modulation has been known previously. This fact has to be carefully examined on the grounds of observational biases. The most extended observations of field RR~Lyrae stars were the photographic observations from the middle of the $20^{th}$~century. These observations were not, however, accurate enough to detect modulation amplitudes smaller than $\\sim$~$0.2$~mag. Photoelectric observations of field RR~Lyrae stars are very sparse and usually time limited, so again, they were not suitable to detect small amplitude, long period modulations. The cluster and extragalactic variable surveys are usually seriously affected by crowding problems and are also often very time limited, thus small amplitude modulations might be missed or attributed to observational inaccuracy.  The detection of hundredths of magnitude amplitude modulation of \\object{RR~Gem} indicates that there might be a serious selection effect against small amplitude Blazhko stars. A systematic search using today's accurate photometric detectors is needed to reveal the true unbiased statistics of the modulation amplitudes.  Among the 731 fundamental mode Blazhko variables of the MACHO database there was only one star with a modulation period shorter than 10 days and only 5 with $P_{Bl}<15$~days \\citep{macho}. In contrast, \\object{RR~Gem} is already the second very short modulation period variable in the much smaller sample of well observed galactic field Blazhko variables (less than 50 stars). The other short period Blazhko star (\\object{AH~Cam}) was studied in detail by \\citet{smithah}. Both \\object{AH~Cam} and \\object{RR~Gem} are relatively metal rich variables with pulsation periods shorter than $0\\fd4$. These are their common features, but the modulation pattern of the two stars are very different. In \\object{AH~Cam} the amplitude of the modulation is $0.3-0.4$~mag, and phase modulation with about 0.1 pulsation phases occurs as well. A simplified explanation of the larger percentage of short modulation period Blazhko stars in our Galaxy than in the LMC would relate the short modulation periods to the larger metal content of these stars. However, the 530 day modulation period of \\object{RS~Boo} (\\cite{kanyo, nagy}) with $P_{puls}=0.377$~d and [Fe/H]=$-0.35$~dex \\citep{layden} indicates that a more complex description is needed.  Though \\citet{chadid}, favouring the nonradial mode explanation of the Blazhko phenomenon, argue against the presence of a strong magnetic field in \\object{RR~Lyrae}, our detailed analysis of \\object{RR~Gem} revealed a challenge to these models as well. The two most important results which would be hard to explain by nonradial modes are the followings.  a) The linear decrease of the amplitudes of the modulation components compared to the exponential decrease in the amplitudes of the radial mode (Fig.~\\ref{amp2}) indicates significant differences between these frequency components. As nonlinearity is proportional to the amplitude, it is expected that if the harmonic terms of the radial mode invoked by nonlinear coupling decrease exponentially, then the nonlinear coupling terms of the much smaller amplitude nonradial modes have to decrease even more steeply. The linear decrease of the amplitudes of the very small amplitude modulation components indicates a different origin of its higher order components than nonlinear coupling. This anomalous behavior of the amplitudes of the modulation frequencies  arises from the high concentration of the modulation in only 0.2 phase intervals of the pulsation.  b) The modulation is centred on the phase when H emission is assumed to occur. Such a connection between a temporal physical phenomenon in the atmosphere and the occurrence of the modulation suggests that during the Blazhko period real changes in the atmosphere are present. This is also supported by the detected colour variations during the Blazhko cycle.  Our detailed photometric observations of \\object{RR~Gem} have given new insight into the Blazhko phenomenon. However, it is an open question whether the Blazhko properties of \\object{RR~Gem} are unique or give general information about the modulation as well. To find the answer, further similar studies of other Blazhko variables are planned."
    },
    "0410/astro-ph0410433_arXiv.txt": {
        "abstract": "The ESO Supernova Progenitor Survey (SPY) has identified 58 (mostly helium-rich) subluminous O stars. We use the Balmer line strength to distinguish sdO from He-sdO (no Balmer lines) and present the results of the analyses of high resolution optical VLT-UVES spectra using an extensive grid of NLTE atmosphere models covering a large range in T$_\\mathrm{eff}$, $\\log{g}$ and helium abundances. The stellar atmospheric parameters are derived from line profile fits using a $\\chi^2$ technique. The resulting distribution in the (T$_\\mathrm{eff}$, $\\log{g}$) diagram as well as the luminosity function are discussed in the context of stellar evolution scenarios. By combining our results with those for the sdB stars from SPY (Lisker et al.~2004) we discuss the implications for binary population synthesis models of Han et al.~(2003). Models with a low CEE efficiency and a constant mass ratio distribution provide a reasonable explanation of the observed properties of the SPY sample of sdB and sdO stars indicating that the sdO stars form the hot and luminous extension of the sdB sequence. However, for the He-sdO stars none of the considered evolution scenarios are in agreement with the measured parameters of our programme stars. We conclude that He-sdO stars are formed by a different process than the sdB and sdO stars. ",
        "introduction": "The ESO Supernova Progenitor Survey (SPY; Napiwotzki et al.~2003) has identified 137 hot subluminous stars. 79 of them were classified as hydrogen rich sdB stars, 58 of them were classified as subluminous O stars. The sdB spectra have been studied by Lisker et al.~(2004, see also these proceedings). They tested several evolutionary scenarios by comparing the results of their spectroscopic analyses to model predictions. In particular, the binary population synthesis models of Han et al.~(2003, hereafter HPMM) were discussed. However, two diagnostic tools used yielded conflicting results. Moreover, the HPMM models predict that some EHB stars should be even hotter than those found in the sdB star sample of Lisker et al.~(2004). This led the authors to the conclusion that the subdwarf sample may not be sufficiently complete to describe the whole parameter range covered by the simulations. By including the subluminous O stars we attempt here to correct for this observational bias in the study of Lisker et al.~(2004). ",
        "conclusions": "We conclude that those binary population synthesis models of HPMM with a low CEE efficiency and a constant mass ratio distribution can explain the observed properties of the SPY sample of sdB and sdO stars. Possible explanations for the evolution of He-sdO stars are more difficult to find. None of the considered evolution scenarios showed good agreement with the measured parameters of our programme stars. The lack of radial velocity variability of He-sdO stars (Napiwotzki et al.~2004) may support the merger hypothesis. However, the measured temperatures are too high. The same holds for the gravities of some of the He-sdO stars. It should be noted that in the HPMM models assumptions about the remaining hydrogen-rich envelope had to be made. The position of a model star in the (T$_\\mathrm{eff}$, $\\log{g}$) plane, however, depends on the envelope mass. By reducing the adopted envelope masses in the HPMM models it may be possible to match the observed positions of the He-sdO stars. The spectra of the He-sdO stars display lots of metal lines. Metal line blanketing, however, has been neglected in the model atmosphere calculations, which may have led to an overestimate of the effective temperatures. A metal line blanketed NLTE model grid needs to be calculated to derive more accurate parameters of He-sdO stars.  A larger sample of subdwarfs is needed to discriminate between the various scenarios and hypotheses outlined above. The Sloan Digital Sky survey (SDSS) will provide a promising source in the near future. Kleinman et al.~(2004) already classified 240 subdwarfs from the first data release. At the end of the SDSS project spectra of more than 1000 hot subdwarfs will become available for spectral analyses."
    },
    "0410/astro-ph0410743_arXiv.txt": {
        "abstract": "{ This is the fourth paper in a series studying star formation rates, stellar components, metallicities, and star formation histories of a blue compact galaxy (BCG) sample. Using  \\ha, \\oii, infrared (IR), radio (1.4\\,GHz) luminosities and  neutral hydrogen (\\hi) gas masses, we estimated star formation rates and gas depletion timescales of 72 star-forming BCGs. The star formation rates of the BCGs in our sample span nearly four orders of magnitude, from approximately 10$^{-2}$ to $10^2$\\,M$_\\odot$\\,yr$^{-1}$, with a median star formation rate of about 3\\,M$_\\odot$\\,yr$^{-1}$. The typical gas depletion timescale of BCGs is about one billion years. Star formation could be sustained at the current level only on a timescale significantly lower than the age of the universe before their neutral gas reservoir is completely depleted.  To assess the possible systematic differences among different star formation rate indicators, we compared the star formation rates derived from \\ha, \\oii, IR, and radio luminosities, and investigated the effects from underlying stellar absorption and dust extinction. We found that subtracting underlying stellar absorption is very important to calculate both dust extinction and star formation rate of galaxies. Otherwise, the intrinsic extinction will be overestimated, the star formation rates derived from \\oii\\ and \\ha\\ will be underestimated (if the underlying stellar absorption and the internal extinction were not corrected from the observed luminosity) or overestimated (if an overestimated internal extinction were used for extinction correction). After both the underlying stellar absorption and the dust extinction were corrected, a remarkably good correlation emerges among \\ha, \\oii, IR and radio star formation rate indicators. Finally, we find a good correlation between the measured star formation rate and the absolute blue magnitude, metallicity, interstellar extinction of BCGs. Our results indicate that faint, low-mass BCGs have lower star formation rates. ",
        "introduction": "Understanding the star formation history of the universe is the primary goal of much current research in astronomy (Madau \\etal 1996; Steidel \\etal 1999). Star formation rate (SFR) is a crucial ingredient to understand the star formation history of galaxies at all redshifts, as well as the global evolution of the Universe. To obtain this understanding, a reliable estimate of the SFR in individual galaxies is required. Many calibrations of SFR depend on the luminosity measured at various wavelengths, including radio, infrared (IR), ultraviolet (UV), optical spectral lines (such as \\ha, \\oii) and continuum (Kennicutt 1998). Using these SFR indicators, a number of studies of the SFR were performed at different redshift regimes (such as Gallego \\etal 1995; Pettini \\etal 1998; Sullivan \\etal 2001). Unfortunately, the agreement among these SFR indicators at different wavelengths is poor; the underlying reasons for the differences are not well understood (Charlot \\etal 2002).  Blue compact galaxies (BCGs) are characterized by their blue color, compact appearance, high gas content, strong nebular emission lines, and low chemical abundances (\\cite{kunth00}). In recent years BCGs have attracted a great deal of interest and have become key in understanding fundamental astrophysical problems. Most of the work carried out has focused on a statistical analysis of BCG samples by means of surface photometry (Papaderos et al. 1996; Vitores \\etal 1996; Doublier \\etal 1997; \\\"{O}stlin \\etal 1999; Cair{\\' o}s \\etal 2001; Gil de Paz \\etal 2003). The measured broadband colors are strongly affected by interstellar extinction and gaseous emission; therefore, optical broadband photometry alone does not allow us to derive their physical properties accurately.  Spectroscopic information is required since spectra of galaxies contain a wealth of additional information concerning the physical properties of galaxies. The continuum and the absorption lines provide information concerning the stellar content, while the nebular emission lines provide a measure of the SFR and the interstellar metallicity. During the last three decades, a number of major objective-prism surveys for extragalactic emission-line objects have been carried out.  Examples include the University of Michigan (UM) survey (MacAlpine, Smith, \\& Lewis 1977), the Universidad Complutense de Madrid (UCM) survey (Zamorano \\etal 1994, 1996), the KPNO International Spectroscopic Survey (KISS) (Salzer \\etal 2000). Some BCG candidates and other emission line objects have been discovered. Since the objective-prism spectra have very low spectral dispersion and small wavelength range, they cannot be used to investigate the physical properties of BCGs. Spectroscopic follow-up observations were carried out to study the characteristics of the survey BCG candidates (\\cite{salzer89}; \\cite{terlevich91}; \\cite{gallego96}). The observations were obtained with different telescopes at different observatories with different equipment and instrument setups. The spectra from these observations have relatively low dispersions and low S/N ratios.  We have undertaken an extensive long-slit spectral observation of 200 BCGs since 1997. Spectra were obtained at the Cassegrain focus of the 2.16 m telescope of National Astronomical Observatory of China. A 300 line mm$^{-1}$ grating was used, giving a dispersion of 4.8 \\AA\\,pixel$^{-1}$. All 200 spectral observations were finished by February 2004, and 97 luminous BCG spectra were studied in this series. The observations have been performed with the same instruments; the reductions and analyses are homogeneous.  The blue color, high gas content and strong nebular emission-line spectrum of BCGs indicate that they have a high star formation activity. One of the key parameters in developing an understanding of the formation and evolution of BCGs is the determination of the SFRs. While SFRs are central to discussions of the evolution of BCGs, there are surprisingly few quantitative SFRs for BCGs available in the literature.  In addition, the values available in the literature suggest that the SFRs of BCGs might vary over five orders of magnitude, and the mean SFRs given in the existing investigations are noted to also vary over two orders of magnitude. It is therefore not clear whether and to what extent the wide range of BCG SFRs results from the use of different techniques, assumptions, or sample selections, employed in the available studies.  Furthermore, the studies of intermediate redshift compact galaxies in the Hubble deep field (\\cite{guzman96}; \\cite{phillips97}) reveal that these include galaxies with properties very similar to those of local luminous BCGs, such as small sizes, high luminosities and blue colors. Because spectroscopic observation for these intermediate redshift galaxies is difficult (for example, \\ha\\ is shifted outside the optical range), some other SFR indicators or methods must be used to determine the SFRs of these intermediate redshift galaxies (\\cite{guzman97}; \\cite{hammer01}; \\cite{guzman03}). Therefore, estimating SFRs of nearby BCGs with different SFR indicators may help us in accurately estimating SFRs of intermediate redshift galaxies and studying their evolutionary history.  Based on our homogeneous BCG optical spectral sample, we calculate SFRs and gas depletion timescales of a sample of star-forming blue compact galaxies, using the \\ha\\ and \\oii\\ emission line fluxes and the IR and radio luminosities. We compare the global SFRs derived from \\oii\\ and \\ha, with or without corrections for dust extinction and underlying stellar absorption, to the SFRs derived from the IR and radio luminosities of these galaxies and investigate the source of the discrepancies among these different SFR indicators. We also investigate the relationship between the intrinsic SFR and other galaxy properties. We found that subtracting underlying stellar Balmer absorptions is very important in estimating dust extinction and SFR of galaxies.  The paper is organized as follows. In Sect. 2 we review the spectroscopic observation and reduction, and describe the {\\it IRAS} data, the radio 1.4\\,GHz observations. In Sect. 3 we calculate multiwavelength SFR estimates for the galaxy sample using the available flux data. In Sect. 4, we compare the results obtained using the different SFR estimators and investigate the underlying reason for their discrepancies. We present some discussion in Sect. 5 and our conclusions in Sect. 6. Throughout this paper we assume $H_0=71\\,\\rm{km\\,s^{-1}\\,Mpc^{-1}}$, and adopt the Salpeter initial mass function with mass ranges $0.1\\,\\Msolar < {\\rm M} < 100\\,\\Msolar$. ",
        "conclusions": "\\subsection{Gas depletion timescales}  \\begin{figure} \\centering \\includegraphics[angle=-90,width=0.45\\textwidth]{aa0852f6.ps} \\caption{The gas depletion time-scale of BCGs from \\ha, \\oii\\ emission lines, far-infrared, and radio luminosities. The typical gas depletion time-scale of BCGs is about a few billion years.} \\label{tau} \\end{figure}  The star formation history of BCGs is still an open question, BCGs are either genuinely young galaxies, or old galaxies that have resumed forming stars at a prodigious rate. In the third paper of this series (Kong \\etal 2003), using a technique of empirical population synthesis based on a library of observed star cluster spectra, we find that BCGs are typically age-composite stellar systems; in addition to young stellar populations, intermediate-age and old stellar populations contribute significantly to the 5870 \\AA\\ continuum emission of most galaxies in our sample. Since all BCGs have been found to contain a population of old stars, our results suggest that BCGs are primarily old galaxies with discontinuous star formation histories.  Since the galaxies in our sample were selected from some large neutral hydrogen (\\hi) surveys (Gordon \\& Gottesman 1981; Thuan \\& Martin 1981), the \\hi\\ gas mass ($M_{\\rm HI}$) of our sample galaxies was measured by these \\hi\\ surveys. The SFR of our sample galaxies were measured in Sec. 3, using the \\ha, \\oii\\ emission lines, far-infrared, and radio luminosities. Therefore, we can calculate the gas depletion time-scale of BCGs, obtained by dividing the available gas mass by the present SFR:  \\begin{equation} \\tau_{gas}= M_{\\rm HI}/{\\rm SFR}, \\end{equation}  where $M_{\\rm HI}$ is the total \\hi\\ mass content and SFR is the ongoing star formation rate.  Figure~\\ref{tau} shows the distribution of gas depletion timescales for the galaxies in our sample. Most galaxies are consuming their interstellar gas at a rate which can only be sustained for another few billion years. The median value of the gas depletion time-scale of BCGs is less than or about one billion years, and the gas depletion time-scale of BCGs is generally much shorter than the age of the universe. It indicates that star formation in BCGs could be sustained at the current level only on a time scale significantly lower than $1/H_0$ (cosmological time) before their neutral gas reservoir will be completely depleted.  To understand the evolutionary connection between blue compact star-forming galaxies observed at intermediate redshifts ($0.1 < z < 1$) and nearby blue compact galaxies, Pisano \\etal(2001) have compared Arecibo \\hi\\ 21 cm spectroscopy and Keck HIRES spectroscopy of 11 nearby blue compact galaxies.  Using the \\sfrha\\ and the \\hi\\ mass, they calculated the gas depletion timescale, $\\tau_{gas}= M_{\\rm HI}/{\\rm SFR}$, of each galaxy. The SFRs of BCGs range from 0.23$-$23.5 M$_\\odot$\\,yr$^{-1}$, and gas depletion timescales range from 0.18$-$7.29 Gyr.  Using the optical spectroscopy, deep optical/near-IR photometry of 4 luminous blue compact galaxies, and spectral evolutionary models, Bergvall \\& {\\\" O}stlin (2002) concluded that star formation in luminous BCGs is intense, corresponding to gas depletion timescales of the order of 100 Myr. Our results are consistent with those works: the luminous BCGs have shorter gas depletion timescale, and the typical gas depletion timescale of BCG is about a few billion years.  Based on surface photometry and spectroscopic analysis, most of BCGs have been found to contain a population of old stars. Based on the gas depletion estimation, the high star formation activities of BCGs must be transient. Therefore, the general star formation history of BCGs is that these galaxies undergo a few or several short bursts of star formation followed by longer more quiescent periods.   \\subsection{SFRs versus M$_B$}  Figure~\\ref{mbsfr} shows the SFR against absolute B-band magnitude for BCGs; the SFRs were derived from the \\ha\\ (circles) and from IR (triangles) luminosities, and $M_B$ were taken from Kong \\& Cheng (2002a). There is a trend between B-band magnitude and SFR of BCGs, the Spearman Rank correlation coefficient being $r_s = -0.9$. The higher SFRs are clearly being hosted by galaxies with higher luminosities although there is a significant scatter in the trend. The least-squares fit of these data, shown by the dash line, yields an increased SFR towards high luminosity galaxies, $ SFR = -(7.65 \\pm 0.25) - (0.42 \\pm 0.01) {\\rm M}_B$. This relation implies that the SFR per solar luminosity in BCGs is constant over a range of 7 magnitudes.  \\begin{figure} \\centering \\includegraphics[angle=-90,width=0.45\\textwidth]{aa0852f7.ps} \\caption{ Correlation between the star formation rates derived from the infrared, \\ha\\ luminosities and the absolute blue magnitude of galaxies. The data suggest that there exists a good correlation between star formation rates of BCGs and $M_B$. The dashed line shows the ordinary least-squares fit for BCGs. The different symbols indicate SFRs derived from \\ha\\ (circles) and IR (triangles) for BCGs in this paper, and IR (stars) for galaxies from Kong \\etal (2004).} \\label{mbsfr} \\end{figure}  In Kong \\etal (2004), the authors have compiled an ultraviolet-selected sample of 115 nearby, non-Seyfert galaxies spanning a wide range of star formation activities, from starburst to nearly dormant, based on observations with various ultraviolet satellites.  To check the correlation between SFR and $M_B$ for other galaxy types, we over-plot these galaxies in Figure~\\ref{mbsfr} as stars.  {\\it IRAS} flux densities $f_{\\nu} (25\\,\\micron)$, $f_{\\nu} (60\\,\\micron)$ and $f_{\\nu}(100\\,\\micron)$ were used to estimate the total far-infrared flux \\lir\\ and \\sfrir. The apparent blue magnitudes from De Vaucouleurs \\etal (1991) and the redshift from the NASA/IPAC Extragalactic Database (NED) were used to calculate the absolute B-band magnitudes. From Figure~\\ref{mbsfr}, we found there is a scattered but good correlation between the $M_B$ and the SFR of other galaxy types. Starburst galaxies lie substantially above the $M_B \\sim {\\rm SFR}$ relation for BCGs, and the less vigorously star forming galaxies are located in  the region below this relation. Our results are consistent with the much large sample of $\\sim10^5$ SDSS galaxies studied by Brinchmann \\etal (2004): they found a clear correlation between SFR and stellar mass (luminosity) over a significant range in log $M_*$.   While optical spectral line fluxes, UV, IR, and radio luminosity are commonly used as SFR indicators, the luminosity at B-band wavelengths is also dominated in starburst galaxies by young stellar populations. In the absence of other observations, the B-band luminosity may thus be used as an SFR indicator for nearby large-area image surveys (such as SDSS) and deep image surveys (such as COSMOS and  Subaru Deep Field , Ouchi \\etal 2003). BCGs can be selected according to their compactness and blue colors from images in these surveys, and the best-fit relation between ${\\rm M}_B$ and SFR can be used to estimate the SFR of BCG. The SFR based on ${\\rm M}_B$ is not very accuracy, since the scatter for the ${\\rm M}_B$ and SFR relation is large.   \\subsection{SFRs and dust extinction}  Star forming galaxies in the local universe have been shown to exhibit a correlation between obscuration and FIR luminosity (Wang \\& Heckman 1996). This implies that the obscuration in a galaxy is related to its SFR. This same general effect can also be seen in other recent studies (such as Takagi, Arimoto \\& Hanami 2003). Buat \\etal\\ (2002) found a good relation between the extinction in the \\ha\\ line and the FIR luminosity for starburst galaxies. Hopkins \\etal\\ (2001) and Sullivan \\etal\\ (2001) report a positive correlation between the dust extinction traced by the Balmer decrement and the SFR of the galaxies traced by their total FIR or \\ha\\ luminosities.  Figure~\\ref{sfrebv} shows the correlation between the dust extinction ($\\tau_V^{cor}$) and the SFR of BCGs in our sample. Using the data of \\sfrha\\ and \\sfrir, we perform a least squares fit to the $\\tau_V^{cor}$ and \\sfr\\ to derive $ SFR = -(0.41 \\pm 0.09) + (0.73 \\pm 0.0.08) \\tau_V^{cor}$.  The Spearman Rank correlation coefficient between $\\tau_V^{cor}$ and SFR is 0.62. There is a clear trend between the dust extinction and the SFR of BCGs; the value of SFR increases with $\\tau_V^{cor}$. We see evidence for a strong correlation between SFR and metallicity (in the next subsection), presumably reflecting a correlation between metallicity and dust content. The correlation between SFR and dust may simply reflect a more fundamental correlation between dust and metallicity. This question clearly merits further investigation. The scatter in these trends may be related to intrinsic differences within galaxy populations. Because greater SFRs correspond to greater extinction in the dust, then the SFRs of high redshift galaxies derived from either Balmer fluxes or from UV continuum measurements would tend to be systematically underestimated, if the dust extinction were not corrected. Since the dispersion between the dust extinction and the SFR of BCGs is large, the situation is worse for these low extinction galaxies (most BCGs have low $\\tau_V^{cor}$). This relation cannot be used to estimate the amount of dust extinction in BCGs.  \\begin{figure} \\centering \\includegraphics[angle=-90,width=0.45\\textwidth]{aa0852f8.ps} \\caption{The dust extinction traced by the Balmer decrement versus the SFR. $\\tau_V^{cor}$ were calculated from the \\fha\\ and \\fhb\\ fluxes where the underlying stellar population absorption were subtracted. } \\label{sfrebv} \\end{figure}   \\subsection{SFRs vs. Metallicity}  Figure~\\ref{sfrz} shows the correlation between oxygen abundance and SFR. The oxygen abundance of BCGs was estimated by the N2 = \\nii$\\lambda$6584\\AA/H$\\alpha$ estimator from the results of Denicol{\\' o}, Terlevich \\& Terlevich (2002). These authors found that the N2 estimator follows a linear relation with log(O/H) that holds for a wide abundance range, from about $1/50th$ to twice the Solar value.  A positive trend of 12+log(O/H) with \\sfrha\\ and \\sfrir\\ can be seen, in the sense that galaxies with higher global SFRs also show higher oxygen abundance. The Spearman Rank correlation coefficient for this correlation is 0.67, an unweighted linear regression yields $12+{\\rm log}(O/H) = 0.24 {\\rm log}(SFR) + 8.52$.  Zaritsky, Kennicutt \\& Huchra (1994) found the luminosity (mass)-metallicity relationship exists over a large range of luminosity, galaxy type, and metallicity. A similar relationship was seen for BCGs in our sample (see Fig. 9 in Kong \\etal 2002). This correlation reflects the fundamental role that galaxy mass plays in galactic chemical evolution. With the luminosity and SFR correlation (see Figure~\\ref{mbsfr}), the most straightforward interpretation of the correlation between SFRs and metallicity is that more massive galaxies form fractionally more stars in a Hubble time (higher SFR) than low-mass galaxies, and then have higher metallicity.  From Figure~\\ref{sfrz}, we also found that the slope of the SFR and metallicity relation for these low-mass galaxies increases more quickly and the slope of the more massive galaxies increases more slowly. It can probably be explained as in massive galaxies, which are already relatively metal rich, it is more difficult to significantly alter $12+{\\rm log}(O/H)$ than it would be in a low-mass, low-metallicity galaxy.    \\begin{figure} \\centering \\includegraphics[angle=-90,width=0.45\\textwidth]{aa0852f9.ps} \\caption{Metallicity, 12 + log(O/H), as a function of SFR (\\sfrha\\ and \\sfrir). Galaxies with higher global SFRs also show higher oxygen abundance.} \\label{sfrz} \\end{figure}"
    },
    "0410/astro-ph0410096_arXiv.txt": {
        "abstract": "We present a  high-resolution (24\\s) study of the HI interstellar gas distribution around the radio-quiet neutron star Geminga. Based on Very Large Array (VLA)  and MPIfR Effelsberg telescope data, we analyzed a 40\\m$\\times$40\\m~ field around Geminga. These observations have revealed the presence of a neutral gas shell, 0.4 pc in radius, with an associated HI mass of 0.8 M$_{\\odot}$, which surrounds Geminga at a radial velocity compatible with the kinematical distance of the neutron star. In addition, morphological agreement is observed between the internal face of the HI shell and the brightest structure of Geminga's tail observed in X-rays. We explore the possibility that this morphological agreement is the result of a physical association. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410213_arXiv.txt": {
        "abstract": "The {\\em Chandra} X-ray observatory monitored the single cool star, AB Doradus, continuously for a period lasting 88\\,ksec (1.98\\,$P_{\\rm rot}$)  in 2002 December with the LETG/HRC-S. The X-ray lightcurve shows significant rotational modulation. It can be represented as having a flat level of emission superimposed with bright flaring regions that appear at the same phases in both rotation cycles. Phase-binned O\\,{\\sc viii} line profiles show centroid shifts that also repeat in consecutive rotation cycles. These Doppler shifts trace a roughly sinusoidal pattern with a a semi-amplitude of $30 \\pm 10$\\kmsec. By taking both the lightcurve and spectral diagnostics into account along with constraints   on the rotational broadening of line profiles (provided by archival  Chandra HETG  {Fe}\\,{\\sc xvii} line profiles) we can construct a simple model of the X-ray corona. The corona can be described as having two components, one component is homogeneously distributed, extending less than 1.75\\rstar; and the other consists of at least two compact emitting regions near the stellar surface. These  compact regions account for 16\\% of the X-ray emission and are likely to be located less than 0.4\\rstar\\ above the stellar surface. ",
        "introduction": "While the thermal properties of active stellar coronae are increasingly well-determined, we have yet to establish where the emitting plasma is located. As active stars cannot be spatially resolved at X-ray wavelengths we have to rely on indirect techniques to infer the distribution of coronal material. Grating spectra from the {\\em Chandra}/LETG configuration are sufficiently resolved to measure radial velocities in stellar coronae (Brickhouse et al. 2001; Hoogerwerf et al. 2004).  Previous studies of AB Dor data taken using ROSAT show evidence of rotational modulation at the 5--13\\% level (K\\\"urster et al. 1997). We set out to study coronal structure in the active single K0 dwarf, AB Dor ($P_{rot}=0.51$\\,d, $v_e \\sin i = 90$\\,km\\,s$^{-1}$) by analysing {\\em Chandra}/LETG data covering almost two rotation cycles. We find strong evidence for rotational modulation in both the spectra and lightcurves and use the results to develop a simple model for the star's X-ray emitting corona. Section~2 summarises the properties of the LETG X-ray lightcurve, Section~3 details measurements of rotational modulation in the LETG spectra and Section~4 shows the upper limits to which the bulk of the X-ray emitting corona extends, as determined using archive {\\em Chandra}/HETG spectra. The main conclusions and a simple coronal model are presented in Section~5. ",
        "conclusions": "Our main conclusions can be itemised as follows. \\begin{itemize}  \\item The X-ray lightcurve has a flat level of emission superimposed with two (or three) bright/flaring regions; one region lies between phases, $0.8<\\phi<1.0$ and one or two regions lie between $0.1<\\phi<0.7$. These bright flaring regions appear to repeat in both (consecutive) rotation cycles observed.  \\item Phase-binned O\\,{\\sc viii} line profiles show centroid shifts that also repeat from cycle to cycle and are likely to be caused by one or more compact emitting regions. If only one region is reponsible for this modulation, it must be caused by a compact mid-to-high latitude region.  \\item We measure  a strong upper limit for rotational broadening using archive {\\em Chandra}/HETG spectra. Taking instrumental and thermal broadening into account, this corresponds to a coronal $v_e \\sin i < 180$\\,km\\,s$^{-1}$ (i.e. $R_{cor}<2$\\rstar).  \\end{itemize} A simple model of AB Dor's X-ray emitting corona consistent with the above diagnostics would have a homogeneously distributed component extending less than $1.75$\\rstar\\ and/or an active region at the pole. The rotational modulation in the lightcurve and spectra indicates the presence of at least two compact emitting regions. Peaks in the lightcurve indicate that the regions are unlikely to be located higher than 0.4\\rstar\\ above the stellar surface.  We have shown that while we still do not have the spectral resolution at X-ray wavelengths to comprehensively ``map'' the coronae of cool stars, we can trace velocity variations and use these to constrain the positions of emitting regions in the coronae of cool stars. Spectro-polarimetric ground-based observations of AB Dor have been carried out simultaneously with these {\\em Chandra } observations. We will use Zeeman Doppler imaging techniques to map the surface magnetic field of AB Dor at this epoch, we will extrapolate these surface maps to produce detailed 3D coronal magnetic field and X-ray emission models. The X-ray emission model will be tested against the results presented here and will be the subject of a future paper."
    },
    "0410/astro-ph0410739_arXiv.txt": {
        "abstract": "The longitudinal profile of the 320~EeV event observed by the Fly's Eye experiment is analysed. A method of testing the hypothesis of a specific primary particle type is described. Results for different particle types are summarized. For hadronic primaries between proton and iron nuclei, the discrepancy between observed and simulated profiles is in the range of 0.6-1.0$\\sigma$ for two different hadronic interaction models investigated. For primary photons, the discrepancy is 1.5$\\sigma$ assuming a standard extrapolation of the photonuclear cross-section with energy. Larger values of the cross-section at highest energies make primary photon showers more similar to hadron-initiated events. The influence of varying the extrapolation of the photonuclear cross-section is studied. ",
        "introduction": "Deciphering the primary particle type of the most-energetic cosmic rays would be an important step towards understanding their origin. Due to the stochastic process of air shower development, it is in general not possible to unambiguously identify the nature of the primary. However, if the expectation for a specific primary type were found to be inconsistent with the data, an {\\it exclusion} of such a primary particle hypothesis might be possible.  In this work, we summarize an analysis method developed for this purpose: By comparing large statistics simulations to the data of an individual event, the level of consistency between data and the showers simulated for a given primary particle type is evaluated. Results of applying this method to the 320~EeV Fly's Eye event~\\cite{flyseye} are presented. A detailed discussion is given in~\\cite{fe04}.  Of particular interest is the primary photon hypothesis. For the simulation of primary photon showers, the photonuclear cross-section $\\sigma_{\\gamma - air}$ has to be extrapolated to highest energies. It is quantified in this work, how the uncertainty in extrapolating $\\sigma_{\\gamma - air}$ influences shower features of primary photons. ",
        "conclusions": "Both primary photons and any hadron between proton and iron nuclei can not be excluded as primary particle of the 320~EeV Fly's Eye event. The method developed for this analysis is further exploited in an investigation of the most energetic AGASA events as presented at this conference~\\cite{piotrek}. The large uncertainty in extrapolating the photonuclear cross-section can considerably influence the predictions of features of primary photon events and deserves a further careful study.  {\\it Acknowledgments.} This work was partially supported by the Polish State Committee for Scientific Research under grants No.~PBZ~KBN~054/P03/2001 and 2P03B~11024 and in Germany by the DAAD under grant No.~PPP~323. MR is supported by the Alexander von Humboldt-Stiftung."
    },
    "0410/astro-ph0410355_arXiv.txt": {
        "abstract": "{We present a quantitative method to classify galaxies, based on multi-wavelength data and elaborated from the properties of nearby galaxies. Our objective is to define an evolutionary method that can be used for low and high redshift objects. We estimate the concentration of light ($C$) at the galaxy center and the 180$^o$-rotational asymmetry ($A$), computed at several wavelengths, from ultraviolet (UV) to I-band. The variation of the indices of concentration and asymmetry with the wavelength reflects the proportion and the distribution of young and old stellar populations in galaxies. In general $C$ is found to decrease from optical to UV, and $A$ is found to increase from optical to UV: the patchy appearance of galaxies in UV with no bulge is often very different from their counterpart at optical wavelengths, with prominent bulges and more regular disks. The variation of $C$ and $A$ with the wavelength is quantified.  By this way, we are able to distinguish five types of galaxies that we call spectro-morphological types: compact, ringed, spiral, irregular and central-starburst galaxies, which can be differentiated by the repartition of their stellar populations. We discuss in detail the morphology of galaxies of the sample, and describe the morphological characteristics of each spectro-morphological type. We apply spectro-morphology to three objects at a redshift $z\\sim1$ in the Hubble Deep Field North, that gives encouraging results for applications to large samples of high-redshift galaxies. This method of morphological classification could be used to study the evolution of the morphology with the redshift and is expected to bring observational constraints on scenarios of galaxy evolution. ",
        "introduction": "The morphology of galaxies is the result of physical processes that acted on these systems since their formation. In an attempt to organize galaxies and to understand their evolution, the Hubble-Sandage classification, based on the bulge-to-disk ratio, was built from the study of a sample of nearby bright galaxies (\\cite{sandage61}). However, it now appears that this classification is not appropriate for faint nearby galaxies or small-sized objects. Moreover, with telescopes reaching out to redshifts $z\\ge1$ the rate of unclassified galaxies seems to increase with the distance, and an important fraction has unusual morphology (\\cite{brinchmann98}, Abraham et al. 1996, van den Bergh 2002). Theoretical models predict this evolution by accounting for an evolution of the galaxy morphology (\\cite{baugh}). The origin of this morphological change with $z$ is likely twofold: the morphological K-correction (Burgarella et al. 2001) and the intrinsic galaxy evolution (\\cite{giavalisco}). A first step in elaborating a method of morphological classification is to be able to classify galaxies  within morphological boxes  and to calibrate the method where the results can be checked by eye, \\textit{i.e.} in the nearby universe. We also want to discriminate between merging and irregular galaxies, in order to confirm the validity of galaxy evolution models. The Hubble classification fails on this point (\\cite{naim97}). By comparing morphological characteristics of local and distant samples, it might be possible to constrain the mechanisms of formation and evolution of galaxies in the early universe. As a result, it is necessary to elaborate a new system of classification, which would be more objective than the Hubble classification. A scheme based on quantitative indices could have the advantage to not implicate human judgment and to be automatisable, in order to classify quickly large samples of galaxies. The quantification also allows us to measure the variation of the morphology with the wavelength ($\\lambda$). Finally, in order to compare low and high-redshift galaxies, these parameters have to be robust against a degradation of spatial resolution and low signal-to-noise ratios ($S/N$).\\\\ For several years, a number of parameters have been proposed. Morgan (1958, 1959) originally used the index of light concentration in the center of galaxy - the basis of the Hubble scheme with the morphology of the spiral structure. \\cite{kent} and \\cite{abraham94} developed two quantitative measures of concentration. Abraham et al. (1996) also used an asymmetry index of galaxies developed  by \\cite{schade95}. The resulting asymmetry-concentration diagram calibrated from a local sample allowed them to distinguish three morphological types in visible: E/SO galaxies, spiral galaxies and irregular/peculiar galaxies (Abraham et al. 1996; Bershady et al. 2000). Moreover, it seems possible to discriminate irregular galaxies from mergers by using a correlation between color index and asymmetry (Conselice et al. 2000).\\\\ \\cite{kuchinski} and \\cite{burgarella} compared rest-frame UV images with visible ones and made pioneering studies where they showed how asymmetry and concentration evolve with the wavelength for several galaxies. Therefore we must pay attention to compare objects that are actually comparable. More specifically, we must understand the behavior of the concentration and the asymmetry as a function of the wavelength before comparing results from samples at different $z$. Finally, a physical interpretation of the morphology and its evolution with $\\lambda$ is necessary to understand the formation and the evolution of galaxies through high-$z$ observations, and to constrain models.\\\\ This work relies on the largest good-quality data presently available from UV to visible wavelengths. The paper aims at describing how we can define a spectro-morphology classification from the local sample to $z\\sim1$ galaxies.  This work outlines the main parameters and analyses their relative strength and weakness.\\\\ This paper is organized as follows: the data are presented in Section 2. We discuss the concentration and asymmetry calculations in Section 3. Section 4 presents the results of multi-$\\lambda$ measurements. We calibrate these parameters from the nearby sample; and in part 4.1 we give an example of application at high $z$. Finally, we present in Section 5 our conclusions and perspectives. ",
        "conclusions": "A multi-wavelength study of a sample of nearby objects  allowed us to elaborate what we call spectro-morphology, \\textit{i.e.} the study of the behavior of asymmetry and concentration with the wavelength. We saw that two-stellar component galaxies (spiral galaxies of the Hubble classification) present the highest variation of both indices, whereas the variation of one-stellar component galaxies (ellipticals/lenticulars and irregulars) are less wavelength-dependent. More precisely these changes in concentration and in asymmetry with the wavelength between UV and R-band led us to define five spectro-morphological types, characterized by their repartition of stellar populations. Compact galaxies are concentrated and symmetric at all wavelengths. This constancy in $C$ and $A$ is the result of the presence of only one stellar component. Ringed galaxies are often early spiral galaxies with a prominent bulge in visible which exhibits a ring of star-forming activity in UV: the locus of young and old stellar populations are very different, which induce high changes in $C$ and $A$ from UV to optical bands. Spiral galaxies show a prominent bulge in visible and star-forming regions on the arms of the spiral structure in UV. The resulting changes in $A$ and $C$ with the wavelength are less dramatic than for ringed galaxies, but characterize the existence of a composite stellar population. Irregular galaxies have a concentration parameter close to 2.7 at all wavelengths, that is the value expected for a galactic disk, and are rather asymmetric. This class contains galaxies with a faint or inexistent central bulge, which are classified in the Hubble classification as late spiral or irregular. Their stellar population is dominated by young stars. Central starburst galaxies possess a very bright nucleus in UV that induces a high concentration. This starburst activity is often induced by gravitational interaction or the existence of an AGN. The correlation between $C(NUV)-C(R)$ and $A(NUV)-A(R)$ allowed us to identify these five spectro-morphological types.  To test the validity of our method at high redshift, we successfully applied it to three distant objects in the HDF at $z\\sim1$. The resulting morphology is consistent with an eye-ball morphology and these galaxies are classified as spiral galaxies."
    },
    "0410/astro-ph0410480_arXiv.txt": {
        "abstract": "We present mid-infrared spectra of the microquasar SS\\,433 obtained with the Infrared Space Observatory (spectroscopic mode of ISOPHOT) and compare them to the spectra of four Wolf-Rayet stars. The mid-infrared spectrum of SS\\,433 shows mainly \\ion{H}{i} and \\ion{He}{i} emission lines and is very similar to the spectrum of WR\\,147, a WN8(h)+B0.5V binary with a colliding wind. The 2--12\\,$\\mu$m continuum emission corresponds to optically thin and partially optically thick free-free emission from which we calculate a mass loss rate of $1.4-2.2\\times 10^{-4} M_\\odot \\,\\mathrm{yr}^{-1}$ if the wind is homogeneous and a third of these values if it is clumped, which is consistent with a strong WN stellar wind. We propose that this strong wind outflows from a geometrically thick envelope of material surrounding the compact object like a stellar atmosphere, imitating the Wolf-Rayet phenomenon. ",
        "introduction": "\\vspace*{-0.2cm} SS\\,433 is a microquasar, i.e. an X-ray binary with relativistic jets, which shows a very peculiar optical spectrum \\citep{margon84}: there are  the so-called ``stationary'' emission lines, including strong H$\\alpha$ lines, showing normal Doppler shift movements with a period of 13~days, and the less intense ``moving'' lines with huge Doppler shifts corresponding to relativistic velocities with a period of about 162~days. The latter are formed in relativistic jets which undergo a precession movement in 162 days. The parameters of this system are: $P_\\mathrm{orb}=13.08$\\,d, $P_\\mathrm{prec}=162.375\\pm0.011$\\,d, velocity of the ejections $v=0.2647\\pm0.0008$\\,c, inclination of the jet axis to the line of sight $i=78.05^\\circ\\pm0.5^\\circ$ and opening angle of the precession cone $\\theta=20.93^\\circ\\pm0.08^\\circ$ \\citep{eiken01}. The ejections and their precession movement are indeed observed in the radio images \\citep{hjellming81}. SS\\,433 is the only microquasar with continuous ejections and which shows evidence of \\emph{ions} accelerated to relativistic velocities (0.26\\,c). Despite intensive studies, the nature of the two stars in this binary system remains unknown. The presence of Wolf-Rayet-like lines added to the very luminous continuum in the visible and near-IR ranges led to associate the donor star to a Wolf-Rayet or Of star \\citep{murdin80,vandenheuvel80}.    \\vspace*{-0.3cm} ",
        "conclusions": "\\vspace*{-0.3cm} We showed that both the emission lines and the continuum of the mid-IR spectrum of SS\\,433 are compatible with a Wolf-Rayet type for the donor star, and preferably a late WN star (WN8 or later). However, the recent discovery of \\citet{gieshuang02} and \\citet{hillwig04} of absorption features in the blue spectrum of SS\\,433 suggests that the donor star is an A3-7\\,I supergiant star with $10.9 \\pm 3.1 \\, M_\\odot$. On another hand, \\citet{lopezmarshall04} modelling the Chandra X-ray spectrum during eclipse constrain the radius of the mass donor star to be $9.1 \\pm 1.0 \\, R_\\odot$ which corresponds to O6-O8 main sequence stars with $29 \\pm 7 \\, M_\\odot$ and $T=41000-35800$\\,K \\citep{lang92}. But such a radius could also correspond to a B4-5 giant star ($R=8\\,R_\\odot$, $M=7\\,M_\\odot$ and $T=15000$\\,K for a B5\\,III). Note that this constraint on the radius found by \\citet{lopezmarshall04} is not incompatible with the radius of a Wolf-Rayet star, however if the results of \\citet{hillwig04} are confirmed then the donor star is not a Wolf-Rayet star and thus {\\bf there is a phenomenon imitating such a star in SS\\,433.}  Before we discuss this point, let us see which constraints on the nature of the donor star we can provide from the mid-IR spectra of SS\\,433. We modeled the possible donor stars with black body emissions at a distance of 4.6\\,kpc. Both a giant B5 star and a main sequence O6-O8 star have negligible emission compared to the optically thin ISOPHOT flux density of the continuum and thus they are compatible with the spectrum of SS\\,433. But the emission of an A7I ($T$=8150\\,K, $R=60\\,R_\\odot$) is not compatible  unless its radius is $<30R_\\odot$. Then this constraint is compatible with the result of \\citet{hillwig04} who find the radius of the Roche lobe volume for the mass donor star to be $R_L = 28 \\pm 2 \\, R_\\odot$. Note that this corresponds to a star with an intermediate size between type I (supergiant) and type II (bright giant) stars.  Now, if the mass donor star of SS\\,433 is not a Wolf-Rayet star, then there is a phenomenon imitating the emission of such a star in this binary system. We propose that this phenomenon comes from the material surrounding the compact object: this material does not form a classical thin accretion disc but rather a thick torus or a thick envelope like a stellar atmosphere, which is ionized by the X-rays emitted in the close vicinity of the compact object and expelled % by radiation pressure, thus imitating the wind of a Wolf-Rayet star. There may be a ``classical'' thin accretion disc but only very close to the compact object and not detectable because of the surrounding material. This would explain that there is no evidence for an accretion disc in the Chandra X-ray spectrum of SS\\,433 according to \\citet{marshall02} and that SS\\,433 is not a strong X-ray source as the others microquasars or X-ray transients. As Cyg\\,X-3 showed very similar results to those of SS\\,433 concerning the IR spectrum and corresponding mass loss \\citep{koch02} we could ask ourselves about the possibility of Wolf-Rayet imitation in this X-ray binary too."
    },
    "0410/astro-ph0410163_arXiv.txt": {
        "abstract": "We present extensive radio observations of SN\\,2003L, the most luminous and energetic Type Ibc radio supernova with the exception of SN\\,1998bw (associated with GRB\\,980425).  Observations from the Very Large Array are well described by a fitting a synchrotron self-absorption model to the emission spectrum.  This model implies a sub-relativistic ejecta velocity, $\\overline{v}\\approx 0.2 c$, and a size of $r\\approx 4.3\\times 10^{15}$ cm at $t\\approx 10$ days.  The circumstellar density is suitably fit with a stellar wind profile, $n_e \\propto r^{-2}$ and a constant mass loss rate of $\\dot{M}\\approx 7.5 \\times 10^{-6}~\\rm M_{\\odot}~yr^{-1}$.  Moreover, the magnetic field follows $B\\propto r^{-1}$ and the kinetic energy of the radio bright ejecta is roughly $E\\approx 10^{48}$ erg assuming equipartition of energy between relativistic electrons and magnetic fields. Furthermore, we show that free-free absorption does not contribute significantly to the radio spectrum, since it implies ejecta velocities which are inconsistent with size constraints derived from Very Long Baseline Array observations.  In conclusion, we find that although SN\\,2003L has a radio luminosity comparable to that seen in SN\\,1998bw, it shows no evidence for a significant amount of energy coupled to relativistic ejecta.  Using SN\\,2003L as an example, we comment briefly on the coupling of ejecta velocity and energy in Type Ibc supernovae. ",
        "introduction": "Radio emission from core-collapse supernovae (SNe) have been sought primarily to study the circumstellar medium \\citep{wsp+86,c98}. However, starting with Type Ibc SN\\,1998bw, associated with the under-luminous gamma-ray burst GRB\\,980425 \\citep{paa+00,gvv+98,kfw+98}, radio observations have also been used a tracer of relativistic ejecta \\citep{kfw+98,lc99}.  This event prompted a storm of publications: SN\\,1998bw was the most luminous radio supernova (at early times), showed strong variability and produced copious gamma-ray emission relative to ordinary SNe Ibc. These data provided the first observational evidence that at least some SNe Ibc are associated with GRBs (e.g.~\\citealt{bkd+99,bkp+02,smg+03}, c.f.~\\citealt{pks+03}).  To further explore the connection between GRBs and SNe Ibc, we began a radio survey of local Type Ibc supernovae. Our survey was motivated by two simple phenomenological questions: (1) What is the prevalence of SN\\,1998bw-like supernovae?, and (2) Is there a continuum between ordinary SNe and GRBs as traced by their relativistic ejecta?  In \\citet{bkf+03} we reported the first three years of the survey.  We found that most SNe Ibc are not detectable with current radio sensitivity, with limits reaching $10^{-3}$ that of the peak brightness of SN\\,1998bw.  Clearly, luminous SN\\,1998bw-like supernovae are rare (less than 3\\%) and those within an order of magnitude of SN\\,1998bw are uncommon (less than 10\\%).  Here, we present radio observations of SN\\,2003L, the first SN Ibc with a radio luminosity comparable to SN\\,1998bw.  While there is no evidence for mildly-relativistic ejecta, the radio emitting ejecta of SN\\,2003L are unusually energetic.  This discovery suggests the existence of a sub-class of SNe Ibc with similar energetics to SN\\,1998bw and even weaker central engines, if any.  The organization of this paper is as follows: observations from the Very Large Array (VLA), the Very Long Baseline Array (VLBA) and the {\\it Chandra} X-ray Observatory (CXO) are described in \\S\\ref{sec:obs} and \\S\\ref{sec:xray}.  Preliminary estimates of the energy and velocity are presented in \\S\\ref{sec:ep}.  Modeling of the radio light-curves is discussed in Sections~\\ref{sec:ssa} and \\ref{sec:ffa} assuming the dominant absorption is from internal synchrotron self-absorption and external free-free absorption, respectively, Modeling of the X-ray emission follows as \\S\\ref{sec:xray_model}. In the Discussion (\\S\\ref{sec:disc}) we conclude by addressing the nature of SN\\,2003L. ",
        "conclusions": "\\label{sec:disc} As the first Type Ibc supernova with a radio and X-ray luminosity within an order of magnitude of SN\\,1998bw, SN\\,2003L is clearly an unusual event.  Here, we summarize our findings.  Both SSA and FFA models (Sections~\\ref{sec:ssa} and \\ref{sec:ffa}) show that the energy of the ejecta is $\\gtrsim 10^{48}$ erg and nearly comparable to that of SN\\,1998bw at a similar epoch.  In Case 1 the equipartition and SSA analyses predict a sub-relativistic velocity, $\\overline{v}\\approx 0.2c$, and a relatively low mass loss rate, $\\dot{M}\\approx 7.5\\times 10^{-6}~M_{\\odot}~\\rm yr^{-1}$ which are roughly consistent with a scenario where inverse Compton scattering produces the observed X-ray emission.  In Case 2, the complex FFA models reproduce the observed radio light-curves well, yet they imply a highly relativistic ejecta speed, $\\overline{v}\\approx 14c$, or a huge energy output of $E\\ge 10^{51}$ erg, as well as a high mass loss rate of $\\dot{M}\\approx 10^{-2}~M_{\\odot}~\\rm yr^{-1}$ which over-predicts the observed X-ray emission assuming a thermal bremsstrahlung emission model.  Our VLBA observation at $t\\approx 65$ days implies an ejecta size which is inconsistent with the extreme velocities required by the complex FFA model in Case 2. Only by deviating far from equipartition and significantly increasing the total energy of the ejecta can this inconsistency be resolved.  Moreover, the predicted mass loss rate is at least two orders of magnitude above the observed rates from Wolf-Rayet stars.  In conclusion, we find Case 2 to be infeasible and we therefore adopt Case 1 the most likely scenario. We note that Case 1 predicts a mass loss rate comparable to the low end of values observed for local Wolf-Rayet stars.  While similar to SN\\,1998bw in radio luminosity and ejecta energy, a main distinguishing characteristic between SN\\,2003L and SN\\,1998bw is their shock velocity: while the radio ejecta of SN\\,1998bw attained mildly-relativistic speeds, that of SN\\,2003L was clearly sub-relativistic as shown in Section~\\ref{sec:ssa}.  This discrepancy is echoed by optical spectroscopy which showed photospheric velocities in excess of $3\\times 10^4~\\rm km~s^{-1}$ for SN\\,1998bw \\citep{pcd+01} and $10^{4}~\\rm km~s^{-1}$ for SN\\,2003L \\citep{vcd+03,mck+03,skg+04}.  Therefore, we see that for SN\\,1998bw, mildly-relativistic ejecta was coupled with bright radio emission, while for SN\\,2003L the emission is bright {\\it despite} the modest ejecta velocities.  The question is thus, why is the radio ejecta of SN\\,2003L especially luminous and energetic?  One hypothesis considers the distribution of kinetic energy according to ejecta velocity.  SN\\,1998bw is a clear case where a significant fraction of the explosion energy has been deposited into mildly-relativistic material.  For SN\\,2003L, however, the bulk of the energy is coupled to sub-relativistic ejecta.  \\citet{tmm01} have shown that in some cases, high velocity ejecta can be produced from a whip effect racing down the ejecta.  The observed diversity in the coupling between velocity and energy could therefore be intrinsic to the SN Ibc population. A similar situation is observed for gamma-ray bursts which partition their energy into ultra-relativistic emission ($\\gamma$-rays) and mildly-relativistic afterglow emission (X-rays, optical, radio) by variable fractions.  Through our ongoing radio surveys of local SNe Ibc and cosmological gamma-ray bursts, we continue to map out the energetics of these explosions.  Our goal is to understand the link between Type Ibc supernovae and GRBs by studying properties of the explosion (ejecta velocity and kinetic energy) in addition to those of the progenitor star (CSM density profile and evolutionary mass loss rate).  Through these further studies, we aim to shed light on the illusive bridge between these two classes of cosmic explosions."
    },
    "0410/astro-ph0410449_arXiv.txt": {
        "abstract": "{We present diffraction-limited bispectrum speckle interferometry observations of four well-known Herbig Ae/Be (HAeBe) stars, \\object{\\lkha}, \\object{Elias~1}, \\object{HK~Ori} and \\object{V380~Ori}. For two of these, \\lkha\\ and Elias~1, we present the first unambiguous detection of close companions. The plane of the orbit of the new \\lkha\\ companion appears to be significantly inclined to the plane of the circumprimary disk, as inferred from the orientation of the outflow.  We show that the Elias~1 companion may be a convective star, and suggest that it could therefore be the true origin of the X-ray emission from this object. In the cases of HK~Ori and V380~Ori, we present new measurements of the relative positions of already-known companions, indicating orbital motion.  For HK~Ori, photometric measurements of the brightness of the individual components in four bands allowed us to decompose the system spectral energy distribution (SED) into the two separate component SEDs.  The primary exhibits a strong infrared excess which suggests the presence of circumstellar material, whereas the companion can be modelled as a naked photosphere. The infrared excess of HK~Ori~A was found to contribute around two thirds of the total emission from this component, suggesting that accretion power contributes significantly to the flux. Submillimetre constraints mean that the circumstellar disk cannot be particularly massive, whilst the near-infrared data indicates a high accretion rate. Either the disk lifetime is very short, or the disk must be seen in an outburst phase. ",
        "introduction": "Herbig Ae/Be stars (HAeBes) are young intermediate-mass (2--8$M_{\\odot}$) pre-main-sequence stars first defined by Herbig (1960). According to his definition, HAeBes should have emission lines in their spectra, an early spectral type and be associated with dark cloud material as an indication of their youth.  Despite the similarities between HAeBes and the lower-mass T Tauri stars, questions such as the existence of circumstellar disks, the multiplicity and clustering of the objects, the origin of X-ray emission (Zinnecker \\& Preibisch, 1995), and the outflow activity are much less clarified in the case of HAeBes.   The frequency and properties of binary systems amongst all stars are of great importance in understanding the star formation process, since they can point to certain star formation scenarios and rule out or strongly constrain others. The configuration of pre-main-sequence multiples is of special importance because the binary properties of the objects can be compared with their formation environment, for example whether in a T association or massive star forming region. There is also less scope for the system configuration to have been altered by stellar interactions or continuing accretion. The angular momentum of the binary can also be compared to that of circumstellar or circumbinary disks which may still exist at this stage.  The coplanarity of the system orbit, the individual disks of the components and a possible circumbinary disk are all major constraints on viable formation mechanisms.   \\begin{table*}[ht] \\caption{Summary of the observations. For each observation, the epoch and observing wavelength is shown, together with the recovered binary separation (in both mas and AU for the assumed distance), position angle (measured anticlockwise from north), and flux ratio (typical uncertainty 10--20\\%). Where photometric calibration was performed, the individual magnitudes of the components are also given.} \\label{postable} \\begin{center} \\begin{tabular}{ccccccccc} \\hline \\hline \\multicolumn{7}{l}{{\\bf \\lkha}} \\\\ \\hline Epoch   & $\\lambda_c$ & $\\Delta\\lambda$ & $\\rho$     & $\\rho$ & PA         & $m_A$ & $m_B$ &  $F_s/F_p$ \\\\ & [nm]        & [nm]            &   mas        &   AU     &            &                             \\\\ 1996.75 & 2191        & 411             &  66.3 $\\pm$4 & 39.8     & 242 $\\pm$2 &                            &       &     0.25         \\\\ 1997.79 & 2165        & 328             &  66.6 $\\pm$3 & 40.0     & 242 $\\pm$2 &                            &       &     0.20             \\\\ 1998.45 & 2110        & 192             &  65.8 $\\pm$3 & 39.5     & 245 $\\pm$2 &                            &       &     0.20          \\\\ 1999.74 & 1648        & 317             &  63.9 $\\pm$2 & 38.3     & 246 $\\pm$1 &                            &       &     0.17        \\\\ 2001.83 & 2115        & 214             &  58.4 $\\pm$3 & 35.0     & 250 $\\pm$1 &                            &       &     0.16           \\\\ 2002.73 & 1648        & 317             &  50.4 $\\pm$2 & 30.2     & 255 $\\pm$2 &                            &       &     0.16          \\\\ 2003.78 & 2115        & 214             &  48.6 $\\pm$3 & 29.2     & 253 $\\pm$2 &                            &       &     0.27           \\\\ \\hline \\multicolumn{7}{l}{{\\bf Elias~1}} \\\\ \\hline 1996.75 & 2191        & 411                &  50.5 $\\pm$4 & 7.07     & 54 $\\pm$3 &                          &       &     0.69     \\\\ 2003.76 & 2115        & 214                &  60.4 $\\pm$1 & 8.46     & 59 $\\pm$1 &                          &       &     1.00        \\\\ \\hline \\multicolumn{7}{l}{{\\bf HK~Ori}} \\\\ \\hline 1995.18 &  550        &  60             &  334.2 $\\pm$11.5  & 153.72        & 44.1 $\\pm$1.5 & 12.13$\\pm$0.2 & 13.01$\\pm$0.2 & 0.45  \\\\ 1995.18 &  656        &  60             &  333.4 $\\pm$11.5  & 153.35        & 44.6 $\\pm$1.5 & 11.62$\\pm$0.2 & 12.53$\\pm$0.2 & 0.43   \\\\ 1997.80 &  1613       &  304            &  339.2 $\\pm$6.2   & 156.03        & 43.8 $\\pm$0.5 & 8.55$\\pm$0.14 & 9.76$\\pm$0.16 & 0.33    \\\\ 1997.80 &  2165       &  328            &  343.1 $\\pm$6.8   & 157.83        & 44.2 $\\pm$0.5 & 7.42$\\pm$0.16 & 9.65$\\pm$0.17 & 0.13   \\\\ 2003.78 &  2115       &  214            &  347.7 $\\pm$2.5   & 159.95        & 41.8 $\\pm$0.7 &               &               & 0.17    \\\\ \\hline \\multicolumn{7}{l}{{\\bf V380~Ori}} \\\\ \\hline 2003.78 &  2115       &  214            &  122.7 $\\pm$2.5   & 56.44         & 224.0 $\\pm$2.0 &              &               & 0.27   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  Numerous studies using a variety of techniques have established that the majority of low-mass young stars are members of binary systems (e.g. Ghez et al. 1993, Leinert et al. 1993, Simon et al. 1995).  The T~Tauri population in the nearby Taurus-Auriga T-association displays a factor of 2 higher binary fraction than that seen amongst main-sequence stars.  At the high end of the stellar mass range, bispectrum speckle interferometry of the massive stars in the Orion nebula cluster revealed companions for 7 out of 13 target objects, which after correction for undetected systems was estimated to represent a binary frequency close to 100\\%, with an average of $>1.5$ companions per primary (Preibisch et al. 1999). This is about three times the multiplicity of low-mass stars, and may suggest an alternative formation scenario for the high mass objects. Studies aimed specifically at intermediate mass stars have been rare. Li et al. (1994) carried out an infrared imaging study of 16 sources and reported companions for nine of them, although of these many were widely separated and are probably not bound.  Leinert et al. (1997, hereafter LRH97) studied a larger sample of 31 stars using a speckle interferometry technique in the near-infrared, and found that 31\\% of the systems were multiple with separations between 50 and 1300~AU. This was estimated to represent an excess binarity among HAeBe stars compared to main-sequence G dwarfs of a factor of two. Pirzkal et al. (1997) used a speckle shift and add technique to search for companions around 39 HAeBe stars. They detected 9 multiple systems, and estimated on statistical grounds that this implied a true binary frequency amongst HAeBes of at least 85\\%.  A survey for spectroscopic binaries amongst HAeBe stars by Corporon \\& Lagrange (1999) found a binary frequency of 17\\%. For short periods ($<$100 days) the binary frequency was found to be 10\\%, compared to the frequency amongst WTTS spectroscopic binaries of $ 11 \\pm 4$\\% (Mathieu 1992) or the frequency amongst main sequence solar-mass stars of $ 7 \\pm 2$\\% (Duquennoy \\& Mayor 1991).     In this paper, we present bispectrum speckle interferometry at multiple epochs of four HAeBe systems.  For two of these, \\lkha\\ and Elias~1, our diffraction-limited images show unambiguously for the first time that they have close binary companions.  For the known binary HK~Ori, we separate the SEDs of the two companions and by considering the different colours show that the brighter object possesses a circumstellar disk which the secondary lacks.  For V380~Ori, we are able to show relative motion compared to the position measured by LRH97. ",
        "conclusions": "We have used bispectrum speckle interferometry to resolve for the first time the binary companions to the HAeBe stars \\lkha\\ and Elias~1. In the case of \\lkha, the new object lies at a separation of approximately 60~mas, or 36~AU from the brighter component. Our observations comprise seven data points taken at intervals of approximately one year over a seven year time period. Motion of the secondary with respect to the primary is seen and the path is suggestive of an arc, although a straight line cannot be excluded. Because the companion may well have an important effect on the interesting and well-studied circumstellar environment of \\lkha, we have fitted several possible orbits to these points and shown that the plane of the orbit cannot be viewed {\\bf exactly} edge-on {\\bf with an exactly east-west orientation}. This may imply that the binary is not coplanar with the inferred disk powering the outflow. The flux ratio of the Elias~1 system indicates that the new companion has a mass of around 2$M_{\\odot}$ and therefore may be a convective star. This opens the possibility that the X-ray emission of Elias~1 originates from the companion and is powered by conventional mechanisms. The SED of HK~Ori was separated into the contributions of the primary and secondary components. The IR excess of HK~Ori~A was found to contribute around two thirds of the total emission from this component, suggesting that {\\it accretion power} contributes significantly to the flux, in contrast to most HAeBe stars whose SEDs can often be well fitted with passive disks.  We fitted an accretion disk model to the HK Ori A SED and determined the accretion rate and disk mass. From these parameters, the inferred disk lifetime is found to be quite short compared to the suspected age of HK~Ori. It is possible that the disk could be replenished from a wider reservoir of material, but there are problems with such a picture; a spherical envelope should lead to extinction which is not observed, whilst in a coplanar system, accretion from a disk should proceed preferentially onto the secondary. Another explanation could be that the HK~Ori disk is seen in an outburst phase. Such a phase need not affect the outer disk, so that the accretion rate may not be uniform throughout.  Our relative positions for this system trace a path over the eight years of observations which may be the segment of an orbit. {Future near-simultaneous multi-wavelength high resolution observations will allow a more detailed picture of the disk}. High-resolution observations in the mid-infrared could determine whether or not component B has a passive disk. Further observations of all four stars will be necessary to build up reliable orbit determinations. \\lkha\\ in particular should be a fascinating subject for future study from this point of view, since the close binary companion may be responsible for much of the complexity in the circumstellar environment of this system."
    },
    "0410/astro-ph0410559_arXiv.txt": {
        "abstract": "Recent results from the KASCADE experiment on measurements of cosmic rays in the energy range of the knee are presented. Emphasis is placed on energy spectra of individual mass groups as obtained from an two-dimensional unfolding applied to the reconstructed electron and truncated muon numbers of each individual EAS. The data show a knee-like structure in the energy spectra of light primaries (p, He, C) and an increasing dominance of heavy ones ($A \\ga 20$) towards higher energies.  This basic result is robust against uncertainties of the applied interaction models QGSJET and SIBYLL which are used in the shower simulations to analyse the data. Slight differences observed between experimental data and EAS simulations provide important clues for further improvements of the interaction models.  The data are complemented by new limits on global anisotropies in the arrival directions of CRs and by upper limits on point sources. Astrophysical implications for discriminating models of maximum acceleration energy vs galactic diffusion/drift models of the knee are discussed based on this data.  \\vspace{1pc} ",
        "introduction": "A puzzling and most prominent feature of the cosmic ray (CR) spectrum is the so-called knee, where the spectral index of the all-particle power-law spectrum changes from approximately $-2.7$ to $-3.1$.  Several models have been proposed in order to explain this feature shown in Fig.\\,\\ref{fig:all-particle-compilation}, but none of them has managed to become broadly accepted.  Some models focus on a possible change in the acceleration mechanism at the knee~\\cite{Lagage83,Biermann93,Drury94}, e.g.\\ due to the limiting energy defined by the size and magnetic field strength of the acceleration region ($E_{\\rm max} \\la Z \\times (B \\times L))$.  Other ones discuss an increased leakage of CRs from the Galaxy due to a change in the confinement efficiency by galactic magnetic fields, e.g.~\\cite{Wdowczyk84,Ptuskin93,Candia02a,Roulet03}. Again, this results in a rigidity scaling of the knee according to the maximum confinement energy.  Finally, a third group of models attributes the effect of the knee to CR interactions at their sources, during their propagation in the Galaxy, or in the upper atmosphere.  Such scenarios include nuclear photo-disintegration processes by UV-photons at the sources~\\cite{Candia02b}, interactions of CRs in dense fields of massive relic neutrinos~\\cite{Wigmans03}, production of gravitons in high-energy pp collisions~\\cite{Kazanas01}, etc.  A recent review about this topic can be found e.g.\\ in Refs.~\\cite{Kampert01a,Haungs03}.  \\begin{figure}[t] \\centerline{\\epsfxsize=\\columnwidth\\epsfbox{allparticle-KASCADE.eps}} \\vspace*{-5mm} \\caption[]{Compilation of the all-particle cosmic ray spectrum showing the knee, the suggested second knee, and the ankle of the CR spectrum (compilation by H. Ulrich). \\label{fig:all-particle-compilation}} \\end{figure}  To distinguish between these models and allowing to answer the long pressing question about the origin of cosmic rays and about the knee in their spectrum, high quality and high statistics data are required over an energy interval ranging from at least 0.5 to 500 PeV. It appears worthwhile to mention that solving the old problem about the origin of CRs in the PeV region is a prerequisite also for an understanding of the highest energies in the GZK-region.  Due to the low flux involved at energies $\\ga 10^{14}$ eV, only extensive air shower (EAS) experiments are able to provide such data.  In EAS experiments, primary CRs are only indirectly observed via their secondaries generated in the atmosphere.  The most important experimental observables at ground are then the electromagnetic (electrons and photons), muonic, and hadronic components.  In addition or alternatively, some experiments also detect photons originating from Cherenkov and/or fluorescence radiation of charged particles in the atmosphere.  For a brief review about EAS observables and their experimental techniques the reader is referred to Refs.\\,\\cite{Kampert01b,Swordy02,Haungs03}.  Unfortunately, progress on interpreting EAS data has been modest mostly because of two reasons: Firstly, the EAS development is driven both by the poorly known high-energy hadronic interactions and their particle production in the very forward kinematical region as well as by uncertainties in the low energy interaction models influencing mostly the lateral particle density distribution functions \\cite{Drescher04}.  Secondly, due to the stochastic nature of particle interactions, most importantly the height of the very first interaction in the atmosphere, EAS are subject to large fluctuations in particle numbers at ground.  To make things even more complicated, the amount of fluctuations depends, amongst others, sensitively on the primary CR energy and mass~\\cite{Kampert01b}.  Here, it is very important to realize that EAS fluctuations are not to be mistaken as random Gaussian errors associated with the statistics in the number of particles observed at ground.  The latter one can be improved by the sampling area of an EAS experiment, while the former one is intrinsic to the EAS itself, carrying - for a sample of events - important information about the nature of the primary particle.  Clearly, both kinds of fluctuations have to be accounted for in the data analysis of steeply falling energy spectra in order to avoid misinterpretations of the observations. ",
        "conclusions": "KASCADE has provided a wealth of new high quality EAS data in the knee region giving important insight into the origin of the knee and of CRs in general.  Conclusive evidence has been reached on the knee being caused by light primaries (p + He) mostly. Furthermore, the data are in agreement with a rigidity scaling of the knee position giving support to an astrophysical origin by either maximum acceleration or diffusion/drift models of propagation.  Astrophysical parameters start to be constrained by the preliminary KASCADE data, as was demonstrated at the example of the GRB models of Refs.\\,\\cite{Wick04,Dar04}.  Further important constrains result from measurements of large scale anisotropies of CRs and by new limits on point sources.  Claims that were made in Ref.~\\cite{chili-source} about an excess of CRs from the Monogem SNR cannot be confirmed.  A particle physics interpretation of the knee appears to be excluded with a high level of confidence.  For example, interactions of CRs with background particles or photo-disintegration in the Galaxy would produce an abundance of secondary protons and result in a light mass composition above the knee energy, a result which is in contradiction to the present data.  Furthermore, $\\bar{\\nu_{e}}$'s with a mass of $\\simeq 0.5$~eV/$c^{2}$ as are needed in the model of Ref.~\\cite{Wigmans03} appear to be excluded by recent WMAP and 2dF data for neutrinos in case of degenerated masses, with the latter being suggested by recent oscillation and possibly by a neutrinoless double-beta experiment (for a recent discussion of this topic the reader is referred to Refs.~\\cite{Hannestad04},\\cite{Fogli04}).  Presently, more data and more observables are being analyzed within KASCADE, particularly in terms of composition analyses employing reconstructions of the muon production height \\cite{Buettner03}.  Together with measurements of energetic hadrons in the central calorimeter, the unfolding technique of electron and muon numbers in EAS has become a powerful tool to reconstruct the properties of primary particles in EAS and it also provides important clues on how to improve the hadronic interaction models employed in CORSIKA air shower simulations.  KASCADE-Grande has just started its routinely data taking and will extend the measurements up to $10^{18}$\\,eV, thereby allowing to verify the existence of the putative Iron knee marking the so-called second knee in the all-particle CR spectrum \\cite{Kampert03,Badea04}.  This, together with improved statistics for anisotropy measurements will allow to confront astro- and particle-physics motivated models of the knee in much more detail to the experimental data as has been possible up to now.  Furthermore, the use of radio antennas complementing the experimental KASCADE-Grande set-up may open a new window to future EAS observations on large scales \\cite{Horneffer03,Badea04}.  Interesting first observations have been made and are presently being analyzed in detail.  \\vspace*{5mm} {\\small {\\bf Acknowledgement:} It is a pleasure to thank the organizers of the meeting for financial support and for their invitation to participate in a very interesting and fruitful meeting conducted in a pleasant atmosphere.  This work is supported by Forschungszentrum Karlsruhe, the German Ministry for Research and Education (Grant 05 CU1VK1/9), and by a Polish KBN grant for the years 2004-6.}"
    },
    "0410/astro-ph0410303_arXiv.txt": {
        "abstract": "We calculate multicomponent radiatively driven stellar wind models suitable for central stars of planetary nebulae. Some of these stellar winds may be adequately modelled using one-component models, however for some of them multicomponent models are necessary. We obtain a range of stellar parameters for different types of mass-loss. ",
        "introduction": "Hot star winds are accelerated mainly by absorption of radiation in resonance lines. This process can be divided into two steps: % \\begin{enumerate} \\itemsep -1pt \\item transfer of light momentum to C, N, O, Fe, etc., by absorption of radiation, and by Thomson scattering to free electrons, \\item transfer of obtained momentum to predominant wind component (H, He). \\end{enumerate} Since the acceleration of different wind species may be different, their wind velocities may differ and so hot star winds have a multicomponent nature. For {\\em high density} winds (e.g. of galactic O stars) this multicomponent nature does not influence the wind structure, however for {\\em low density} winds multicomponent effects occur, for example frictional heating, decoupling of wind components, etc. (see Krti\\v{c}ka \\& Kub\\'at 2001). To estimate the importance of multicomponent effects for the winds of central stars of planetary nebulae we calculate multicomponent wind models for these stars. ",
        "conclusions": ""
    },
    "0410/astro-ph0410629_arXiv.txt": {
        "abstract": "In this short review, the predictions of inflation are presented and compared to the most recent measurements of the Cosmic Microwave Background (CMB) anisotropy. It is argued that inflation is compatible with these observations but that these ones are not yet accurate enough to probe the details of the scenario. ",
        "introduction": "Inflation is presently our most convincing scenario of the very early universe, not only because it is physically attractive, but also because it makes a series of definite predictions that can be tested concretely by means of astrophysical observations. In this short article, we quickly review what are these predictions and describe how the on-going high accuracy Cosmic Microwave Background (CMB) anisotropy measurements can probe the physics of inflation. It is argued that inflation is presently compatible with all the available data but that these ones are not yet accurate enough to test the details of the inflationary scenario. This review is organized as follows. In the next section, we present the inflationary predictions and, in particular, describe the theory of cosmological perturbations of quantum-mechanical origin. In the next section, we compare these predictions with the recently released CMB anisotropy observations. Finally, we quickly present our conclusions. ",
        "conclusions": ""
    },
    "0410/astro-ph0410135_arXiv.txt": {
        "abstract": "We investigate the viability of $f(R)$ theories in the framework of the Palatini approach as solutions to the problem of the observed accelerated expansion of the universe. Two physically motivated popular choices for $f(R)$ are considered\\,: power law, $f(R) = \\beta R^n$, and logarithmic, $f(R) = \\alpha \\ln{R}$. Under the Palatini approach, both Lagrangians give rise to cosmological models comprising only standard matter and undergoing a present phase of accelerated expansion. We use the Hubble diagram of type Ia Supernovae and the data on the gas mass fraction in relaxed galaxy clusters to see whether these models are able to reproduce what is observed and to constrain their parameters. It turns out that they are indeed able to fit the data with values of the Hubble constant and of the matter density parameter in agreement with some model independent estimates, but the today deceleration parameter is higher than what is measured in the concordance $\\Lambda$CDM model. ",
        "introduction": "The Hubble diagram of type Ia supernovae (hereafter SNeIa) \\cite{SNeIa,Riess04}, the anisotropy spectrum of the cosmic microwave background radiation (hereafter CMBR) \\cite{CMBR,WMAP,VSA}, the matter power spectrum determined by the large scale distribution of galaxies \\cite{LSS,Teg03} and by the data on the Ly$\\alpha$ clouds \\cite{Lyalpha} are all convincing evidences in favour of a new picture of the universe, unexpected only few years ago. According to this nowadays standard scenario, the universe is flat and undergoing an accelerated expansion driven by a mysterious fluid with negative pressure nearly homogeneously distributed and making up to $\\sim 70\\%$ of the energy content. This exotic component is what is called {\\it dark energy}, while the model we have just depicted is usually referred to as the {\\it concordance model}.  Even if strongly supported by the bulk of the available astrophysical data, this new picture is not free of problems. Actually, while it is clear how dark energy works, its nature remains an unsolved problem. The simplest explanation claims for the cosmological constant $\\Lambda$ thus leading to the so called $\\Lambda$CDM model\\footnote{It is common in literature to make no distinction between the concordance and the $\\Lambda$CDM model even if, strictly speaking, in the concordance model the dark energy may also be provided by a different mechanism.} \\cite{Lambda}. Although being the best fit to most of the available astrophysical data \\cite{WMAP,Teg03}, the $\\Lambda$CDM model is also plagued by many problems on different scales. If interpreted as vacuum energy, $\\Lambda$ is up to 120 orders of magnitudes smaller than the predicted value. Furthermore, one should also solve the {\\it coincidence problem}, i.e. the nearly equivalence of the matter and $\\Lambda$ contribution to the total energy density. As a response to these problems, much interest has been devoted to models with dynamical vacuum energy, dubbed {\\it quintessence} \\cite{QuintFirst}. These models typically involve a scalar field rolling down its self interaction potential thus allowing the vacuum energy to become dominant only recently (see \\cite{PR02,Pad02} for good reviews). Although quintessence by a scalar field is the most studied candidate for dark energy, it generally does not avoid {\\it ad hoc} fine tuning to solve the coincidence problem. Moreover, it is not clear where this scalar field comes from and how to choose the self interaction potential.  On the other hand, it is worth noting that, despite the broad interest in dark matter and dark energy, their physical properties are still poorly understood at a fundamental level and, indeed, it has never been shown that they are in fact two different ingredients. This observation motivated the great interest recently devoted to a completely different approach to quintessence. Rather than fine tuning a scalar field potential, it is also possible to explain the acceleration of the universe by introducing a cosmic fluid with an exotic equation of state causing it to act like dark matter at high density and dark energy at low density. An attractive feature of these models is that they can explain both dark energy and dark matter with a single component (thus automatically solving the coincidence problem) and have therefore been referred to as {\\it unified dark energy} (UDE) or {\\it unified dark matter} (UDM).  Some interesting examples are the generalized Chaplygin gas \\cite{Chaplygin}, the tachyonic field \\cite{tachyon}, the condensate cosmology \\cite{Bruce} and the Hobbit model \\cite{Hobbit}. It is worth noting, however, that such models are seriously affected by problems with structure formation \\cite{Sandvik} so that some work is still needed before they can be considered as reliable alternatives to dark energy.  Actually, there is still a different way to face the problem of cosmic acceleration. As  stressed in Lue et al. \\cite{LSS03}, it is possible that the observed acceleration is not the manifestation of another ingredient in the cosmic pie, but rather the first signal of a breakdown of our understanding of the laws of gravitation. From this point of view, it is thus tempting to modify the Friedmann equations to see whether it is possible to fit the astrophysical data with a model comprising only the standard matter. Interesting examples of this kind are the Cardassian expansion \\cite{Cardassian} and the DGP gravity \\cite{DGP}.  In this same framework, there is also the attractive possibility to consider the Einsteinian general relativity as a particular case of a more fundamental theory. This is the underlying philosophy of what are referred to as $f(R)$ theories \\cite{capozcurv,MetricRn,review,PalRn,lnR,ABF04}. In this case, the Friedmann equations have to be given away in favour of a modified set of cosmological equations that are obtained by varying a generalized gravity Lagrangian where the scalar curvature $R$ has been replaced by a generic function $f(R)$. The usual general relativity is recovered in the limit $f(R) = R$, while completely different results may be obtained for other choices of $f(R)$. While in the weak field limit the theory should give the usual newtonian gravity, at cosmological scales there is an almost complete freedom in the choice of $f(R)$ thus leaving open the way to a wide range of models.  The key point of $f(R)$ theories is the presence of modified Friedmann equations obtained by varying the generalized Lagrangian. However, here lies also the main problem of this approach since it is not clear how the variation has to be performed. Actually, once the Robertson\\,-\\,Walker metric has been assumed, the equations governing the dynamics of the universe are different depending on whether one varies with respect to the metric only or with respect to the metric components and the connections. It is usual to refer to these two possibilities as the {\\it metric} and the {\\it Palatini approach} respectively. The two methods give the same result only in the case $f(R) = R$, while lead to significantly different dynamical equations for every other choice of $f(R)$ (see \\cite{MFF1987,FFV,ABF04,ACCF} and references therein).  It is worth noting $f(R)$ theories were initially investigated using the metric approach \\cite{capozcurv,MetricRn,review}. Even if some interesting and successful results have been obtained \\cite{curvfit}, this way to $f(R)$ theories is plagued by serious mathematical difficulties. Actually, even for the simplest $f(R)$, the metric approach leads to a fourth order nonlinear differential equation for the scale factor that is impossible to solve analytically and is affected by several problems that greatly complicate the search for numerical solutions. Moreover, some doubts have been cast on the consistency among the weak field limit of the theory and the newtonian gravity as tested at the Solar system scale \\cite{newtlimitno} even if some interesting different results have also been obtained \\cite{newtlimitok}.  On the other hand, theoretical considerations about the stability of the equations and the newtonian limit argue in favor of the Palatini approach to $f(R)$ theories. Moreover, the dynamics of the universe may be analytically determined from the cosmological equations obtained with this method for some interesting cases. To this aim, a clear mathematical machinery has been presented in Ref.\\,\\cite{ABF04} (hereafter ABF04) that allows to determine analytic expressions for the Hubble parameter as function of the redshift. As we will see later, this is all what is needed to test a given cosmological model.  The Palatini approach to $f(R)$ theories has been widely studied in literature \\cite{PalRn,lnR,MFF1987,FFV,ABF04,ACCF} and the dynamics of the cosmological models obtained by applying this method to different choices of $f(R)$ has been investigated in detail. Here we adopt an observational point of view on the Palatini approach. Assuming that this is the correct way to treat $f(R)$ theories, we investigate the viability of two classes of models obtained by two popular choices for $f(R)$, namely the power law $f(R) = \\beta R^n$ and the logarithmic $f(R) = \\alpha \\ln{R}$. To this aim, we compare the model predictions against the SNeIa Hubble diagram and the data on the gas mass fraction in relaxed galaxy clusters. This analysis will allow us to constrain the model parameters and to see whether $f(R)$ theories are indeed reliable alternatives to dark energy. Moreover, this will be an observational validation of the theoretically motivated Palatini approach.  The paper is organized as follows. Sect.\\,II details the method we employ to constrain the models and present the dataset we will use. The two classes of models we consider are briefly discussed in Sect.\\,III where we also individuate the parameters that are better suited to both assign the model and be constrained by the data. A detailed discussion of the results is the subject of Sect.\\,IV, while we summarize and conclude in Sect.\\,V. ",
        "conclusions": "Assuming that the Palatini (first order) approach is the correct way to treat $f(R)$ theories, we have investigated the viability of two different class of cosmological models corresponding to two popular choices of $f(R)$, namely a power law in the scalar curvature and a logarithmic function of $R$. The expansion rate $H = \\dot{a}/a$ may be analytically expressed as a function of the redshift $z$ for both classes of models so that it is possible to contrast the model predictions against the observations. In particular, we have used the SNeIa Hubble diagram and the data on the gas mass fraction in relaxed galaxy clusters to investigate the viability of each class as dark energy alternative and to constrain their parameters. The main results are sketched below.  \\begin{enumerate}  \\item{Both classes of models provide very good fits to the data even if the choice $f(R) = \\alpha \\ln{R}$ leads to a Hubble diagram that is in better agreement with the trend shown by the highest redshift SNeIa. However, the paucity of the data does not allow us to eventually prefer one model to the other.}  \\item{Eqs.(\\ref{eq: bfpl}) and (\\ref{eq: bflog}) give the best fit parameters for the power law and logarithmic Lagrangian models respectively. The confidence regions have been determined from the marginalized likelihoods and are reported in Eqs.(\\ref{eq: nrpl})\\,-\\,(\\ref{eq: obompl}) for the models with $f(R) = \\beta R^n$ and in Eqs.(\\ref{eq: oharlog})\\,-\\,(\\ref{eq: obomlog}) for those with $f(R) = \\alpha \\ln{R}$. To better compare the model predictions with previous results in literature, we have evaluated the present day deceleration parameter $q_0$ and the matter density parameter $\\Omega_M$ (assuming the estimate of $\\Omega_b h^2$ in Ref.\\,\\cite{Kirk}). For both classes of models, $q_0$ turns out to be higher than what is predicted for the concordance $\\Lambda$CDM model, i.e. $f(R)$ theories lead to models that accelerate less than what is observed. This result is however somewhat weakened by comparing with model independent estimates of $q_0$ even if these latter may be affected by systematic errors. As far as the matter content is concerned, for both classes of models $\\Omega_M$ is in very good agreement with what is inferred from galaxy clusters and estimated by fitting the $\\Lambda$CDM model to the available astrophysical data.}  \\item{To ameliorate the agreement with the observed $q_0$, one should increase the value of $n$ for the models with power law Lagrangians or decrease that of $\\Omega_M H_0^2 \\alpha^{-1}$ for models with $f(R) = \\alpha \\ln{R}$. In this case, a good fit to the data may still be obtained provided that both $h$ and $\\Omega_b/\\Omega_M$ are increased. While higher values of $h$ could still be compatible with the local estimates of the Hubble constant, increasing $\\Omega_b/\\Omega_M$ leads to lower values of $\\Omega_M$. Actually, the very good agreement among the estimated $\\Omega_M$ in Eqs.(\\ref{eq: ompl}) and (\\ref{eq: omlog}) and the results in literature is a strong evidence against this choice. Therefore, we conclude that it is not possible to recover the same value of $q_0$ in the concordance model by using power law or logarithmic Lagrangians.}  \\item{A model independent estimate of the present day age of the universe $t_0$ allows one to break the matter/geometry degeneracy inherent in $f(R)$ theories recovering the value of the coupling constant. For power law Lagrangians, this is indeed the only way to determine $\\beta$ thus offering the possibility to always recover a model that both fits the SNeIa Hubble diagram and the data on the gas mass fraction in relaxed galaxy clusters and also has the correct age. On the other hand, $t_0$ is an independent check for models with logarithmic Lagrangian since, in this case, it may be evaluated as a function of the two parameters $\\Omega_M H_0^2 \\alpha^{-1}$ and $h$ and compared with previous results in literature. It turns out that the predicted $t_0$ is lower than the value estimated for the $\\Lambda$CDM model and only marginally consistent with what is inferred from globular clusters and nucleochronology.}  \\end{enumerate}  The results summarized above may pave the way to the solution of an intriguing dilemma\\,: is Einsteinian general relativity the correct theory of gravity? If yes, then dark energy is absolutely needed to explain the accelerated expansion of the universe and hence all the theoretical efforts of cosmologists have to be dedicated to understanding its nature. On the contrary, if $f(R)$ theories are indeed able to explain the accelerated expansion, then it is time to investigate in more detail what is the right choice for the function $f(R)$ and how the variation has to be performed (higher order metric or first order Palatini approach).  From the observational point of view we have adopted here, there are no strong evidences against models with power law or logarithmic Lagrangians in the framework of the Palatini approach. On the contrary, we have seen that both classes of models successfully fit the data with values of the Hubble constant and matter content in good agreement with some model independent estimates. However, there are some hints that could lead to reject both choices for $f(R)$. Models with power law Lagrangians have a constant $q(z)$ so that they are always accelerating. This is not consistent with the (tentatively) observed transition from acceleration to deceleration at $z_T \\simeq 0.5$. Moreover, a constant $q(z)$ could give rise to problems with nucleosynthesis and structure formation. On the contrary, models with a logarithmic Lagrangian are not affected by such problems and indeed they predicts a transition redshift which is in good agreement with the estimates for the $\\Lambda$CDM model. On the other hand, these models turn out to be too young, i.e. $t_0$ is lower than what is expected.  Actually, a more general remark is in oder here. Let us suppose we have found that a given choice for $f(R)$ leads to models that are in agreement with the data so that we should conclude that this class of models correctly describe the present day universe. What about the early universe? One could expect that the functional expression of $f(R)$ is not changing during the evolution of the universe, even if $R$ may evolve with cosmic time. If this were the case, then the correct choice for $f(R)$ should be the one that leads to models that are not only able to reproduce the phenomenology we observe today, but also give rise to an inflationary period in the early universe. Therefore, we should reject logarithmic Lagrangians since it is well known they do not predict any inflationary period. On the other hand, the choice $f(R) = \\beta R^n$ is able to explain inflation provided one sets $n = 2$, not too far from our estimate in Eq.(\\ref{eq: nrpl}). From this point of view, it is worth noticing that the astrophysical data we have considered probe only the present day universe, while $t_0$ depends on the full evolutionary history. Indeed, logarithmic Lagrangians fail to reproduce the correct $t_0$ in the same way as they fail to give rise to inflation, while both inflation and $t_0$ are correctly predicted by models with power law $f(R)$. This may argue in favour of this choice for $f(R)$, but actually there is no reason to exclude the possibility that also the functional expression of $f(R)$ changes with time so that neither class of models may be definitively rejected or deemed as the correct one from this point of view.  Summarizing, the likelihood analysis presented here allows us to conclude that the Palatini approach to $f(R)$ theories leads to models that are able to reproduce both the SNeIa Hubble diagram and the data on the gas mass fraction in galaxy clusters. From an observational point of view, this means that both power law and logarithmic $f(R)$ are viable candidates to explain the observed accelerated expansion without the need of any kind of dark energy. However, open questions are still on the ground.  First, we have not yet been able to discriminate between the two classes of models. Theoretical considerations and some hints from the age of the universe could argue in favour of the power law $f(R)$, while the observed transition from acceleration to deceleration in the past disfavors this choice. To solve this issue, one has to resort to high redshift probes such as the CMBR anisotropy spectrum. While the data are of superb quality, the underlying theory is still to be developed so that fitting the CMBR anisotropy temperature and polarization spectra with $f(R)$ theories will be  quite a demanding task.  Second, we have only considered two physically motivated and popular choices for $f(R)$. Several other models are possible and are worth of being tested against the data. In particular the $R\\ln R$ Lagrangian which is related to the Straobinsky inflationary model \\cite{starobinsky} and to the limit $R^n\\rightarrow R$ for $n\\simeq 1$ being \\cite{CF} \\begin{equation} R^{1+\\epsilon}=RR^{\\epsilon}=R\\left(e^{\\epsilon\\ln R}\\right)\\simeq R+\\epsilon R\\ln R +{\\cal O}(\\epsilon^2)\\,.\\end{equation}    However, rather than being confused by a plethora of successfull models, it is desiderable to develop a method that allows to directly reconstruct $f(R)$ from the data with as less as possible aprioristic assumptions. This will be the subject of a forthcoming paper \\cite{CCF}.  Last but not least, whether the Palatini approach is indeed the correct method to treat $f(R)$ theories or the metric approach should be preferred is still an unsolved problem. We have shown here that the Palatini approach is not rejected by the data, but a similar analysis for the same models considered in the framework of the metric approach is still lacking. However, it is worth noticing that even this test will not be conclusive. Let us consider, for instance, two choices $f_1(R)$ and $f_2(R)$ and let us suppose that $f_1(R)$ fit the data if considered in the framework of the metric approach, but not if the Palatini approach is used. Let us further assume that the opposite holds for $f_2(R)$. From an observational point of view, it is impossible to select between $f_1(R)$ and $f_2(R)$. Hence, observations could never suggest what is the correct way of performing the variation of a $f(R)$ Lagrangian. The answer to this question is outside the possibilities of an astronomer and lies fully in the field of a theoretician.  As a final comment, we would like to stress the need for synergy between theory and observations. While it is possible to build a physically motivated and mathematically elegant theory, it is not so easy to fit the significant amount of astrophysical data now available. Since the words {\\it observational} and {\\it cosmology} may today be joined together in a single meaningful term ({\\it observational cosmology}), it is time to look at every theoretician's proposal from an observational point of view before drawing any conclusion about the validity of a whatever model. Even if not always conclusive, in our opinion, this is the only way to shed light on the dark side of the universe."
    },
    "0410/astro-ph0410379_arXiv.txt": {
        "abstract": "We present a re--analysis of the European Large Area \\ISO\\ Survey (ELAIS) 90\\,$\\mu$m observations carried out with {\\em ISOPHOT}, an instrument on board the ESA's Infrared Space Observatory ({\\em ISO}). With more than 12\\,deg$^2$, the ELAIS survey is the largest area covered by \\ISO\\ in a single program and is about one order of magnitude deeper than the \\IRAS\\ 100\\,$\\mu$m survey. The data analysis is presented and was mainly performed with the Phot Interactive Analysis software (PIA, Gabriel et al. 1997) but using the pairwise method of Stickel et al. (2003) for signal processing from ERD (Edited Raw Data) to SCP (Signal per Chopper Plateau). The ELAIS 90\\,$\\mu$m catalogue contains 229 reliable sources with fluxes larger than 70 mJy and is available at http://www.blackwell-synergy.com. Number counts are presented and show an excess above the no-evolution model prediction. This confirms the strong evolution detected at shorter(15\\,$\\mu$m) and longer (170\\,$\\mu$m) wavelengths in other \\ISO\\ surveys. The ELAIS counts are in agreement with previous works at 90\\,$\\mu$m and in particular with the deeper counts extracted from the Lockman hole observations (Rodighiero et al. 2003). Comparison with recent evolutionary models show that the models of Franceschini et al. (2001) and Guiderdoni et al. (1998) which includes a heavily-extinguished population of galaxies give the best fit to the data. Deeper observations are nevertheless required to better discriminate between the model predictions in the far-infrared and are scheduled with the Spitzer Space Telescope (e.g. Lonsdale et al. 2003) which already started operating and will also be performed by ASTRO-F (e.g. Pearson et al. 2004). ",
        "introduction": "Strong evolution has been detected in the infrared regime based on \\IRAS\\ number counts at 12, 25, 60 and 100\\,$\\mu$m (Hacking \\& Houck 1987, Hacking, Condon \\& Houck 1987, Hacking   \\& Soifer 1991 , Oliver, Rowan-Robinson \\& Saunders 1992, Bertin,  Dennefeld \\& Moshir 1997) which show an excess of galaxies compared to the no-evolution scenario. These findings were recently confirmed with much deeper surveys carried out with the {\\em ISOPHOT\\/} instrument on-board the Infrared Space Observatory ({\\em ISO}) (Kessler et al. 1996) at 90 and 170\\,$\\mu$m  (Kawara et al. 1998,  Puget et al. 1999,  Efstathiou et al. 2000, Linden-V{\\o}rnle et al. 2000, Juvela, Mattila \\& Lemke 2000, Matsuhara et al. 2000,  Dole et al. 2001). \\ISO\\ also detected a substantial number of faint sources, consistent with strong evolution from 15\\,$\\mu$m number counts (see e.g. Elbaz et al. 1999, Gruppioni et al. 2002). Differential counts obtained from several independent 15\\,$\\mu$m {\\em ISOCAM\\/} surveys show a remarkable upturn at $S_{15}<3$ mJy and an excess of a factor 10 at the faintest flux above the no-evolution predictions.  In addition to the excess of galaxies detected by \\ISO\\ surveys from the mid-infrared to the far-infrared, the observational constraints set by the discovery of the cosmic infrared background (CIB) (see Hauser \\& Dwek 2001 for a review and references therein) together with deep submillimetre surveys (Hughes and Dunlop 1998, Barger et al. 1998, Eales et al. 2000, Scott et al. 2002, Webb et al. 2003) are dramatically increasing the development of new scenarios of galaxy formation and evolution (Pearson \\& Rowan-Robinson 1996, Guiderdoni et al. 1998, Devriendt \\& Guiderdoni 2000, Rowan-Robinson 2001, Franceschini et al. 2001, Takeuchi et al. 2001, Pearson 2001, Wang 2002, Lagache et al. 2003, Xu et al. 2003).  The ELAIS survey (for an overview see Oliver et al. 2000, Paper I) was the largest open time project conducted by {{\\em ISO}}. This survey consists of more than 12\\,deg$^2$ of the sky surveyed at 15 and 90\\,$\\mu$m, nearly 6\\,deg$^2$ at 6.7\\,$\\mu$m and 1\\,deg$^2$ at 175\\,$\\mu$m (i.e. the FIRBACK survey, see Puget et al. 1999) in four high Ecliptic latitude ($|\\beta |$ $>$40$^{\\rm o}$) regions with low \\IRAS\\ 100\\,$\\mu$m sky brightness ($<$ 1.5 MJy sr$^{-1}$). In this work, we present a 90\\,$\\mu$m analysis and source counts limited to the four large areas, three in the northern hemisphere (N1, N2 and N3), and one in the southern hemisphere (S1). Preliminary results of the ELAIS survey at 90\\,$\\mu$m based on the Quick Look analysis and the brightest sources were presented in Efstathiou et al. 2000 (hereafter referred to as Paper III).  The paper is organized as follows: in Section \\ref{sec_obs} we describe the observations and the data reduction based on the analysis of the distribution of consecutive read-outs of the detector instead of using the whole ramp. After the source extraction (Section \\ref{sec_extraction}), we search in vain for solar system objects to remove them from the source list and since some could be useful for calibration purposes. In Section \\ref{src_reliability}, we estimate the completeness of the survey, source flux and position accuracies and the Eddington bias correction from Monte-Carlo simulations of artificial sources on the final maps. The final catalog of sources is presented in Section \\ref{sec_cata}. We compare the {\\em ISOPHOT\\/} calibration for all standard stars observed at 90\\,$\\mu$m with model predictions and for the sources detected in the survey with \\IRAS\\ values (Section \\ref{calib_comparison}). Temperatures from colour ratios between 90 and 170\\,$\\mu$m for sources also detected in the FIRBACK survey are computed in Section \\ref{sec_90_170}. After computing the structure noise (Section \\ref{confusion_noise}) in the ELAIS fields we present number counts (Section \\ref{sec_counts}) which are compared with other works at 90\\,$\\mu$m and to evolutionary models before the summary and discussion of our results in section \\ref{sec_conclusion}. ",
        "conclusions": "\\label{sec_conclusion} We have used a new method to reduce {\\em ISOPHOT\\/} measurements in the 4 main areas of the ELAIS survey at 90\\,$\\mu$m. With a total area of more than 12\\,deg$^2$, the ELAIS survey represents the largest area covered in a single programme with \\ISO.  The relative uncertainty in flux coming from the FCS calibration estimated from the sky background level differences of all rasters is 7 per cent.  On the one hand, the comparison of measured fluxes with models for standard stars shows a strong correlation with a mean ratio of {\\em ISOPHOT\\/} to model values of $1.06\\pm0.02$. On the other hand, the comparison with the \\IRAS\\/FSC catalogue for \\IRAS\\ sources detected in the survey gives a mean ratio of {\\em ISOPHOT\\/} to \\IRAS\\ values equal to 0.76$\\pm$0.17.  Simulations of artificial sources on the final maps spanning a wide range of flux were used to estimate flux and positional uncertainties, completeness and the Eddington bias corrections. The completeness of the survey is about $80\\%$ at 100 mJy.  We present a source list of 237 reliable sources with fluxes larger than 70 mJy, signal-to-noise $\\geq 3$ for the 4 large ELAIS fields. The full version of the catalogue is available at http://www.blackwell-synergy.com.  Sources detected at 90 and 170\\,$\\mu$m in the FIRBACK survey (Dole et al. 2001) have an average colour temperature of $T_{C} = 19K$ with all sources lying in the range 13-25K in agreement with Stickel et al. (2001) in the \\ISO\\ Serendipity Survey (Stickel et al. 1998).  The ELAIS counts extend the \\IRAS\\ counts by more than one order of magnitude in flux and show significant departure from the no-evolution model as detected in other \\ISO\\ surveys from the mid- to the far-infrared.  There is in general a good agreement between ELAIS and other 90\\,$\\mu$m source counts and in particular with the deeper counts measured in the Lockman Hole (Rodighiero et al. 2003). Differential number counts measured in the ELAIS regions at 90\\,$\\mu$m are compared to recent evolutionary models. Among few models which were compared to our counts, the model of Franceschini et al. (2001) and the scenario E of Guiderdoni et al. (1998) give the best agreement with the observations.  However, the latter model is a factor  of $\\sim 2.5$ {\\it below} the counts measured at 170\\,$\\mu$m in two of the ELAIS regions (Dole et al. 2001) related to the present paper. On the other hand Matsuhara et al. (2000) found that the scenario E model prediction of Guiderdoni et al. is in close agreement with the 170\\,$\\mu$m number counts in the small area of the ``Lockman Hole'' (see also Kawara et al. 1998) but their 90\\,$\\mu$m integral counts are significantly above the model.  The nature and redshift distributions of the ELAIS galaxies can test the various models and their hypothesis e.g. distinguishing the different galaxy populations on which these models are built. This will also help to clarify the origin of the differences seen in the number counts of the various \\ISO\\ surveys at different wavelengths.  The 90\\,$\\mu$m luminosity function is presented in Serjeant et al. (2004) and Rowan-Robinson et al. (2004) present results based on the ELAIS final-band merged catalogue combining the \\ISO\\ and ground-based observations in the ELAIS fields."
    },
    "0410/astro-ph0410653_arXiv.txt": {
        "abstract": "Recent observations suggest that the behaviour of tracer species such as $\\mathrm{N_2H^+}$ and $\\mathrm{CS}$ is significantly different in regions of high and low mass star formation. In the latter, $\\mathrm{N_2H^+}$ is a good tracer of mass, while $\\mathrm{CS}$ is not. Observations show the reverse to be true in high-mass star formation regions. We use a computational chemical model to show that the abundances of these and other species may be significantly altered by a period of accelerated collapse in high mass star forming regions. We suggest these results provide a potential explanation of the observations, and make predictions for the behaviour of other species. ",
        "introduction": "Molecular lines are effective probes of physical conditions in star forming regions. The abundances of molecules and their line profiles give information on the physical state of the gas and its likely evolution during the collapse to form a star \\citep{b10}. Methods based on molecular lines have been used together with dust continuum emission to probe regions of low-mass and of high-mass star formation \\citep{b9}. Recent work has shown that tracer species may be significantly depleted in the central parts of dense cores. In low-mass star forming regions, CS and CO do not trace the gas mass in the central cores \\citep{b1,b3,b11,b12,b2}, but $\\mathrm{N_2H^+}$ is an excellent tracer of the dust continuum emission \\citep{b19,b4,b3}. One possible explanation, therefore, is that the $\\mathrm{N_2}$ molecules which drive the nitrogen chemistry do not freeze-out on to dust grains as readily as some other species, and the species derived from $\\mathrm{N_2}$, notably $\\mathrm{N_2H^+}$, remain present in the gas to higher densities and trace its mass, as does the dust continuum emission.  In high-mass star forming regions, however, the situation is unclear. The epoch most readily studied is the so called hot core phase, in which a characteristic chemistry signifies the operation of processes unlike those in quiescent clouds. This chemistry arises from the evaporation of molecules that had been frozen out on dust grains during the long cold period of collapse; they are then released back to the gas phase when the newly formed massive star warms its environment so that the ices sublimate.  Recent molecular line observations \\citep{b8,b9,b14} of high-mass-star forming cores show differences between the CS and $\\mathrm{N_2H^+}$ distributions.  In a survey of sources in regions of high-mass star formation, Pirogov et al. find that the CS and dust emission distributions are very similar, and differ significantly in many cases from the $\\mathrm{N_2H^+}$ distributions; the $\\mathrm{N_2H^+}$ abundance falls near massive young stellar objects, on scales up to a parsec where the temperature, relative to the rest of the cloud and to the CS peak, does not change significantly. The variation in relative intensity can be about one order of magnitude, larger than could be accounted for by excitation or opacity effects, and is seen in approximately one-third of the sources. {\\bf A detailed study} is, however, expected to reveal significant differences in the majority of the sources. Zinchenko et al. have also discussed the behaviour of some other molecules in dense cores in high-mass star forming regions. In summary, in contrast to its behaviour in low-mass star-forming regions where $\\mathrm{N_2H^+}$ traces mass very well, in regions of high-mass star formation $\\mathrm{N_2H^+}$ is a poor tracer of mass.   Here, we focus on the exciting new results concerning CS and $\\mathrm{N_2H^+}$. Why is there a complete reversal of the behaviour between regions of formation of high-mass stars (in which CS is apparently a good tracer of core mass) and of low mass stars (in which $\\mathrm{N_2H^+}$ is apparently a good tracer of core mass)? We offer a possible explanation in terms of the differing dynamical state of the two types of region, and with a detailed computational model of high-mass star formation we obtain results that are consistent with the observations. ",
        "conclusions": "Recent observational data on cores in regions of high-mass star formation show the exciting and unexpected results opposite in character to that found in regions of low-mass star formation. We have suggested that the enhancement of CS and reduction in $\\mathrm{N_2H^+}$ abundance found in regions of high-mass star formation may be related to the high dynamical activity in these regions which could enhance the rate of collapse of cores above the free-fall rate. Consequently, high gas densities would be achieved before freeze-out had removed the molecules responsible for $\\mathrm{N_2H^+}$ loss, while the high densities promote CS formation. We have demonstrated using a computational chemical model that the predicted effects are indeed recovered for plausible parameter choices. We have shown that these results are not particularly sensitive to the particular enhancement rate adopted, above a factor of twice the free-fall rate. The observational results also give some information about variations of other molecules with respect to high density tracers. Our model confirms that HCN, for example, is similar to CS in its distribution, as found by Pirogov et al. However, our model predicts SO would have a distribution similar to $\\mathrm{N_2H^+}$, whereas the observations of S255 show it to be much closer to CS. (See Table 2). {\\bf As yet, data on the distribution of SO abundance is available only for S255 and so it is not clear whether this result is significant or restricted to this source.} As with most molecules, the effect of accelerated collapse on SO will be both to increase the rate of production and to prevent the removal of other possible reaction partners from the gas phase. This discrepancy between theory and observation for SO suggests that our classification of molecules into two classes of species, each of which is primarily affected by only one of these mechanisms, is over-simplistic for molecules at the centre of complex reaction networks. Should further SO observations confirm the S255 results, a detailed analysis of the behaviour of sulfur-bearing species in our model will be necessary.   The temperatures found in these cores are somewhat warmer than normally found in dense cores. We have examined the results from computational chemical models in which the initial warming stages of the hot core phase is permitted, up to temperatures of around 30 - 50 K, according to the prescription contained in Viti et al.(2004) The results reported here are not significantly changed.  We conclude that the results of the model investigated are consistent with the rather limited observational data.  This first investigation of the effects of accelerated collapse on the chemical abundances of species in high mass star formation regions suggests a possible explanation of the observations considered, and provides an important route for future observational and theoretical work."
    },
    "0410/astro-ph0410465_arXiv.txt": {
        "abstract": "{A conceptually novel model - one in which stimulated ($i.e.$ induced) resonance Rayleigh scattering causes ions of select species to become accelerated to very high terminal velocities in the solar winds of OB-type stars - is proposed to explain the general appearance of so-called P Cygni profiles that often dominate the vacuum ultraviolet (VUV) spectra of such stars. In the unit step of the proposed scattering process, a radially outgoing photon (frequency $\\nu _1 )$ that is arbitrarily blueshifted with respect to the rest frame frequency $\\nu _o $ of the resonance line associated with a P Cygni profile is absorbed from the illuminating star's continuum, a new photon (frequency $\\nu _2 )$ that is very close to $\\nu _o $ and that propagates in the \\textit{backwards }direction ($i. e.$ towards the star) is created, and the outwardly directed velocity $v$ of the ion being accelerated is increased by an amount $\\Delta v\\approx 2h\\nu _o /cm_I ,$ with all three events occurring simultaneously. A monochromatic wave at $\\nu _2 $ initially forms at some distance from the star, and becomes enormously amplified via the stimulated scattering process as it propagates radially inwards towards the star, all the while retaining a relatively high degree of monochromaticity. The high rate of stimulated scattering enables ions of the resonant species in the solar wind to become accelerated to terminal velocities as high as $\\sim $1000 km/sec. Since stimulated scattering processes are characterized by pump power thresholds, the model readily explains why only select species are accelerated to very high velocities in a given star's solar wind, and also why a dramatic P Cygni profile for a given ion species can often discontinuously be present or absent when spectra of stars varying only slightly in spectral type are compared. ",
        "introduction": "So-called P Cygni profiles of select ion resonances ($e.g.$ C IV $\\lambda \\lambda $1548, 1551; Si IV $\\lambda \\lambda$1394, 1403; NV $\\lambda \\lambda $1239, 1243 -- see Fig. 1) are often the dominant features in the vacuum ultraviolet (VUV) spectra of O-type (Walborn $et$ $al.$ \\cite{walborn85}) and B-type (Walborn $et$ $al.$ \\cite{walborn95}) stars, especially giants and supergiants. These features were originally discovered in the late nineteen-sixties by Morton (Morton \\cite{morton67}) and his colleagues (Morton $et$ $al.$ \\cite{morton69}) with the use of rocket-based instrumentation, and have since been studied extensively by many astronomers utilizing several orbiting satellites that were successfully launched in the years that followed. Especially productive in obtaining P Cygni spectral data were the fully dedicated \\textit{Copernicus} and \\textit{International Ultraviolet Explorer (IUE)} satellites, and the \\textit{Space Telescope Imaging Spectrometer (STIS)} aboard the \\textit{Hubble Space Telescope.}  In a typical P Cygni profile, the star's continuum level in a relatively broad spectral region that is \\textit{blueshifted} with respect to a given ion resonance line $\\nu _o $ is heavily absorbed. (Frequently the ion resonances are doublets, as is the case with the three prominent P Cygni profiles shown in Fig.~\\ref{fig1}.) The measured frequency displacement from $\\nu _o $ of the short wavelength limit $\\nu _B $ of this absorbed region is normally used by astronomers to determine the maximum velocity $v_{\\max } $ of the corresponding ions in the solar wind via the standard relationship $\\nu _B -\\nu _o =\\nu _o v_{\\max } /c.$ From such Doppler-shift-based analyses of P Cygni profiles, astronomers have inferred that the corresponding ions are strongly accelerated radially away from the stars, reaching in some cases maximum velocities as high as two or three thousand km/sec. For a hot star, the maximum velocity deduced in this manner is frequently greater than the calculated escape velocity $V_e =\\left( {2GM/R} \\right)^{1/2}$ at the radius $R$ of the photosphere, implying an actual ejection of matter from the star into interstellar space.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig1.eps}} \\caption{VUV spectrum of the O9 supergiant HD 210809 recorded with the \\textit{Space Telescope Imaging Spectrometer (STIS) }aboard the \\textit{Hubble Space Telescope}. (Spectrum downloaded from the \\textit{MAST Scrapbook }interactive web site.) The three P Cygni profiles dominating the spectrum are (from left to right) due to N V, Si IV, and C IV ions, respectively. The absorption at 1216 {\\AA} is that of H-atom Lyman alpha. A multitude of narrow interstellar absorption lines are present. Wider stellar photospheric absorptions can also be seen. } \\label{fig1} \\end{figure}  Included in Morton (\\cite{morton76}) are special \\textit{Copernicus }VUV spectral scans of the O4 If supergiant $\\zeta $ Pup, recorded at very high spectral resolution (0.051 {\\AA} FWHM), and with an exceptionally high signal-to-noise ratio. For this very hot star, which has $R/R_\\odot \\approx $ 20.3 and $M/M_\\odot \\approx $ 100, it was calculated in Morton (1976) that $V_e =$1370 $\\pm $ 160 km/sec. Between 920 {\\AA} and 1750 {\\AA}, 13 different ion species were found to display P Cygni profiles in $\\zeta $ Pup, many of them appearing very strong. For example, the ion C III was observed to have a strong P Cygni profile associated with its resonance line at $\\approx $ 977 {\\AA}. The blue edge of the heavily absorbed region appeared shifted from 977 {\\AA} by $\\approx $ 2710 km/sec, prompting the conclusion that was made in Morton (1976) that C III ions in the solar wind of $\\zeta $ Pup are continually being ejected into space. However, as will be shortly pointed out, according to the nonlinear P Cygni model here being proposed, use of the relationship $\\nu _B -\\nu _o =\\nu _o v_{\\max } /c$ to determine the maximum ion velocity $v_{\\max } $ attained in a star's solar wind results in an overestimation of this quantity by almost exactly a factor 2.  As noted in Lamers {\\&} Cassinelli (\\cite{lamers}), current stellar wind theories fall into three broad classes: radiative ($i.e.$ line-driven) models, coronal models, and hybrid models. In radiative models, transfer of photon momentum to the gas is assumed to occur through the opacity of the many strong VUV resonance lines that are present. Increased acceleration is believed to result from the progressive Doppler shifting of the line opacity into the unattenuated photospheric radiation field. The stellar wind in existing radiative models is thus assumed to be driven by a purely \\textit{linear} effect, commonly known as \\textit{radiation pressure}.  In the model here being proposed, a \\textit{nonlinear} photomechanism drives the acceleration of the fast moving ions in the solar wind, simultaneously producing the P Cygni spectral profiles one sees in the line-of-sight to the star. Perhaps the most attractive feature of such a nonlinear mechanism would be that it would possess a definite pump power threshold. Only those ion species for which this threshold is reached would be accelerated to very high velocities in a given star's solar wind. Such selectivity is very hard to explain with linear radiative models for solar winds.  In the present paper the focus is placed upon the basic physics of the proposed nonlinear model. No attempt is made to account in detail for fine-grained features appearing in P Cygni profiles of specific stars. The paper is organized as follows. In Sect. 2, C II $\\lambda \\lambda $1334, 1336 ion resonance spectra recorded in four different stars of roughly the same spectral type are compared to demonstrate the abrupt manner in which P Cygni profiles often appear in star spectra, thus strongly suggesting that a threshold effect of some kind must be involved. In Sect. 3, the proposed nonlinear scattering mechanism is outlined. Undoubtedly of some significance is the fact that this mechanism is easily shown to be consistent with the conservation of both energy and momentum. Simple consideration of the manner in which spontaneous resonance Rayleigh scattering ($i.e.$ elastic scattering) would occur in the rest frame of an ion being accelerated in a star's solar wind directly leads to identification of the proposed nonlinear scattering mechanism with stimulated ($i.e.$ induced) resonance Rayleigh scattering. In the unit step of this stimulated process, three distinct events \\textit{simultaneously} occur. (\\ref{eq1}) An outwardly propagating photon from the star's continuum that is blueshifted with respect to the rest frame frequency $\\nu _o $ of a strong ion resonance line becomes absorbed by the ion. (\\ref{eq2}) A photon at exactly $\\nu _o $ that propagates radially backwards ($i. e. $ towards the star) is emitted by the ion, contributing to the intensity of a backward propagating light wave at $\\nu _o $ that is already present at the position of the ion. (For analytical purposes, this wave is termed the $\\nu _2 $ wave, but its frequency is $\\approx \\nu _o .)$ (\\ref{eq3}) The outwardly directed radial velocity $v$ of the ion is increased by $\\Delta v\\approx 2h\\nu _o /cm_I $, \\textit{irrespective} of the value of $v$.  Since, at any point in the solar wind, the transition probability for the unit step to occur is proportional to the product of the outwardly propagating continuum flux and the inwardly propagating light beam intensity $I_2 $, the likely occurrence of a stimulated scattering regime is therefore here expected, with the $\\nu _2 $ wave continuously gaining in intensity as it propagates inwardly towards the star. During the entire time it undergoes amplification, the $\\nu _2 $ wave remains relatively monochromatic, a result of the unit step in the scattering process becoming stimulated by the same wave. However, as the $\\nu _2 $ wave propagates inwardly towards the star, the frequency of the pump light that effectively drives the stimulated scattering process continually changes. In this way, continuum light from the star spanning a broad spectral range can contribute to the intensity of the monochromatic $\\nu _2 $ wave as the latter impinges upon the photosphere of the star.  In Sect. 4, equations that should describe the stimulated scattering process are briefly discussed. However, efforts to provide numerical solutions for these equations have not yet been attempted.  In the next section of the present paper (Sect. 5), two types of spectroscopic evidence that could help to substantiate the nonlinear model are discussed -- fluorescence spectra and two-photon absorption bands. Most P Cygni profiles are observed to possess a prominent fluorescence component. Via simple extension of the stimulated Rayleigh scattering model used to explain P Cygni absorption, one can also comprehend how P Cygni fluorescence is excited. By considering the effect of having the absorption and the fluorescence excitation mechanisms operate in tandem, one can then easily account for the observed fact that the strongest apparent absorption in a P Cygni profile occurs at frequencies considerably blueshifted from $\\nu _o .$  In Sect. 5 it is also suggested that one might seek to identify Doppler broadened absorption bands representing resonantly enhanced two-photon absorption of outwardly propagating continuum light from a star by various ion species present in its photosphere, or just outside. Such two-photon absorption could in principle be induced by the powerful, inwardly propagating, monochromatic beams at the P Cygni ion rest frame frequencies $\\nu _o $, provided that near resonances exist between some of the frequencies $\\nu _o $ and some transitions of the two-photon-absorbing ions. Details of a search for such near resonances in the case of the P Cygni star $\\zeta $ Pup are described. However, with one possible intriguing exception, this search did not result in the clear identification of any two-photon absorption bands in this star.  In Sect. 6 nonlinear and linear P Cygni mechanisms are briefly compared. ",
        "conclusions": ""
    },
    "0410/astro-ph0410186_arXiv.txt": {
        "abstract": "Interaction of cosmological gamma ray burst radiation with the dense interstellar medium of host galaxy is considered. Gas dynamical motion of interstellar medium driven by gamma ray burst is investigated in 2D approximation for different initial density distributions of host galaxy matter and different total energy of gamma ray burst. The maximum velocity of motion of interstellar medium is $1.8\\cdot10^4$ km/s. Light curves of gamma ray burst afterglow are calculated for set of non homogeneous density, distribution gamma ray burst total energy, and different viewing angles. Spectra of gamma ray burst afterglow are modeled taking into account conversion of hard photons (soft X-ray, hard UV) to soft UV and optics photons. ",
        "introduction": "Although gamma-ray bursts (GRBs) were discovered more then thirty years ago \\cite{Kleb73}, their origin is still unclear. The most extensive data on the detection of GRBs have been obtained by the Compton Gamma Ray Observatory BATSE experiment \\cite{Briggs95, Fishman95,Meegan92}. Analysis of the GRBs detected showed that their apparent distribution on the sky was isotropic, but that there was a significant departure of their (logN -- logS) curve from the $N\\sim S^{-3/2}$, low corresponding to a spatially uniform source distribution \\cite{Briggs95,Kouv94}.  Observation of optical afterglows  of GRB, following after identification of GRB with a transient X-ray source by Beppo-SAX, and discovery of large (up to $z=4.5$) redshifts in the spectra of optical transients had confirmed the cosmological origin of long GRB.  The cosmological model suggests that the GRB source are located in distant galaxies (within $\\sim 10^3$ Mpc). In the framework of this model, the observed fluxes $\\sim 10^{-4}$ erg s$^{-1}$ require the release of enormous amounts of energy ($\\sim 10^{51} - 10^{53}$ erg) within a fairly short time interval (of the order of several tens of seconds). Such a powerful energy release should have a strong effect on a large volume of matter around the parent galaxy and should give rise to the formation of GRB counterpart at other wavelengths. Without specifying the mechanism for GRB formation, we assume only the existence of a strong flux gamma radiation and consider the interaction of this radiation with the interstellar medium on large spatial and temporal scales.  We investigate the response of an interstellar medium of standard chemical composition to the passage of a short-term powerful pulse of gamma radiation (that is, the dynamical behaviour and radiative cooling after heating by the gamma rays).  The spherical symmetric model was investigated by Bisnovatyi-Kogan and Timokhin (1997), 2D model, permitting to study different matter distribution and reaction to anizotropic GRB had been studied by Barkov and Bisnovatyi-Kogan (2004a, 2004b) using numerical simulations by PPM method. In what follows we represent the results from Barkov and Bisnovatyi-Kogan (2004a, 2004b). ",
        "conclusions": "Optical afterglow, connected with the reradiation  of the GRB by the dense enough molecular cloud could be observed as optical transient. Indications that GRB explosions take place in the region of star formation filled with dense gas clouds \\cite{Sokol,Pacz} make this possibility as very probable. Observation of plato in the optical afterglow of GRB 030329 during a month between 64 and 94 days after GRB detection \\cite{Poz} may be connected with such kind of reradiation.  The dense molecular clouds have very low temperature, at which dust is formed. Estimations made by Barkov and Bisnovatyi-Kogan (2004b) show, that dust is evaporated by GRB pulse in the cloud near GRB impulse direction, or in the whole cloud by isotropic GRB with $R\\sim 1$ pc, and does not influence the light curve of optical, UV and X-ray GRB afterglows."
    },
    "0410/astro-ph0410715_arXiv.txt": {
        "abstract": "We consider a stationary axisymmetric force-free degenerate magnetosphere of a rotating Kerr black hole surrounded by a thin Keplerian infinitely-conducting accretion disk. We focus on the closed-field geometry characterize by a direct magnetic coupling between the disk and the hole's event horizon. We first argue that the hole's rotation necessarily limits the radial extent of the force-free link on the disk surface: the faster the hole rotates, the smaller the magnetically-connected inner part of the disk has to be. We then show that this is indeed the case by solving numerically the Grad--Shafranov equation---the main differential equation describing the structure of the magnetosphere. An important element in our approach is the use of the regularity condition at the inner light cylinder to fix the poloidal current as a function of the poloidal magnetic flux. As an outcome of our computations, we are able to chart out the maximum allowable size of the portion of the disk that is magnetically connected to the hole as a function of the black hole spin. We also calculate the angular momentum and energy transfer between the hole and the disk that takes place via the direct magnetic link. We find that both of these quantities grow rapidly and that their deposition becomes highly concentrated near the inner edge of the disk as the black hole spin is increased. ",
        "introduction": "\\label{sec-intro}  This paper is devoted to the subject of magnetic interaction between a rotating black hole and an accretion disk around it---a topic that has enjoyed a lot of attention among researchers in recent years. Magnetic fields are believed to play an important role in the dynamics of accreting black hole systems (e.g., Begelman, Blandford, \\& Rees~1984; Krolik~1999b; Punsly~2001). In particular, they can be very effective in transporting angular momentum and the associated rotational energy of either the hole or the disk.  Where and how this transport takes place and to what observational consequences it can lead, is partly determined by the global geometry of the magnetic field lines. Conceptually, one can think of two basic types of geometry. The first type is the open-field configuration shown schematically in Figure~\\ref{fig-geometry-open}. The main topological feature here is that there is no direct magnetic link between the hole and the disk. Instead, all the field lines are open and extend to infinity. Historically, this configuration was the first to have been considered, and it has been studied very extensively during the past three decades (see, e.g., Lovelace~1976; Blandford~1976; Blandford \\& Znajek~1977, hereafter BZ77; MacDonald \\& Thorne~1982, hereafter MT82; Phinney~1983; Macdonald~1984; Thorne, Price, \\& Macdonald~1986; Punsly 1989, 2001, 2003, 2004; Punsly \\& Coroniti 1990; Beskin \\& Par'ev 1993; Beskin~1997; Ghosh \\& Abramowicz 1997; Beskin \\& Kuznetsova 2000; Komissarov~2001, 2002b, 2004a). The reason for this popularity is that this configuration is related to the famous Blandford--Znajek mechanism (BZ77) now widely regarded as the primary process powering jets in active galactic nuclei (AGN) and micro-quasars. As Blandford and Znajek showed, a large-scale, ordered open magnetic field can extract the rotational energy from a spinning black hole and transport it to large distances via Poynting flux (a similar process works along the field lines connected to the disk).  The second type of magnetic field geometry is the closed-field configuration, shown in Figure~\\ref{fig-geometry-closed}. Although it has been occasionally discussed in the literature before the last decade (e.g., Zeldovich \\& Schwartzman, quoted in Thorne~1974; MT82; Thorne~et~al. 1986; Nitta, Takahashi, \\&  Tomimatsu~1991; Hirotani~1999), it is only in the last five years that it has attracted serious scientific attention (e.g., Blandford~1999, 2000, 2002; Gruzinov~1999; van~Putten 1999; van~Putten \\& Levinson 2003; Li~2000, 2001, 2002a, 2002b, 2004; Wang, Xiao, \\& Lei 2002; Wang, Lei, \\& Ma 2003a; Wang et al. 2003b, 2004). The basic topological structure of magnetic field in this configuration is very different from that of the open-field configuration. The field lines are closed and directly connect the black hole to the disk. In this so-called Magnetically-Coupled configuration (Wang et al. 2002), the energy and angular momentum are not taken away to infinity, but instead are exchanged between the hole and the disk by the magnetic field. Therefore, magnetic coupling, together with the accretion process, controls the spin evolution and the spin equilibrium of the black hole (Wang et al. 2002, 2003a). In addition, the rotational energy of the hole can be magnetically extracted (just like in the Blandford--Znajek process) and deposited onto the disk leading to a change in the disk energy-dissipation profile and hence its observable spectral characteristics (Gammie 1999; Li 2000, 2001, 2002a, 2002b, 2004; Wang et al. 2003a,b). Finally, if the rotating field lines are strongly twisted and become unstable to a non-axisymmetric kink-like instability, a strong variability of the energy release may result, which would be a possible explanation for quasi-periodic oscillations (QPOs) in micro-quasar systems (e.g., Gruzinov~1999; Wang et al.~2004). All these phenomena make the closed-field configuration astrophysically very interesting.  Most of the work that has been done on studying the magnetic field structure around an accreting Kerr black hole, including the seminal paper by Blandford \\& Znajek (BZ77), has been performed under the assumption that the magnetosphere above the thin disk is ideally conducting and force-free. Then, if one also assumes that the system is stationary and axisymmetric, the magnetic field is governed by the general-relativistic version of the force-free Grad--Shafranov equation (e.g., MT82; for the full-MHD generalization of this equation see Nitta~et~al. 1991; Beskin \\& Par'ev 1993; Beskin~1997). Since this is a rather nontrivial nonlinear partial differential equation (PDE) with singular surfaces and free functions, it is generally not tractable analytically, except in some special simple cases, such as the slow-rotation limit (BZ77). However, over the past 20 years, a number of force-free solutions for the magnetosphere have been obtained numerically, either by solving the Grad--Shafranov equation itself (MacDonald~1984; Fendt~1997) or as an asymptotic steady state of force-free degenerate electrodynamics (FFDE) evolution (Komissarov~2001, 2002b, 2004a). Until now, most of these studies have been done in the context of the open-field configuration, primarily because of its relevance to the jet problem.  In contrast, most of the work on closed-field configurations has been limited to analytic and semi-analytic studies of the effects that magnetic link has on the disk radiation profile and on the spin evolution of the black hole. The structure of the magnetosphere has not in fact been computed self-consistently. These studies have just assumed the existence of the link and made some simplified assumptions about the field distribution on the horizon.  The only exception to this deficiency is the recent work by Uzdensky (2004) where a force-free magnetosphere linking a Keplerian disk to a Schwarzschild black hole has been numerically computed for the first time.  In the present paper, we make the next logical step by extending this previous work to the more general case of a rapidly rotating Kerr black hole. This is indeed the most important case, not only because real astrophysical black holes are believed to be rotating, but also because the nonlinear terms in the Grad--Shafranov equation, especially the toroidal field pressure, become large in this case. As a result, even the existence of closed-field solutions is not guaranteed. And indeed, one of the main goals of our present study is to determine the conditions for existence of such solutions in Kerr geometry. In other words, we aim at determining the limitations that the rotation of the black hole imposes on the direct magnetic link between the hole and the disk. In addition, by computing the global magnetic field structure, we will be able to study the effect of the black hole rotation on the magnetic field distribution on the horizon, the poloidal electric current as a function of poloidal magnetic flux, and the location of the inner light cylinder, as well as such astrophysically-important processes as angular momentum and energy transfer between the hole and the disk.  In order to achieve the goal of obtaining numerical solutions of the force-free Grad--Shafranov equation in Kerr goemetry, we first analize the mathematical structure of this equation. In particular, we pay special attention to its singular surfaces (the event horizon and the light cylinder) and the corresponding regularity conditions. Thus, we use the light-cylinder regularity condition to determine the poloidal current as a function of poloidal magnetic flux, similar to way it was done by Contopoulos, Kazanas, \\& Fendt (1999) for the case of the pulsar magnetosphere (see also Beskin \\& Kuznetsova 2000; Uzdensky~2003; Uzdensky~2004). The event-horizon regularity condition, also known as Znajek's (1977) horizon boundary condition, is then used to determine the poloidal flux distribution on the horizon. Thus, one does not have the freedom to arbitrarily specify any extra boundary conditions at the horizon, and hence there is no problem with causality, in line with the reasoning presented by Beskin \\& Kuznetsova (2000) and by Komissarov (2002b, 2004a) (see also Levinson~2004).  Finally, although in this paper we deal exclusively with large-scale, ordered magnetic fields, we acknowledge the difficulty in justifying the existence of such fields around accreting black holes (e.g., Livio, Ogilvie, \\& Pringle 1999), especially in the closed-field configuration. Also, as recent numerical simulations (e.g., Hawley \\& Krolik 2001; Hirose et al.~2004), there may be a significant deposition of energy and angular momentum at the inner edge of the disk due to small-scale, intermittent magnetic fields connecting the disk to the plunging region (see also Krolik~1999a; Agol \\& Krolik~2000).  The paper is organized as follows. \\S~\\ref{sec-equations} describes the mathematical formalism of force-free axisymmetric stationary magnetospheres in Kerr geometry. In particular, in \\S~\\ref{subsec-kerr} we introduce the Kerr metric tensor in Boyer--Lindquist coordinates and list several general geometric relationships for the future use. In \\S~\\ref{subsec-GS-eqn} we consider steady-state, axisymmetric, degenerate electro-magnetic fields and then discuss the force-free condition and the Grad--Shafranov equation. In \\S~\\ref{subsec-EH} we consider the black hole's event horizon as a singular surface of this equation and discuss the associated regularity condition, which is also known as Znajek's horizon boundary condition. In \\S~\\ref{sec-idea} we present a simple but robust physical argument that demonstrates that a force-free magnetic link between a rotating black hole and the disk cannot extend to arbitrarily large distances on the disk, we also argue that the maximal radial extent of the magnetic link should scale inversely with the black hole's rotation rate in the slow-rotation limit. We confirm these propositions in \\S~\\ref{sec-numerical}, where we present our numerical solutions of the grad--Shafranov equation. Then, in \\S~\\ref{sec-implications} we discuss the magnetically-mediated angular-momentum and energy exchange between the hole and the disk. We then close by summarizing our findings in \\S~\\ref{sec-conclusions}. ",
        "conclusions": "\\label{sec-conclusions}  In this paper we investigated the structure and the conditions for the existence of a force-free magnetosphere linking a rotating Kerr black hole to its accretion disk. We assumed that the magnetosphere is stationary, axisymmetric, and degenerate and that the disk is thin, ideally conducting, and Keplerian and that it is truncated at the Innermost Stable Circular Orbit. Our main goal was to determine under which conditions a force-free magnetic field can connect the hole directly to the disk and how the black hole rotation limits the radial extent of such a link on the disk surface.  We first introduced (in \\S~\\ref{sec-idea}) a very simple but robust physical argument that shows that, generally, magnetic field lines connecting the polar region of a spinning black hole to arbitrarily remote regions of the disk cannot be in a force-free equilibrium. The basic reason for this can be described as follows. Since the field lines threading the horizon have to first cross the inner light cylinder, and since they generally rotate at a rate that is different from the rotation rate of the black hole, these field lines have to be bent somewhat. In other words, they develop a toroidal magnetic field component, just like the open field lines crossing the outer light cylinder in a pulsar magnetosphere. In the language of the Membrane Paradigm (see Thorne~et~al. 1986), this toroidal field is needed so that the field lines could slip resistively across the stretched event horizon. The next step in our argument is to look at those field lines that connect the polar region of the horizon to the disk somewhere far away from the black hole. In a force-free magnetosphere, toroidal flux spreads along field lines to keep the poloidal current $I\\sim B_{\\hat{\\phi}}\\alpha\\varpi$ constant along the field. Then one can show that the outward pressure of the toroidal field generated due to the black hole rotation turns out to be so large that it cannot be confined by the poloidal field tension at large enough distances. In other words, the field lines under consideration cannot be in a force-free equilibrium. Furthermore, one can generalize this argument to the case of closed magnetospheres of finite size and derive a conjecture that the maximal radial extent $R_{\\rm max}$ of the magnetically-coupled region on the disk surface should scale inversely with the black hole spin parameter~$a$ in the limit $a\\rightarrow 0$.  In order to verify this hypothesis and to study the detailed structure of magnetically-coupled disk--hole configurations, we have obtained numerical solutions of the general-relativistic force-free Grad--Shafranov equation corresponding to partially-closed field configurations (shown in Fig.~\\ref{fig-geometry-kerr}). This is an nonlinear 2nd-order partial differential equation for the poloidal flux function~$\\Psi(r,\\theta)$ and it is the main equation governing the system's behavior.  An additional complication in this problem arises from the need to specify two free functions that enter the force-free Grad--Shafranov equation; these are the field-line angular velocity $\\Omega_F(\\Psi)$ and the poloidal current~$I(\\Psi)$. Because all the field lines are assumed to be frozen into the disk, the first of this functions is determined in a fairly straightforward way. Namely, for any given field line~$\\Psi$, $\\Omega_F(\\Psi)$ is just the Keplerian angular velocity at this line's footpoint on the disk. Specifying the poloidal current, on the other hand, is a much more difficult and nontrivial task. The reason for this is that it cannot be just prescribed explicitly on any given surface and one should look more thoroughly into the mathematical nature of the Grad--Shafranov equation itself to determine~$I(\\Psi)$. In particular, the most important feature of the Grad--Shafranov equation in this regard is that it becomes singular on two surfaces, the event horizon and the inner light cylinder. This observation is very useful because one can impose a physically-motivated regularity condition at each of these surfaces. One of the most important ideas in our analysis is that one can use the light-cylinder regularity condition to determine, using an iterative procedure, the poloidal current $I(\\Psi)$, similar to the way it was done by CKF99 in the context of pulsar magnetospheres.  As for the singularity at the event horizon, it is also very important. Basically, it tells us that it is not possible to prescribe an arbitrary boundary condition at the horizon; instead, one can only impose a certain physical condition of regularity there. When combined with the Grad--Shafranov equation itself, this regularity requirement results in a single relationship (historically known as the horizon boundary condition, first derived by Znajek 1977) between three functions: the horizon flux distribution $\\Psi_0(\\theta)$, and the two free functions, $\\Omega_F(\\Psi)$ and~$I(\\Psi)$ (e.g., Beskin~1997; Beskin \\& Kuznetsova~2000). What's important is that there are no other independent relationships that can be specified on this surface. In practical terms, this means that this condition should be used to determine the function $\\Psi_0(\\theta)$ in terms of $\\Omega_F(\\Psi)$ and~$I(\\Psi)$, which therefore must be determined outside the horizon. This fact helps to alleviate some of the causality issues raised by Punsly (1989, 2001, 2003) and by Punsly \\& Coroniti (1990).  Since one of the goals of this work was to study the dependence $R_{\\rm max}(a)$, we performed a series of computations corresponding to various values of two parameters: the black hole spin parameter $a$ and the the magnetic link's radial extent $R_s$ on the disk surface (the field lines anchored to the disk beyond $R_s$ were taken to be open and non-rotating). At the same time, the disk boundary conditions were kept the same in all these runs, namely, $\\Psi_d(r)=\\Psi_{\\rm tot} r_{\\rm in}(a)/r$. Therefore, varying the value of $R_s$ for fixed $a$ was equivalent to varying the amount $\\Psi_s$ of open magnetic flux threading the disk.  Whereas for some pairs of values of $a$ and $\\Psi_s$ we were able to achieve a convergent force-free solution, for others we were not. Thus, as one of the main results of our computations, we were able to chart out the allowed and the forbidden domains in the two-parameter space $(a,\\Psi_s)$. The boundary between these two domains is a curve $a_{\\rm max}(\\Psi_s)$, which can be easily remapped into the curve $R_{\\rm max}(a)$. As can be seen in Figure~\\ref{fig-a_of_Psi_s}, this is a monotonically rising curve with the asymptotic behavior $a_{\\rm max}\\propto \\Psi_s$ as $\\Psi_s\\rightarrow 0$, which is in line with our predictions.  We also computed the total angular momentum and red-shifted energy exchanged in a unit of global time~$t$ between the hole and the disk through magnetic coupling. We studied the dependence of these quantities on the black hole spin parameter~$a$ and found that the angular momentum transfer rate rises roughly linearly with~$a$; it is negative for small~$a$ (meaning the angular momentum transfer to the hole) and reverses sign around $a\\approx 0.23$ (for $\\Psi_s=0.5\\Psi_{\\rm tot}$). The total energy transfer increases with~$a$ at an accelerated (i.e., faster than linear) rate, especially at larger values of~$a$; it is also negative at small~$a$, but becomes positive around $a=0.26$. This means that there is a narrow range $0.23..<a<0.26..$ where the integrated angular momentum flows from the hole to the disk, whereas the integrated red-shifted energy flows in the opposite direction.  Finally, we note that, in the case of open or partially-open field configuration responsible for the Blandford--Znajek process, one has to consider magnetic field lines that extend from the event horizon out to infinity. Since these field lines are not attached to a heavy infinitely conducting disk, their angular velocity $\\Omega_F(\\Psi)$ cannot be explicitly prescribed; it becomes just as undetermined as the poloidal current~$I(\\Psi)$ they carry. Fortunately, however, these field lines now have to cross two light cylinders (the inner one and the outer one). Since each of these is a singular surface of the Grad--Shafranov equation, one can impose corresponding regularity conditions on these two surfaces. Thus, we propose that one should be able to devise an iterative scheme that uses the two light-cylinder regularity conditions in a coordinated manner to determine the two free functions $\\Omega(\\Psi)$ and $I(\\Psi)$ simultaneously, as a part of the overall solution process. At the same time, the regularity conditions at the event horizon and at infinity could be used to obtain the asymptotic poloidal flux distributions at $r=r_H$ and at $r\\rightarrow\\infty$, respectively. We realise of course that iterating with respect to two functions simultaneously may be a very difficult task. This purely technical obstacle (in addition to having to deal with the separatrix between the open- and closed-field regions) is the primary reason why, in this paper, we have restricted ourselves to a configuration which has no open field lines extending from the black hole to infinity. We leave this problem as a topic for future research.  It is possible that, instead of solving the Grad--Shafranov equation itself, the easiest and most practical way to achieve a stationary solution will be to use a {\\it time-dependent} relativistic force-free code, such as one of those being developed now (Komissarov 2001, 2002a, 2004a; MacFadyen \\& Blandford 2003; Spitkovsky 2004; Krasnopolsky 2004, private communication).   I am very grateful to V.~Beskin, O.~Blaes, R.~ Blandford, S.~Boldyrev, A.~K{\\\"o}nigl, B.~C.~Low, M.~Lyutikov, A.~MacFadyen, V.~Pariev, B.~Punsly, and A.~Spitkovsky for many fruitful and stimulating discussions. I also would like to express my gratitude to the referee of this paper (Serguei Komissarov) for his very useful comments and suggestions that helped improve the paper. This research was supported by the National Science Foundation under Grant No.~PHY99-07949."
    },
    "0410/astro-ph0410523_arXiv.txt": {
        "abstract": "Mid-infrared photometry provides a robust technique for identifying active galaxies.  While the ultraviolet to mid-infrared ($\\lambda \\simlt 5 \\mu$m) continuum of normal galaxies is dominated by the composite stellar black body curve and peaks at approximately 1.6 $\\mu$m, the ultraviolet to mid-infrared continuum of active galaxies is dominated by a power law.  Consequently, with sufficient wavelength baseline, one can easily distinguish AGN from stellar populations. Mirroring the tendency of AGN to be {\\em bluer} than galaxies in the ultraviolet, where galaxies (and stars) sample the blue, rising portion of stellar spectra, AGN tend to be {\\em redder} than galaxies in the mid-infrared, where galaxies sample the red, falling portion of the stellar spectra.  We report on {\\it Spitzer Space Telescope} mid-infrared colors, derived from the IRAC Shallow Survey, of nearly 10,000 spectroscopically-identified sources from the AGN and Galaxy Evolution Survey.  Based on this spectroscopic sample, we find that simple mid-infrared color criteria provide remarkably robust separation of active galaxies from normal galaxies and Galactic stars, with over 80\\% completeness and less than 20\\% contamination.  Considering only broad-lined AGN, these mid-infrared color criteria identify over 90\\% of spectroscopically-identified quasars and Seyfert~1s.  Applying these color criteria to the full imaging data set, we discuss the implied surface density of AGN and find evidence for a large population of optically-obscured active galaxies. ",
        "introduction": "The dominant sources of energy production in the universe are fusion in stars and gravitational accretion onto supermassive black holes.  The tight correlation between nuclear black hole mass and bulge mass \\markcite{Magorrian:98, Tremaine:02}(\\eg Magorrian {et~al.} 1998; Tremaine {et~al.} 2002) implies the processes are intimately connected.  However, identifying an unbiased census of black holes in the universe remains challenging, hampering our ability to fully probe this connection.  For example, the spectral index of the X-ray background is significantly harder than the soft X-ray spectrum of bright, unobscured quasars, implying that there exists a large population of heavily-obscured AGN.  Identifying the sources responsible for the (hard) X-ray background was one of the primary motivators for the {\\it Chandra X-ray Observatory}, and {\\it Chandra} has successfully identified a population of high-redshift, heavily-obscured, luminous {\\em type II quasars} \\markcite{Stern:02a, Norman:02}(\\eg Stern {et~al.} 2002a; Norman {et~al.} 2002) as well as a significant population of apparently-normal galaxies with optical spectra dominated by starlight but hosting significant X-ray emission and implying the presence of an AGN \\markcite{Barger:01, Hornschemeier:01, Stern:02b}(\\eg Barger {et~al.} 2001; Hornschemeier {et~al.} 2001; Stern {et~al.} 2002b).  We show here that the {\\it Spitzer Space Telescope} \\markcite{Werner:04}(Werner {et~al.} 2004) also provides a valuable probe of AGN demographics.  The obscuring dust which hides AGN from ultraviolet, optical, and soft X-ray surveys should be a strong, largely isotropic emitter in the mid- to far-infrared ($\\lambda \\simgt 20 \\mu$m).  Furthermore, the different spectral energy distributions (SEDs) of stars and AGN produces different colors for normal and active galaxies in the mid-infrared ($\\lambda \\simlt 10 \\mu$m).  In particular, the composite black body spectra of the stellar population of normal galaxies produces an SED which peaks at approximately 1.6~$\\mu$m, while quasars have a power-law SED, $f_\\nu \\propto \\nu^{-\\alpha}$.  One of the early, traditional techniques to identify quasars relied on identifying sources with an ultraviolet excess, targeting sources whose SED does not plummet on the blue side of the stellar peak \\markcite{Schmidt:83}(\\eg Schmidt \\& Green 1983).  Though fruitful, this technique fails in two cases: (1) obscured quasars will preferentially suffer attenuation of their ultraviolet flux, and (2) high-redshift quasars disappear from the ultraviolet due to absorption from the \\lya\\ forest.  By contrast, we seek to exploit the many-fold recent improvement in mid-infrared sensitivity afforded by {\\it Spitzer} to target sources having SEDs that do not decline on the {\\em red} side of the stellar peak.  This provides an approach to identifying AGN that is relatively insensitive to extinction by dust or gas, and it has been demonstrated with a variety of instrumentation, including the Two Micron All-Sky Survey \\markcite{Cutri:01, Glikman:04}(Cutri {et~al.} 2001; Glikman {et~al.} 2004), the {\\it Infrared Space Observatory} \\markcite{Laurent:00, Haas:04}(Laurent {et~al.} 2000; Haas {et~al.} 2004), and the First Look Survey by {\\it Spitzer} \\markcite{Lacy:04}(Lacy {et~al.} 2004).  We present here a study of the mid-infrared colors of nearly 700 spectroscopically-confirmed quasars from the 9 deg$^2$ Bo\\\"otes field of the NOAO Deep Wide-Field Survey \\markcite{Jannuzi:99, Dey:05, Jannuzi:05}(NDWFS; Jannuzi \\& Dey 1999; Dey {et~al.} 2005; Jannuzi {et~al.} 2005).  The mid-infrared colors are from the IRAC Shallow Survey \\markcite{Eisenhardt:04}(Eisenhardt {et~al.} 2004), a guaranteed-time program with {\\it Spitzer}. The spectroscopy comes from the AGN and Galaxy Evolution Survey (AGES; Kochanek \\etal, in preparation).  AGES includes redshifts and spectral classifications for nearly ten thousand normal galaxies and nearly one thousand AGN.  Based on this spectroscopic sample, we show that mid-infrared photometry provides a robust technique for identifying AGN.  Using empirical color criteria to isolate AGN, we use the entire IRAC Shallow Survey to study the surface density of active galaxies, and evaluate the efficacy of mid-infrared selection of AGN.  A coordinated paper discusses the mid-infrared colors of X-ray sources in this field \\markcite{Gorjian:05}(Gorjian {et~al.} 2005).  Section~2 briefly summarizes the two surveys used.  Section~3 describes the mid-infrared properties of the spectroscopically-identified sources, followed by a discussion of mid-infrared selection of AGN in \\S4.  The results are summarized in \\S5.  Unless otherwise noted, all magnitudes refer to the Vega system. ",
        "conclusions": "We analyze mid-infrared photometry from the IRAC Shallow Survey of 4693 sources with well-determined spectroscopic redshifts from version 1.11 of AGES.  This sample includes 576 broad-lined AGN and 99 narrow-lined AGN.  We find that mid-infrared photometry provides remarkably robust separation of normal and active galaxies:  simple mid-infrared color criteria identify over 90\\% of the spectroscopically-identified broad-lined AGN and 40\\% of the spectroscopically-identified narrow-lined AGN.  Our selection criteria rely on a mid-infrared excess in AGN SEDs, mirroring the traditional ultraviolet excess technique which has long been exploited to identify quasars.  Only 17\\% of the AGES sources which fall within our mid-infrared color criteria for AGN are not spectroscopically classified as AGN.  Many of these sources likely indeed host a luminous, active nucleus, where either dust obscuration hides the signature of an AGN at observed optical wavelengths, or the AGN signature is drowned out by the stellar emission from the host galaxy \\markcite{Moran:02}(\\eg Moran, Filippenko, \\& Chornock 2002).  Simple mid-infrared selection criteria imply a surface density of 250 AGN candidates per deg$^2$ to an 8~$\\mu$m flux limit of 76~$\\mu$Jy, nearly triple the surface density inferred from optical surveys to this flux limit.  AGN missing from optical surveys may be of the obscured, {\\it type II} variety which are implied by unified models of AGN and are invoked by models of the X-ray background.  We show that the {\\it Spitzer Space Telescope} is a powerful tool for studying AGN demographics; {\\it Spitzer} will help identify a less-biased sample of super massive black holes in the universe, allowing us to probe the interconnection of fusion-driven and accretion-driven energy production over cosmic history."
    },
    "0410/hep-th0410270_arXiv.txt": {
        "abstract": "We suggest that spontaneous eternal inflation can provide a natural explanation for the thermodynamic arrow of time, and discuss the underlying assumptions and consequences of this view.  In the absence of inflation, we argue that systems coupled to gravity usually evolve asymptotically to the vacuum, which is the only natural state in a thermodynamic sense.  In the presence of a small positive vacuum energy and an appropriate inflaton field, the de~Sitter vacuum is unstable to the spontaneous onset of inflation at a higher energy scale.  Starting from de~Sitter, inflation can increase the total entropy of the universe without bound, creating universes similar to ours in the process.  An important consequence of this picture is that inflation occurs asymptotically both forwards and backwards in time, implying a universe that is (statistically) time-symmetric on ultra-large scales. ",
        "introduction": "The role of initial conditions in cosmology is unique within the physical sciences.  We only have a single observable universe, rather than the ability to change boundary conditions and run experiments multiple times.  A complete theory of cosmology therefore involves not only a set of dynamical laws, but a specification of the particular initial conditions giving rise to the universe we see.  One could certainly argue that the origin of our initial conditions is not an answerable scientific question.  Given the state of our universe at the present time, and a complete set of dynamical laws describing its evolution, we can in principle solve for the entire history, including whatever the initial state was.  Ultimately, we are stuck with the boundary conditions we have.  Similarly, however, we are stuck with the laws of physics that we have, but this constraint doesn't stop us from searching for deep principles underlying their nature. It therefore seems sensible to treat our initial conditions in the same way, and try to understand why we have these conditions rather than some others.  In particular, we would like to know whether the initial conditions for our universe are in some sense ``natural.''  The notion of naturalness in this context is not precisely defined, but in simple cases we know it when we see it.  Our current universe exhibits a pronounced matter/antimatter asymmetry, which seems unnatural, so we search for mechanisms of baryogenesis to create this asymmetry dynamically.  The large-scale homogeneity and isotropy of our universe is unstable, so we search for mechanisms in the early universe to generate this smoothness dynamically.  One aspect of the boundary conditions for our universe that remains puzzling, or at least controversial, is the arrow of time (for discussions see \\cite{price,albrecht03,Coule:2002zb}).  The Second Law of Thermodynamics tells us that the entropy of a closed system is non-decreasing as a function of time.  A comoving patch corresponding to our observable universe is not precisely a closed system, but is pretty close, since homogeneity of the universe near the boundary implies there should be no net entropy flux into or out of our patch.  And indeed, we observe increasing entropy in our universe; the thermodynamic arrow of time is the direction picked out by this increase of entropy.  Long ago Boltzmann understood the increase of entropy as a statistical phenomenon, as evolution from a very small region of phase space to a larger region.  The origin of the Second Law is traced back to initial conditions:  the early universe had an extremely low entropy, allowing it to continue to increase thereafter.  Even after 13.7 billion years of entropy growth, it remains low in our current universe.  \tWithin our observable patch, the entropy in ordinary matter is of order \\be S_M(U) \\sim 10^{88}\\ . \\label{meu} \\ee This figure is dominated by photons and neutrinos; dark matter may also contribute, but wouldn't significantly change the final answer.  However, Penrose has emphasized \\cite{penrose} that this entropy could be much larger. Our universe remains quite smooth on large scales, which is a finely-tuned condition to be in.  To see this, we can simply compare the current entropy to what it could have been, for example if more energy were stored in black holes. The Bekenstein-Hawking entropy of a black hole \\cite{Bekenstein:1975tw, Hawking:1974sw} is proportional to its horizon area, \\be S_{BH} = {A\\over 4G}\\ . \\ee We believe that galaxies similar to ours contain supermassive black holes at their centers, of order a million solar masses; the entropy of a single such black hole is \\be S_{BH} \\sim 10^{89}\\left({M_{BH}\\over 10^6 M_\\odot}\\right)^2\\ . \\ee Any one such black hole therefore contains more entropy than all of the ordinary matter in the visible universe.  But there are perhaps $10^{10}$ such black holes in the universe; their total entropy is thus \\be S_{BH}(U) \\sim 10^{99}\\ , \\label{bheu} \\ee which represents most of the entropy within our observable universe. This seems like a large number, but should be compared to how large it could be.  The total amount of mass in the observable universe is \\be M_U \\sim \\rho H_0^{-3}\\sim G^{-1}H_0^{-1} \\sim 10^{22} M_\\odot\\ . \\ee If all of this mass were collected into a single black hole, the entropy would be \\be S_{\\rm max}(U) \\sim 10^{121}\\ , \\label{eumax} \\ee larger than the actual entropy by a factor of $10^{22}$.  Moreover, in the early universe there weren't any black holes, and the entropy was dominated by the matter entropy (\\ref{meu}), smaller than the equivalent black-hole entropy by a factor of $10^{33}$.  So the entropy of our universe seems to be very small, and seems to have evolved from an era where it was significantly smaller.  The initial conditions for our observable universe thus have extremely low entropy.  We would like to understand these conditions as arising through dynamical evolution from a natural starting point; it is not immediately obvious how this could be achieved, however, since dynamical evolution tends to increase the entropy.  The low entropy of our current universe is clearly related to its homogeneity and isotropy, features which are supposed to be explained by the inflationary universe scenario \\cite{Guth:1980zm,Linde:1981mu, Albrecht:1982wi,Starobinsky:1979ty,Starobinsky:1980te}. It is therefore tempting to invoke inflation as the origin of the thermodynamic arrow of time, and indeed this move has been made \\cite{Davies:1983nf}. The converse argument, however, has also been advanced: that inflation never works to decrease the entropy, and posits that our observable universe originates in a small patch whose entropy is fantastically less than it could have been, so that in fact there is a hidden fine-tuning of initial conditions implicit in the inflationary scenario, which consequently (in the extreme version of this reasoning) doesn't explain anything at all \\cite{penrose,Page:1983uh,unruh, Hollands:2002yb}.  In this paper we suggest a resolution of this apparent tension between inflation and the Second Law.  We first discuss in general terms what sort of evolution for the universe would qualify as naturally giving rise to an arrow of time through dynamical processes.  We then review the arguments concerning the role of inflation in setting up the early universe in a low-entropy state, concluding that ordinary inflation does not directly explain the origin of the arrow of time if one believes in unitary evolution. We then argue that, in the absence of a vacuum energy or inflation of any sort, systems coupled to gravity generically evolve to flat, empty space.  In the presence of a small positive vacuum energy (such as seems to exist in our universe \\cite{Riess:1998cb,Perlmutter:1998np,Carroll:2000fy}), systems will tend to initially empty out to an approximate de~Sitter vacuum. If there is also at least one scalar inflaton field with an appropriate potential, however, there is an instability: quantum fluctuations in this field will eventually create an inflating patch, and according to the traditional arguments for eternal inflation \\cite{Vilenkin:1983xq,Linde:1986fc, Linde:1986fd,Goncharov:1987ir,Guth:2000ka} the total physical volume of inflating space will continue to increase thereafter, while ``pocket universes'' drop out of the inflating background and reheat. The idea of inflation beginning spontaneously has been suggested previously \\cite{Vilenkin:1982de,Vilenkin:1983xq,Linde:1983mx,Vilenkin:1984wp, Starobinsky:1986fx,Goncharov:1986ua, Farhi:1986ty,Vilenkin:1987kf,Farhi:1989yr,Fischler:1989se,Fischler:1990pk, Linde:1991sk}.  Recently, Garriga and Vilenkin \\cite{Garriga:1997ef} have suggested that spontaneous inflation in a low-energy de~Sitter background could lead to a recycling universe, while Dyson, Kleban and Susskind and Albrecht and Sorbo have examined spontaneous inflation in the context of a causal-patch view inspired by the holographic principle \\cite{Dyson:2002pf, Albrecht:2004ke}; we discuss the relationship between our picture and these ideas in greater detail below.  The crucial features of this scenario are thus the following: \\begin{enumerate} \\item Generic initial states empty out and approach de~Sitter space. \\item In the presence of an inflaton field, the de~Sitter vacuum will be unstable to the onset of eternal inflation. Once eternal inflation starts, the entropy of the universe can grow without bound, never reaching thermal equilibrium. \\item In the process of increasing the entropy, eternal inflation creates large regions of space similar to our observable universe. \\end{enumerate} As an important consequence of our reasoning, we obtain a picture of the ultra-large-scale structure of the universe. As in any model of eternal inflation, the ultimate state of the universe is a fractal distribution of inflating and post-inflating regions \\cite{Linde:1986fd,Goncharov:1987ir,Aryal:1987vn,Winitzki:2001np}, in which our local Big Bang is not a particularly unique event. We also predict that this structure should be recovered infinitely far into the past, with a {\\it reversed} thermodynamic arrow of time.  Our overall universe is therefore statistically time-symmetric about some Cauchy surface of minimum entropy.  The details of the state on this surface are not important; starting from a generic state, inflation eventually occurs towards both the past and future, erasing any information about the ``initial'' conditions. ",
        "conclusions": "\\label{discussion}  In this paper we have addressed the question of why the entropy of our observable universe appears to have been extremely small in the past.  It has long been suspected that inflation might have something to do with the answer, although the fantastically low entropy of the proto-inflationary universe was a significant obstacle to constructing a convincing picture.  The answer to this conundrum must lie in the process by which inflation begins; we have proposed a scenario in which this happens naturally from the evolution of some arbitrarily-chosen state.  It is interesting to note that that the recently discovered cosmological constant plays a crucial role in our scenario for generating the arrow of time.  The basic ingredients of our picture are as follows.  We consider an arbitrary state of the universe, specified on some Cauchy surface.  We argued in Section~\\ref{typicalstates} that the generic evolution of a system coupled to gravity is to dilute excitations via the expansion of spacetime.  Black holes and other inhomogeneities may form, but they will eventually decay away, so that the universe approaches empty space.  However, in the presence of a positive vacuum energy and an appropriate inflaton field, the resulting de~Sitter phase is unstable to the spontaneous onset of inflation, instigated by the thermal fluctuations of the inflaton. If the inflaton fluctuates sufficiently high that eternal inflation can begin, it will continue forever, and new pocket universes will be brought into being in those places where the field rolls down the potential and reheats.  This chain of events happens both to the past and the future of the specified Cauchy surface, leading to a statistically time-symmetric universe as portrayed in Figure~\\ref{ultralarge}.  An arrow of time is dynamically generated in both the past and the future, as the universe continually acts to increase its entropy.  In addition to the desire to understand the origin of the arrow of time, a primary motivation for our study has been to understand the onset of inflation.  As discussed in Section~\\ref{critique}, the unitarity critique argues that a proto-inflationary patch of spacetime is much lower entropy than that of an ordinary radiation-dominated universe, and hence is less likely to arise as a random fluctuation.  In our picture, the answer to this conundrum lies in the fact that the beginning of our observable Big Bang cosmology does not arise as a random choice in a large phase space of initial conditions, but rather comes via a quantum fluctuation from a very specific prior state -- an empty de~Sitter universe that is the natural consequence of evolution from generic initial conditions.  By taking seriously the ability of spacetime to expand and dilute degrees of freedom, we claim to have shown how an arrow of time can naturally arise dynamically in the course of the evolution from a generic boundary condition.  In the classification introduced in Section~\\ref{arrow}, our proposal imagines that there do not exist any maximum-entropy equilibrium states, but rather that the entropy can increase from any starting configuration.  This is not, of course, sufficient; it is also necessary to imagine that the path to increasing the entropy naturally creates regions of spacetime resembling our observable universe.   In the presence of a nonzero vacuum energy and an appropriate inflaton field, we suggest that thermal fluctuations from de~Sitter space into eternal inflation provide precisely the correct mechanism.  A number of other cosmological scenarios have been proposed in which the Big Bang is not a boundary to spacetime, but simply a phase through which the universe passes.  These include the pre-Big-Bang scenario \\cite{Gasperini:2002bn,Lidsey:1999mc}, the ekpyrotic and cyclic universe scenarios \\cite{Khoury:2001zk,Khoury:2001bz,Steinhardt:2001st}, the Aguirre-Gratton scenario of eternal inflation \\cite{Aguirre:2003ck}, and Bojowald's loop-quantum-gravity cosmology \\cite{Bojowald:2003cc,Bojowald:2004kt}. To the best of our understanding, each of these proposals invokes special low-entropy conditions on some Cauchy surface, either asymptotically in the far past or at some moment of minimum size for the universe.  In our picture, on the other hand, there is a slice of spacetime on which the entropy is minimized, but that entropy can be arbitrarily large. The Big Bang in our past is not a unique moment in the history of the universe; it is simply one of the many times that inflation spontaneously began from a background de~Sitter phase, similar to the proposal of Garriga and Vilenkin \\cite{Garriga:1997ef}. Along with the fractal distribution of pocket universes to the far past and far future, this feature is another reminder of the importance of overcoming the Robertson-Walker intuition we naturally develop by thinking about the patch of the universe we are actually able to observe.  There are a number of points in our scenario that have yet to be perfectly understood.  One obvious point is how the ``initial'' state is chosen.  We have argued that the ultimate evolution to the past and future is essentially insensitive to the details of this state, but it is nevertheless interesting to ask what it may have been. On a more technical level, the idea of spontaneous onset of inflation (and related ideas within the paradigm of eternal inflation) deserves closer investigation.  We have calculated the renormalized fluctuations of a scalar field in a fixed de~Sitter background, and imagined that the back-reaction of the metric would lead to inflation if the fluctuation were sufficiently large.  A more rigorous calculation would have to involve quantum gravity from the start, at least at a semiclassical level. Finally, the issues concerning the number and evolution of degrees of freedom clearly warrant further research.  Our assumption has been that the number of degrees of freedom is infinite but fixed, not increasing during inflation.  We therefore require a literally infinite number of degrees of freedom to be put in their ground states during the evolution toward de~Sitter, before the onset of inflation.  It is important that we understand this issue more fully, although doing so may require a deep knowledge of quantum gravity.  An interesting implication of our scenario is the non-privileged role played by the Big Bang of our observable universe.  As in other models of eternal inflation, the future history of spacetime takes on a fractal structure of ever-increasing volume; in addition we suggest that a similar structure is found in a time-reversed sense in the far past.  The boundary conditions of the universe on ultra-large scales are unlikely to resemble simple Robertson-Walker geometries in any way. Indeed, the cacophony of matter and radiation in our observable patch of universe appears in this picture as a mere byproduct of the relentless evolution of the larger spacetime toward states of higher entropy.  Stars and galaxies are seen as exaptations -- structures that have found a use other than that for which they were originally developed by evolution."
    },
    "0410/cs0410071_arXiv.txt": {
        "abstract": "We present results from the first geological field tests of the `Cyborg Astrobiologist', which is a wearable computer and video camcorder system that we are using to test and train a computer-vision system towards having some of the autonomous decision-making capabilities of a field-geologist and field-astrobiologist. The Cyborg Astrobiologist platform has thus far been used for testing and development of these algorithms and systems: robotic acquisition of quasi-mosaics of images, real-time image segmentation, and real-time determination of interesting points in the image mosaics.  The hardware and software systems function reliably, and the computer-vision algorithms are adequate for the first field tests. In addition to the proof-of-concept aspect of these field tests, the main result of these field tests is the enumeration of those issues that we can improve in the future, including: first, detection and accounting for shadows caused by 3D jagged edges in the outcrop; second, reincorporation of more sophisticated texture-analysis algorithms into the system; third, creation of hardware and software capabilities to control the camera's zoom lens in an intelligent manner; and fourth, development of algorithms for interpretation of complex geological scenery. Nonetheless, despite these technical inadequacies, this Cyborg Astrobiologist system, consisting of a camera-equipped wearable-computer and its computer-vision algorithms, has demonstrated its ability of finding genuinely interesting points in real-time in the geological scenery, and then gathering more information about these interest points in an automated manner. ",
        "introduction": "\\begin{spacing}{1} {\\it ``With great autonomy comes great responsibility''}  --with apologies to Peter Parker's Uncle Ben (Lee {\\it et al.}\\ 2002, 2004). \\end{spacing}  \\vspace{0.5cm} Outside of the Mars robotics community, it is commonly presumed that the robotic rovers on Mars are controlled in a time-delayed joystick manner, wherein commands are sent to the rovers several if not many times per day, as new information is acquired from the rovers' sensors.  However, inside the Mars robotics community, they have learned that this process is rather cumbersome, and they have developed much more elegant methods for robotic control of the rovers on Mars, with highly significant degrees of robotic autonomy. Particularly, the Mars Exploration Rover (MER) team has demonstrated autonomy for the two robotic rovers Spirit \\& Opportunity to the level that: practically all commands for a given Martian day (1 `sol' $=$ 24.6 hours) are delivered to each rover from Earth before the robot wakens from its power-conserving nighttime resting mode (Arvidson {\\it et al.}\\ 2003; Maki {\\it et al.}\\ 2003; Bell {\\it et al.}\\ 2003; Squyres {\\it et al.}\\ 2003; Crisp {\\it et al.}\\ 2003). Each rover then follows the commanded sequence of moves for the entire sol, moving to desired locations, articulating its arm with its sensors to desired points in the workspace of the robot, and acquiring data from the cameras and chemical sensors. From an outsider's point of view, these capabilities may not seem to be significantly autonomous, in that all the commands are being sent from Earth, and the MER rovers are merely executing those commands. But upon closer inspection, what the MER team has achieved is truly amazing: \\begin{spacing}{1} \\begin{trivlist} \\raggedright \\item $\\bullet$ Firstly, the rovers can move to points $50-150$ meters away in one sol with autonomous obstacle avoidance enabled for the uncertain or dangerous parts of the journey (Goldberg, Maimone \\& Matthies 2002; Nesnas, Maimone \\&  Das 1999). \\item $\\bullet$ Secondly, prior to a given sol, based upon information received after the previous sol, the MER team has the remarkable capabilities to develop a command sequence of tens or hundreds of robotic commands for the entire sol. As of July 4, 2004, this was taking 4-5 hours per sol for the mission team to complete, rather  than the 17 hours per sol that it took at the beginning of the MER missions (Squyres, unpublished\\footnote{lecture notes from the course on the MER rovers from Astrobiology Summer School {\\it ``Planet Mars''}, Universidad Internacional Menendez Pelayo Santander, Spain}, 2004). \\end{trivlist} \\end{spacing}     Such capabilities for semi-autonomous teleoperated robotic `movement and discovery' are a significant leap beyond the capabilities of the previous Mars lander missions of Viking I \\& II and of Pathfinder \\& Sojourner, particularly in the large improvements of the mobility of both the cameras and the instrumentation onboard the robot, allowing significant discoveries to be made, and new insights to be obtained (Chan {\\it et al.}\\ 2004; Squyres {\\it et al.}\\ 2004; Squyres \\& Athena Science Team 2004; Moore 2004; Catling 2004; Crumpler {\\it et al.}\\ 2004; Golembek {\\it et al.}\\ 2004; Squyres 2004; McSween {\\it et al.}\\ 2004; Christensen {\\it et al.}\\ 2004; Arvidson {\\it et al.}\\ 2004; Greeley {\\it et al.}\\ 2004; Herkenhoff {\\it et al.}\\ 2004; Bell {\\it et al.}\\ 2004).  Nonetheless, we would like to build upon this great success of the MER rovers by developing enhancing technology that could be deployed in future robotic and/or human exploration missions to the Moon, Mars, and Europa.  \\markright{The Cyborg Astrobiologist: First Field Experience}   One future mission deserves special discussion for the technology developments described in this paper: the Mars Science Laboratory, planned for launch in 2009 (MSL'2009). A particular capability desired for this MSL'2009 mission will be to rapidly traverse to up to three geologically-different scientific points-of-interest within the landing ellipse (Heninger {\\it et al.}\\footnote{Heninger, B., M. Sander, J. Simmonds, F. Palluconi, B. Muirhead, C. Whetsel, ``Mars Science Laboratory Mission 2009 Landed Science Payload'', Proposal Information Package (draft version),   http://centauri.larc.nasa.gov/msl/PIP-Drft\\_FBO-RevA-031113.pdf}, unpublished, 2003).  These three geologically-different sites will be chosen from Earth by analysis of relevant remote-sensing imagery from a possible co-launched, satellite-based, Synthetic Aperture Radar (SAR) system or high-resolution maps from prior missions (i.e.\\ 30 cm/pixel spatial resolution from the HiRISE imager on the Mars Reconnaisance Explorer with imagery expected in 2006-2007, McEwen {\\it et al.}\\ 2003). Possible desired maximal traversal rates could range from 300-2000 meters/sol in order to reach each of the three points-of-interest in the landing ellipse in minimum time.  These are our rough estimates of the typical maximum traversal rates, based upon current capabilities, mission goals, increased rover size and expectations of the state of technology for the possible 2009 mission;  our estimate is partially based upon such work as described in: Goldberg, Maimone \\& Matthies (2002); Olson, Matthies, Schoppers \\& Maimone (2003); Crawford (unpublished\\footnote{``NASA Applications of Autonomy Technology'', NASA Ames Research Center, Intelligent Data Understanding Seminar,  http://is.arc.nasa.gov/IDU/slides/Crawford\\_Aut02c.pdf}, 2002); and Crawford \\& Tamppari (unpublished\\footnote{``Mars Science Laboratory -- Autonomy Requirements Analysis'', NASA Ames Research Center, Intelligent Data Understanding Seminar,  http://is.arc.nasa.gov/IDU/slides/Crawford\\_MSL02c.pdf}, 2002).     Given these substantial expected traversal rates of the MSL'2009 rover, autonomous obstacle avoidance (Goldberg, Maimone \\& Matthies 2002) and autonomous visual odometry \\& localization (Olson, Matthies, Schoppers \\& Maimone 2003) will be essential to achieve these rates, since otherwise, rover damage and slow science-target approach would be the results.  Given such autonomy in the rapid traverses, it behooves us to enable the autonomous rover with sufficient scientific responsibility. Otherwise, the robotic rover exploration system might drive right past an important scientific target-of-opportunity along the way to the human-chosen scientific point-of-interest. Crawford \\& Tamppari (previous footnote, 2002) and their team summarize possible `autonomous traverse science', in which every 20-30 meters during a 300 meter traverse (in their example), science pancam and Mini-TES (Thermal Emission Spectrometer) image mosaics are autonomously obtained. They state that ``there {\\it may be} onboard analysis of the science data from the pancam and the mini-TES, which compares this data to predefined signatures of carbonates or other targets of interest. If detected, traverse may be halted and information relayed back to Earth.'' This onboard analysis of the science data is precisely the technology issue that we have been working towards solving. This paper is the first detailed published report describing our progress towards giving a robot some aspects of autonomous recognition of scientific targets-of-opportunity (see McGuire {\\it et al.}\\ (2004) for a brief description and motivation of our work).   \\begin{figure}[ht] \\center{\\includegraphics[height=5.0cm]{IJAmcguire1a.ps} \\includegraphics[height=5.0cm]{IJAmcguire1b.ps}} \\caption{ On the left, D\\'iaz Mart\\'inez \\& McGuire with the Cyborg Astrobiologist System at 1:30PM on 3 March 2004, 10 meters from the outcrop cliff that is being studied during the first geological field mission near Rivas Vaciamadrid. We are taking notes prior to acquiring one of our last-of-the-day mosaics and its set of interest-point chips. On the right is a close-up view of the same scene, but focusing on the half-seated/half-tripod Cyborg Astrobiologist and the assisting human Geologist. Photo copyright: D\\'iaz Mart\\'inez, Orm\\\"o \\& McGuire} \\end{figure}  Our general aim of study for this project is to develop autonomous computer vision algorithms for natural terrestrial geological scenery in uncontrolled lighting conditions (see our `Cyborg Astrobiologist' research platform in Figure 1), with the expectation that such algorithms will mature in the future so that they can be used in the human or robotic astrobiological exploration of Mars and some of the moons of our solar system.  This project is not only aimed at giving more scientific autonomy to robotic explorers of the surfaces of other planetary bodies, but it is also meant to take some of the burden off of the large team of human geologists at mission control here on the Earth by performing some of the low-level geological image analysis autonomously, thus freeing the human geologists to do more high-level analysis of the scenery. Such a system could also assist human geologists here in the field on the Earth for their studies, as well as astronauts in their geological field studies (who may be less-practiced in geology and more burdened with other duties and concerns). Our long-term aim of study is for the Cyborg Astrobiologist system to be able to understand the geological history of an outcrop in order to decide which area should be approached for sampling. Our short-term aim of study is to demonstrate that the Cyborg Astrobiologist system can (in the field) find the most unique regions in an image mosaic, and then to autonomously study those interesting regions in more detail.  In this report, we give additional introductory background in Section 1.1, a discussion of the hardware and software for our ``Cyborg Astrobiologist'' system in Section 2, summaries \\& results of the first two geological expeditions of our Cyborg Astrobiologist system in Sections 3 \\& 4, followed by a more general discussion \\& conclusions in Section 5.   \\subsection{Autonomous Recognition of Scientific Astrobiological Targets-of-Opportunity} We discuss here two of the related efforts in the development of autonomous recognition of scientific targets-of-opportunity for astrobiological exploration: firstly, the work on developing a Nomad robot to search for meteorites in Antartica led by the Carnegie Mellon University Robotics Institute (CMU-RI) (Apostolopoulos {\\it et al.}\\ 2000; Pederson 2001; Wagner 2000; Vandapel {\\it et al.}\\ 2000), and secondly, the work by a group at NASA Ames Research Center (NASA-ARC) on developing a Geological Field Assistant (GFA) (Gulick {\\it et al.}\\ (2001, 2002a, 2002b, 2003, 2004)).  There are several other efforts in the community which also deserve mention (including Volpe 2003; Huntsberger {\\it et al.}\\ 2002; Storrie-Lombardi, Grigolini {\\it et al.}\\ 2002; Corsetti \\& Storrie-Lombardi 2003; Cabrol {\\it et al.}\\ 2001; Wettergreen {\\it et al.}\\ 1999; Whittaker {\\it et al.}\\ 1997; Cheeseman {\\it et al.}\\ 1988; Cheeseman \\& Stutz\\ 1996). Outside of the astrobiology community, there is some similar work on automated obstacle avoidance (instead of target discovery) by Batavia \\& Singh (2001), among others.  The CMU-RI team has programmed the capabilities to autonomously distinguish between rocks and ice/snow in various lighting conditions (direct sunlight, shadows, and diffuse light in overcast conditions) into the science computer of its meteorite-searching Nomad robot (Apostolopoulos {\\it et al.}\\ 2000). This autonomous decision between whether a pixel is ice/snow or if it is rock is decided by how `blue' the pixel is: if the pixel has its `blue' ratio ($blue/(blue + green)$) above a certain lighting-dependent threshold, then it is considered to be ice/snow; otherwise, it is considered to be rock/meteorite. This system works rather well except when the rocks are in shadows (but we naively suggest that a non-linear separator in the 2D space of `blue' and intensity could solve this problem even in shadows). After a rock/meteorite is segmented from the ice/snow based upon how `blue' it is, it can then trigger the Nomad robot's Bayesian classifier network system to compute the probability of the rock being any one of six different meteor types or of 19 different terrestrial rock types, based upon data from a metal detector, from a color camera detector, and from a spectrometer. The Nomad robot ``[autonomously] found and [autonomously] classified five indigenous meteorites [of 42 rocky targets studied and classified over a 2500 m$^2$ search area] during an expedition to a remote site of Elephant Moraine [in Antarctica] in January 2000.'' (Apostolopoulos {\\it et al.}\\ 2000).  The NASA-ARC GFA team have developed a number of algorithms for assisting an astronaut in doing geology or in adding capabilities to an autonomous robotic geologist (Gulick {\\it et al.}\\ (2001, 2002a, 2002b, 2003, 2004)). They have been building an image and spectral database of rocks and minerals [now there are over 700 such rock and mineral entries in their database] for continued development of advanced classifier algorithms, including future development work to autonomously identify `key igneous, sedimentary and some metamorphic rocks' (Gulick {\\it et al.}\\ 2003). As of 2003, they were able to identify ``some igneous rocks from texture and color information; quartz, silica polymorphs, calcite, pyroxene and jarosite from both visible/near-infrared and mid-infrared spectra; and high-iron pyroxenes and iron-bearing minerals using visible/near-infrared spectra only'' (Gulick {\\it et al.}\\ 2003). As of 2004, the NASA Jet Propulsion Laboratory (JPL) CLARAty system team (Volpe 2004) had incorporated the NASA-ARC GFA team's rock, layer and horizon detectors (Gulick {\\it et al.}\\ 2001) for ``possible use by MSL'2009 and [other] future Mars robotics missions'' (Gulick {\\it et al.}\\ 2004). Work reported in 2004 includes significant progress on their igneous rock detection system: the Decision Tree method has yielded better and more decipherable results than the Bayesian Network approach, giving ``correct identification of greater than $80\\%$ of granites and granodiorites and greater than $80\\%$ of andesites and basalts, using color and texture algorithms combined [without spectral information]'' (Gulick {\\it et al.}\\ 2004). ",
        "conclusions": "Both the geologists on our team concur with the judgement of the Cyborg Astrobiologist software system, that the two dark \\& wet spots on the cliff wall were the most interesting spots during the second mission. However, the two geologists also state that this largely depends on the aims of study for the geological field trip; if the aim of the study is to search for hydrological features, then these two dark \\& wet spots are certainly interesting; however if faulting between layers is of interest, then the dark \\& wet spots would not retain as much interest. One question which we have thus far left unstudied is ``What would the Cyborg Astrobiologist system have found interesting during the second mission if the two dark \\& wet spots had not been present during the second mission?'' It is possible that it would again have found some dark shadow particularly interesting, but with the improvements made to the system between the first and second mission, it is also possible that it could have found a different feature of the cliff wall more interesting.  \\subsection{Thought Experiment: Blueberry Fields in the Meridiani Planum} \\markright{The Cyborg Astrobiologist: Discussion: Blueberry Fields at Meridiani} One of the remarkable discoveries made by the MER team this year is the existence of small, dark \\& round pebbles (a.k.a. `blueberries') in many places near the landing site of the Opportunity rover Meridiani Planum. The Opportunity rover was sent to Meridiani Planum because this was the main site on Mars that had been observed by the Thermal Emission Spectrometer on the Mars Global Surveyor orbiter to have a large quantity of gray hematite (Christensen {\\it et al.}\\ 2000). Gray hematite is an iron-bearing mineral which often forms in the presence of water, particularly as a precipitate in cold standing water (Glotch {\\it et al.}\\ 2004). The so-called blueberries that Opportunity discovered at Meridiani Planum were also observed to have significant hematite in them, based upon studies from both of the mini-TES and M\\\"ossbauer spectrometer instruments onboard the Opportunity rover (Squyres, unpublished\\footnote{lecture notes from the course on the MER rovers from Astrobiology Summer School {\\it ``Planet Mars''}, Universidad Internacional Menendez Pelayo Santander, Spain}, 2004). There is significant excitement in the planetary geology community about this discovery, partly due to the existence of a place on Earth (Utah) where similar blueberries have been formed (Chan {\\it et al.}\\ 2004; Squyres \\& Athena Science Team, 2004; Moore 2004; Catling 2004).  Such a geological phenomenon suggests a thought experiment that could be applied to our Cyborg Astrobiologist system as it currently exists: ``If the Cyborg Astrobiologist system was taken to a field covered with Meridiani/Utah blueberries, would it find the blueberries to be interesting?'' The brief, unqualified answer to this thought experiment is: ``No''. If the blueberries covered the landing site, so that their coverage exceeded some percentage (assume for argument's sake that the areal coverage of blueberries at the location of the Cyborg Astrobiologist is 80\\%), then an image mosaic taken by the Cyborg Astrobiologist would be dominated by dark blueberries. Therefore, the Cyborg Astrobiologist's uncommon map algorithm, as it now exists, would likely search the remaining 20\\% of the area for the smallest color class of pixels, and it would ignore the blueberries.  However, if the Cyborg Astrobiologist was allowed to approach the field of blueberries from a distance (i.e., from orbit or from a neighboring area without blueberries), then it is probable that the Cyborg Astrobiologist would at some point find the blueberries to be new and interesting. This is because blueberries don't cover all of Mars, and so at some spatial scale, they only have a small areal coverage. So the qualified answer to the thought experiment is: ``Yes''.  Furthermore, we could program more advanced renditions of the software inside of the Cyborg Astrobiologist so that it could find the blueberries interesting even if the Cyborg Astrobiologist was not allowed to approach the blueberry field from a distance. Such algorithm enhancements could include the texture \\& edge information suggested above, as well as a possible algorithm to search for uncommon shapes. With spectrometer information (i.e., visible/near-infrared, mid-infrared, Raman, and/or M\\\"ossbauer) and with biased searching for particularly-interesting spectra, textures, rock shapes and/or rock juxtapositions, the odds of success could also be maximized, either in hunting for blueberries or in serendipitously finding such blueberries.  These advances in our system could take some years to implement, but we will surely implement some of them in the near-term. In the meantime, as discussed above, the valued concept of a geological approach will remain our solution to the `blueberry-field' thought experiment. Nevertheless, such a thought experiment does point out a significant inadequacy of our simple computer vision algorithm which searches for small regions of the color imagery, as discussed in the preceding three paragraphs. Such simple algorithms need to be enhanced and brought into relation with other algorithms, and perhaps with a visual memory of past imagery (i.e., the VAMPIRE project (Wrede {\\it et al.}\\ 2004)).  A visual memory would also be useful to record and track changes in the areal coverage of components in the scenery, as mosaics or images are acquired from different camera positions or for different outcrops. In a different outcrop, the dark wet spots might dominate the whole outcrop, or the whitish color from the gypsum might not dominate the outcrop. The visual memory could keep track of outcrops studied in the recent past, and inform the user when an uncommon area of a particular image or mosaic is truly novel in areal coverage, when compared with an area of similar color or texture characteristics in prior imagery.  \\subsection{Outlook} \\markright{The Cyborg Astrobiologist: Outlook}  The NEO programming for this Cyborg Astrobiologist/Geologist project was initiated with the SONY Handycam in April 2002. The wearable computer arrived in June 2003, and the head mount\\-ed display arrived in November 2003. We now have a reliably functioning human and hardware and software Cyborg Geologist system, which is partly robotic with its Pan Tilt camera mount. This robotic extension allows the camera to be pointed repeatedly, precisely \\& automatically in different directions.  Prior to each of the two field-trips to Rivas Vaciamadrid, the Cyborg Geologist system was tested on the grounds of our Spanish research institute (CAB). These tests were mostly on a  wide prairie field: by winter, this prairie field was plowed and full of 3-10 cm diameter stones of different, dull colors; and by summer, this prairie field was full of prairie grasses and differently-colored wildflowers. In these nearby field-trips, the Cyborg Geologist has been providing us with constructive criticism of the basic capabilities of our image-segment\\-ation and interest-map algorithms.  As of our initial field-trips to Rivas, the Cyborg Geologist system has accompanied two Human field geologists to a site of geological interest (i.e. layered rock outcrops) for more advanced testing of these computer-vision algorithms.  We plan several more such trips in the coming months.    Based upon the performance of the Cyborg Astrobiologist system during the 1st mission to Rivas in March 2004 on the outcropping cliffs near Rivas Vaciamadrid, we have decided that the system was paying too much attention to the shadows made by the 3D structure of the cliffs. We hope to improve the Cyborg Astrobiologist system in the next months in order to detect and to pay less attention to shadows. We also hope to upgrade our system to include: image-segmentation based upon micro-texture; and adaptive methods for summing the uncommon maps in order to compute the interest map.    Based upon the significantly-improved performance of the Cyborg Astrobiologist system during the 2nd mission to Rivas in June 2004, we conclude that the system now is debugged sufficiently so as to be able to produce studies of the utility of particular computer vision algorithms for geological deployment in the field. We have outlined some possibilities for improvement of the system based upon the second field trip, particularly in the improvement in the systems-level algorithms needed in order to more intelligently drive the approach of the Cyborg or robotic system towards a complex geological outcrop. These possible systems-level improvements include: a better interest-map algorithm, with adaptation and more layers; hardware \\& software for intelligent use of the camera's zoom lens; a memory of the image segmentation performed at greater distance or lower magnification of the zoom lens; and high-level image-interpretation capabilities.    Now that we have demonstrated that this software and hardware in the Cyborg Astrobiologist system can function for developing and testing computer-vision algorithms for robotic exploration of a geological site, we have some decisions to make as to future directions for this project, options for these future directions include: \\begin{spacing}{1} \\begin{trivlist} \\raggedright \\item $\\bullet$ Performing further offline analysis and algorithm-development for the imagery obtained at Rivas Vaciamadrid: several of the parameters of the algorithms need testing for their optimality, and further enhancements of the algorithms could be made. \\item $\\bullet$ Optimizing the image-processing \\& robotic-control code for the current Cyborg Astrobiologist system for speed and memory utilization. \\item $\\bullet$ Further testing of the existing Cyborg geological exploration system at other geological sites with different types of imagery. \\item $\\bullet$ Speeding up the algorithm development by changing the project from being partly a hardware project with cameras and pan-tilt units and fieldwork to being entirely a software project without robotically-obtained image mosaics and without robotic interest-map pointing; with such a change in focus, our algorithms could be significantly enhanced by studying many more types of imagery: for example, from human geologist field studies on the Earth, from robotic geologist field studies on Mars, and from orbiter or flyby studies of our solar system's moons. \\end{trivlist} \\end{spacing}"
    },
    "0410/astro-ph0410537_arXiv.txt": {
        "abstract": "{I reexamine recent observations of velocity profiles across jets blown by young stellar objects, and argue that the observations do not support the interpretation of jets rotating around their symmetry axes. Instead, I propose that the interaction of the jets with a twisted-tilted (wrapped) accretion disk can form the observed asymmetry in the jets' line of sight velocity profiles. The proposed scenario is based on two plausible assumptions. (1) There is an inclination between the jet and the outer parts of the disk; the jet is perpendicular to the inner part of the disk; Namely, there is a twisted-tilted (wrapped) disk. (2) The disk-jet interaction slows down the jet as the jet entrains mass from the disk, with larger decelaration of jet segments closer to the tilted disk. The proposed scenario can account for the basic properties of the observed velocity profiles, while having the advantage that there is no need to refer to any magnetic jet launching model, and there is no need to invoke jet rotation with a huge amount of angular momentum. } ",
        "introduction": "\\label{sec:intro}  Asymmetry in the line of sight velocity across some jets ejected by young stellar objects (YSOs) have been interpreted as being due to a large scale rotation of the material in each jet around the jet's symmetry axis (e.g., Bacciotti et al.\\ 2002 and Coffey et al.\\ 2004, hereafter B2002 and C2004). A rotation around the jet axis is predicted by the magneto-centrifugal models for jet launching. In these models the magnetic fields that are anchored into the accretion disk-star system play a dominate role in accelerating the jet's material from the accretion disk. Most models in YSOs are based on the operation of large scale magnetic fields driving the flow from the disk; either via the ``centrifugal wind'' mechanism, first proposed by Blandford \\& Payne (1982), or from a narrow region in the magnetopause of the stellar field via an ``X-wind mechanism'' introduced by Shu et al. (1988, 1991) and in a somewhat different setting, by Ferreira \\& Pelletier (1993, 1995). See recent reviews and papers by K\\\"onigl and Pudritz (2000), Shu et al. (2000), Ferreira (2002), and Krasnopolsky et al. (2003). It should, however, be pointed out that the origin of the large scale magnetic fields and the manner that open field lines of sufficiently strong magnitude persist (in the centrifugal wind models), or the manner by which a stellar field interacts with the disk, allowing inflow and at the same time driving an outflow (in the X-wind models) are still open key issues of the theory (e.g. Heyvaerts et al.\\ 1996). In addition, it seems that thermal pressure is needed in some of these models (e.g., Ferreira \\& Casse 2004).  In a recent paper Cerqueira \\& de Gouveia Dal Pino (2004) conduct magnetohydrodynamics (MHD) 3D numerical simulations of YSO jets with rotation. They use parameters suitable for the jet of DG Tauri, a YSO studied by B2002, and obtain radial velocity maps which are in good agreement with the observations of B2002. In the appendix I point out some difficulties with this interpretation of jet rotation.  Magnetic fields might also play a role in triggering the jet's ejection event, e.g., by causing instabilities in the disk, and in collimating the jets (e.g., Heyvaerts \\& Norman 1989). The collimation issue is, however, still quite controversial and while magnetic collimation is certainly plausible, its exact nature is probably quite involved and still not fully understood (see the recent works of Bogovalov \\& Tsinganos 2001 and Li 2002). In some models external pressure are required to maintain collimation, e.g., Kato et al. (2004), and in the magnetic tower model of Lynden-Bell (2003). Although some of the arguments in these papers apply to relativistic jets, I mention them here as I am looking for a unified collimation machanism in different jets environments. In section (2) I discuss some aspects of disk-jet interaction in YSOs.  In some models, magnetic fields do not play a major role in launching the jet. These include the thermally-launched jet model (Torbett 1984; Torbett \\& Gilden 1992), that was applied to YSOs (Soker \\& Regev 2003) and cataclysmic variables (Soker \\& Lasota 2004). Motivated by non-magnetic models, in section (3) I show that the asymmetric radial velocity field of these YSO jets may be interpreted in an alternative way, where the magnetic field have no major role. In this proposed model, instead of jet rotation, the velocity field is attributed to the interaction of the jet with the accretion disk wind and/or corona. A summary of the main results is in section (4). ",
        "conclusions": "\\label{sec:sum}  The goal of the present paper is to provide an interpretation to the velocity profiles of some YSO jets without referring to any magnetic jet launching model. These observations (B2002; C2004) show that the two sides of some YSO jets have different projected velocity (line of sight velocity), and that the sense of the velocity shift is the same in the blue (approaching) and red (receding) jets. B2002 and C2004 interpret the velocity profiles as rotation of the jets around their symmetry axis. Such a rotation is expected in some magnetic jet launching models (Cerqueira \\& de Gouveia Dal Pino 2004).  In the appendix I show that the required angular momentum for the rotation explanation is on the upper limit of what the accretion disk can supply. In other words, to account for the required angular momentum and the kinetic energy of the jets, the mass outflow must be smaller than observed.  Alternatively, I propose a scenario where the velocity profile results from the interaction of the jets with the accretion disk. In section (2) I argue that such an interaction is likely to occur, as accretion disks around YSOs might have a thick surface layer. I then (section 3) presented a phenomenological scenario based on two assumptions. (1) There is an inclination between the jet and the disk further out. Namely, the inner disk is perpendicular to the jet, but then the disk flares in a point-symmetric manner: one side up and the opposite side (relative to the central star) down (figure 2). (2) The disk slows down the jet, with larger decelaration of jet segments closer to the tilted disk, and less efficient deceleration closer to the jets' axis. The velocity of the gas inside the jet as function of distance from the jet's symmetry axis $r$, and the angle relative to the direction where the disk flares up (in the $+z$ direction) $\\phi-\\phi_0$, was parameterized in equation ({\\ref{vjet1}}). The different values appearing in that equation are defined in figures 1 and 2. The calculated projection velocity profiles for 4 cases are plotted in figure 3.  The conclusions driven from figure 3 are summarized in the last paragraph of section 3. In short, these are: (1) The proposed non-magnetic scenario explains the decline in radial velocity with increasing distance from the jet's axis. (2) The inclined jet-disk scenario might explain the velocity shift between the two sides of each jet, with the same sense of shift in the blue shifted (approaching) and red shifted (receding) jets. (3) The large velocity variations between different spectral lines can be accounted for if different lines originate in different regions in the jets. For example, instabilities in the jet-disk interaction leads to the formation of dense clumps where some lines are favored over other lines. (4) There is no need to invoke jet rotation with a huge amount of angular momentum.  The interaction of the warrped-tilted disk with the jet will cause a point-symmetric structure. High spatial resolution imaging may reveal such a structure close to the star. In Th 28, the H$\\alpha$ image of the two HH objects possesses a point symemtric structure (Krauter 1986), which most likely comes from precessing jets. The two jet in RW Aur are highly symemtric, and show no point-symmetry  (Dougados et al.\\ 2000). The one jet of DG Tau as presented in the high veloicty (HV) map by Lavalley-Fouquet et al.\\ (2000; the upper row of their fig. 1) possesses some wiggling. In any case, the predicted point-symmetry morphology is of small degree, but still might be noticed.  Another prediction of the proposed scenario is that in some cases, where disk changes twisting (wrapping) direction in a relatively short time period, the sense of asymmetry will change. Namely, the side which presently has a lower radial velocity, might slowly, say on a time scale of $\\sim 100 \\yr$, change to have a higher radial velocity. However, in most disks the expected time scale is very long, according to the different mechanisms that cause such twisting (or titling or warping; Armitage \\& Pringle 1997; Larwood et al. 1996)."
    },
    "0410/gr-qc0410093_arXiv.txt": {
        "abstract": "The joint NASA-ESA mission LISA relies crucially on the stability of the three spacecraft  constellation. Each of the spacecraft is in heliocentric orbits forming a stable triangle. The principles of such a formation flight have been formulated long ago and analysis performed, but seldom presented if ever, even to LISA scientists. We nevertheless need these details in order to carry out theoretical studies on the optical links, simulators etc. In this article, we present in brief, a model of the LISA constellation, which we believe will be useful for the LISA community. ",
        "introduction": "LISA is a ESA-NASA mission for observing low frequency gravitational waves in the frequency range from $10^{-5}$ Hz to 1 Hz \\cite{RIP}. In order for LISA to operate successfully, it is crucial that the three spacecraft which form the hubs of the laser interferometer in space maintain nearly constant distances between them, though their order of magnitude is $5\\times 10^6$ km. The existence of orbits having this property was firstly reported by Bender \\cite{PB} as the basis of LISA. In order to thoroughly study the optical links and light propagation between these moving stations, we however need a detailed model of the LISA configuration. We therefore find it useful to recall explicitly the not so trivial principles of a stable formation flight. In this brief work, we firstly study three Keplerian orbits around the Sun with small eccentricities and adjust the orbital parameters so that the spacecraft form an equilateral triangle with nearly constant distances between them. Then we find that to the first order in the parameter $\\a = l/2R$, where $l \\sim 5 \\times 10^6$ km, is the distance between two spacecraft and $R = 1$ A. U. $\\sim 1.5 \\times 10^8$ km, the distances between spacecraft are exactly constant; any variation in arm-lengths should result from higher orders in $\\a$ or from external perturbations of Jupiter and the secular effect due to the Earth's gravitational field. (The eccentricity $e$ is related in a simple way to $\\a$ and is proportional to $\\a$ to the first order in $\\a$.) In fact our analysis shows that such formations are possible with any number of spacecraft provided they lie in a magic plane making an angle of $60^{\\circ}$ with the ecliptic. We establish this general result with the help of the Hill's or Clohessy-Wiltshire (CW) equations \\cite{CW}. ",
        "conclusions": "We have explicitly constructed three heliocentric spacecraft orbits which to the first order in eccentricities maintain equal distances between them which is taken to be 5 million km. We have shown with the help of a more sophisticated formalism - the CW equations - that there are two planes in the CW frame, in which particles obeying the CW equations and satisfying stability requirements, namely, no offsets (and hence no runaway behaviour) and maintaining equal distance from the origin, maintain their relative distances in the CW approximation which is equivalent to a first order calculation in the eccentricities. This has the implication that formations not necessarily triangular and with any number of spacecraft are possible as long as they obey the stability constraints and lie in any one of these planes; their relative distances will be maintained within the CW approximation. This result opens up new possibilities of spacecraft constellations with various geometrical configurations and any number of spacecraft which would be useful to future space missions. \\ack SD would like to thank IFCPAR for travel support and the Observatoire de la C\\^ote D'Azur, Nice, France for local hospitality, where substantial part of this work was carried out. SK would like to acknowledge DST, India for the WOS-A. All the authors would like to thank E. Chassande-Mottin for detailed discussions. \\vspace{24pt}"
    },
    "0410/astro-ph0410701_arXiv.txt": {
        "abstract": "H1504+65 is a young white dwarf with an effective temperature of 200\\,000\\,K and is the hottest post-AGB star ever analysed with detailed model atmospheres. {\\it Chandra} LETG+HRC-S spectra have revealed the richest X-ray absorption line spectrum recorded from a stellar photosphere to date. The line forming regions in this extremely hot photosphere produce many transitions in absorption that are also observed in emission in cool star coronae. We have performed a detailed comparison of {\\it Chandra} spectra of H1504+65 with those of Procyon and $\\alpha$\\,Cen~A and B. State of the art non-LTE model spectra for the hot white dwarf have enabled us to identify a wealth of absorption lines from highly ionised O, Ne and Mg. In turn, these features have allowed us to identify coronal lines whose origins were hitherto unknown. ",
        "introduction": "High-resolution X-ray spectroscopy performed with {\\it Chandra} and XMM-Newton allows very detailed studies of coronae about cool stars. While many individual emission lines were detected for the first time in stellar spectra by the Extreme Ultraviolet Explorer Satellite (EUVE; see, e.g., Drake \\etal 1996), the resolving power of $\\lambda/\\Delta\\lambda\\sim 200$ of the EUVE spectrographs was a quite modest compared with that of present day X-ray observatories.  In particular, the unprecedented resolution capabilities of the {\\it Chandra} X-ray Observatory Low Energy Transmission Grating Spectrograph (LETG) in the 30-170~\\AA\\ range ($\\lambda/\\Delta\\lambda\\sim 1000$) that overlaps with the EUVE Short Wavelength spectrometer (70-170~\\AA), have revealed many more weak spectral lines.  Owing largely to a historical lack of attention to the 30-170~\\AA\\ range, a large fraction of these remain unidentified.  Identification of these features is desirable because they could be used as a spectroscopic diagnostics, because they potentially contribute to the flux of diagnostic lines currently employed, and because they contribute to the overall plasma radiative loss.  Two nearby stars that have illuminated the forest of lines in the 30-170~\\AA\\ range are \\alp\\ (G2V+K1V) and \\pp\\ (F5IV).  All three stars exhibit classical solar-like X-ray emitting coronae.  Indeed, analogues of the relatively X-ray faint Sun are difficult to observe because they become unreachable with current instrumentation beyond a few parsecs, and \\alp\\ and \\pp\\ represent the nearest and brightest coronal sources with solar-like activity.  Only a small fraction of the multitude of lines between 30-170~\\AA\\ seen their {\\it Chandra} LETG spectra could be identified based on current radiative loss models (Raassen \\etal\\ 2002, 2003).  Drake \\etal (in prep.) have estimated that these models underestimate the true line flux in the range 30-70~\\AA\\ in these stars by factors of up to 5 or so.  The ``missing lines'' are predominantly transitions involving $n=2$ ground states in abundant elements such as Ne, Mg, Si, S and Ar---the analogous transitions to the Fe ``L-shell'' lines between $\\sim 8$-18~\\AA, together with Fe $n=3$ (the ``M-shell'') transitions (Drake 1996, Jordan 1996).  Some of these lines have been identified based on Electron Beam Ion Trap experiments (Beiersdorfer \\etal 1999, Lepson \\etal 2002, 2003).  In the present paper we approach this problem from a new perspective, namely through a {\\it Chandra} observation of the photosphere of the hottest white dwarf (WD) known, \\hh, and its quantitative analysis by means of detailed non-LTE model atmospheres.  \\hh\\ has an effective temperature of 200\\,000\\,K. It belongs to the PG1159 spectral class, which are hot, hydrogen-deficient (pre-) white dwarfs. Their surface chemistry (typical abundances: He=33\\%, C=48\\%, O=17\\%,\\\\ Ne=2\\%, mass fractions) suggests that they exhibit matter from the helium-buffer layer between the H- and He-burning shells in the progenitor AGB star (Werner 2001). This is likely because the PG1159 stars have suffered a late He-shell flash, a phenomenon that drives the fast evolutionary rates of such famous stars like FG\\,Sge and Sakurai's object.  \\hh\\ is in fact a peculiar member of this class, because it is also helium-deficient.  Its atmosphere is mainly composed of carbon and oxygen plus neon and magnesium (C=48\\%, O=48\\%, Ne=2\\%, O=2\\%, mass fractions).  \\hh\\ is a unique object, considering its high \\Teff\\ and chemical surface composition, and we have speculated that it represents the naked C/O core of a former red giant (Werner \\etal 2004, W04).  {\\it Chandra} LETG+HRC-S spectra from \\hh\\ have revealed the richest X-ray absorption line spectrum record\\-ed from a stellar photosphere to date. We have recently performed a detailed analysis of this spectrum (W04, Fig.~\\ref{fig_chandra_fit}) and we use in the paper in hand the photospheric spectrum of \\hh\\ together with an appropriate model atmosphere to identify a number of emission lines in the coronae of \\ala, \\alb, and Procyon. The difference in particle densities in the WD photosphere and in the coronae amounts to many orders of magnitude (roughly n$_{\\rm e}$=$10^{10}$ and $10^{13}-10^{18}$\\,cm$^{-3}$, respectively), however, the temperature in the line forming regions of the WD (up to 300\\,000\\,K) is comparable to the low-temperature component of multi-temperature fits to coronae, required to account for the lines of low-ionisation stages (e.g.\\ 630\\,000\\,K for Procyon; Raassen \\etal 2002). As a consequence, numerous lines from \\ion{O}{vi}, \\ion{Ne}{vi-viii} and \\ion{Mg}{vi-ix} are visible in the soft X-ray spectra of both, the cool star coronae (in emission) and the hot WD photosphere (in absorption). Lines from higher ionisation stages are formed in the high-temperature regions of the coronae (T of the order 1--2.5 million K for the stars studied in this paper), hence, their respective absorption line counterparts cannot be formed in the WD photosphere. Fig.\\,\\ref{fig_ion} shows the temperature and particle density structure of a model for \\hh.  In the following, we first introduce briefly the characteristics of the objects studied here.  We describe our model atmosphere calculation for the hot WD, concentrating on the atomic data employed. We then perform a detailed comparison of the absorption and emission line spectra and suggest a number of new line identifications for the cool star coronae.  \\begin{figure}[tbp] \\resizebox{\\hsize}{!}{\\includegraphics{kwernerf3.ps}} \\caption[]{ Comparison of {\\it Chandra} X-ray spectra of \\hh\\ and Procyon. Lines from \\ion{Mg}{vii} and \\ion{Ne}{vii} are in absorption in \\hh\\ and in emission in Procyon. Top: photosphere model for \\hh\\ with line identifications for \\ion{Mg}{vii} and \\ion{Ne}{vii}. Middle: Degraded model spectrum (i.e.\\ folded with a 0.05\\AA\\ FWHM Gaussian) plotted over \\hh\\ observation. Bottom: Procyon spectrum with line identifications from Raassen \\etal (2002). {\\it Chandra} spectra were smoothed with a 0.03\\AA\\ boxcar.  } \\label{fig_84} \\end{figure}  \\begin{figure*}[tbp] \\includegraphics[width=\\textwidth]{kwernerf4.ps} \\caption[]{Comparison of {\\it Chandra} X-ray spectra of \\hh\\ with Procyon and \\ala, similar to Fig.\\,1. All shown lines from highly ionised Ne and Mg are identified for the first time in the cool star corona, except for \\ion{Mg}{viii}~86.85/87.02\\AA\\ and the \\ion{Ne}{viii} doublet at 88.1\\AA, which were identified by Raassen \\etal (2002, 2003). {\\it Chandra} spectra of \\hh\\ and the coronae were smoothed with 0.03\\AA\\ and 0.05\\AA\\ boxcars, respectively.  } \\label{fig_88} \\end{figure*}   \\begin{table*} \\caption{ List of X-ray multiplets observed in both the \\hh\\ photosphere and in the coronae of either \\ala\\ (``A''), \\alb\\ (``B''), or \\pp (``P''). In the last column, ``R'' denotes that the line identification was performed in Raassen \\etal (2002, 2003), whereas ``N'' denotes a new identification suggested in this paper. ``R, N'' means that at least one component of the multiplet is newly identified here. Expressions in brackets denote doubtful cases. The column ``Source'' gives the reference to the level energies of the transition. After each transition we have marked, if the lower level is a ground state (``G'') or a metastable state (``M''). \\label{lines_tab} } \\begin{tabular}{lllr@{\\,\\,--\\,\\,}llll} \\hline \\hline \\noalign{\\smallskip} $\\lambda$/\\AA\\ (\\hh\\ model)     & seen in  & Ion & \\multicolumn{2}{c}{Transition} & & Source & Remark \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 69.41, .47, .57                & A,B,P  & \\ion{Mg}{viii} & 2p\\hse\\,$^2$P$^{\\rm o}$  & 3p\\hse\\,$^2$D\\hse       &G & {\\sc Nist} & N\\\\ 74.27, .32, .34, .37, .41, .43 & A,B,P  & \\ion{Mg}{viii} & 2p$^2$\\,$^4$P\\hse        & 3d\\hse\\,$^4$D$^{\\rm o}$ &M & {\\sc Nist} & N\\\\ 74.78, .81, .87                & A,B    & \\ion{Ne}{vii}  & 2p\\hsed\\,$^3$P$^{\\rm o}$ & 4p\\,$^3$D\\hse           &M & {\\sc Nist} & N\\\\ 74.86, 75.03, .04              & A,B,P  & \\ion{Mg}{viii} & 2p\\hse\\,$^2$P$^{\\rm o}$  & 3d\\hse\\,$^2$D\\hse       &G & {\\sc Nist} & R\\\\ (78.34), 78.41, 78.52          & A,B,P  & \\ion{Mg}{vii}  & 2p$^3$\\,$^3$P\\hse        & 3p\\hse\\,$^3$P$^{\\rm o}$ &G & Kelly      & N\\\\ 80.23, .25                     & A,B,P  & \\ion{Mg}{viii} & 2p$^2$\\,$^2$D\\hses       & 3d\\hse\\,$^2$D$^{\\rm o}$ &  & {\\sc Nist} & R\\\\ 80.95, 81.02, .14              & A,(B,P)& \\ion{Mg}{vii}  & 2p$^3$\\,$^3$P\\hse        & 3p\\hse\\,$^3$S$^{\\rm o}$ &G & Kelly      & N\\\\ 81.37                          & (A),B,P& \\ion{Ne}{vii}  & 2p\\hsw\\,$^1$P$^{\\rm o}$  & 4p\\,$^1$P\\hse           &  & Kelly      & N\\\\ 81.73, .79, .84, .87, .94, .98 & A,B,P  & \\ion{Mg}{viii} & 2p$^2$\\,$^4$P\\hse        & 3s\\hse\\,$^4$P$^{\\rm o}$ &M & {\\sc Nist} & N\\\\ 82.17, .20, .27                & A,B,(P)& \\ion{Ne}{vii}  & 2p\\hse\\,$^3$P$^{\\rm o}$  & 4d\\,$^3$D\\hse           &M & Kelly      & N\\\\ (82.60), .82                   & A,B,P  & \\ion{Mg}{viii} & 2p\\hse\\,$^2$P$^{\\rm o}$  & 3s\\hse\\,$^2$S\\hse       &G & {\\sc Nist} & R\\\\ 83.51, .56, .59, .64, .71, .76 & A,B,P  & \\ion{Mg}{vii}  & 2p$^3$\\,$^3$P\\hse        & 3d\\hse\\,$^3$P$^{\\rm o}$ &G & Kelly      & R, N\\\\ 83.91, .96, .99, 84.02, .09, .11&A,B,P  & \\ion{Mg}{vii}  & 2p$^3$\\,$^3$P\\hse        & 3d\\hse\\,$^3$D$^{\\rm o}$ &G & Kelly      & R, N\\\\ (84.19, .23,) .30              & A      & \\ion{Ne}{vii}  & 2p\\hse\\,$^3$P$^{\\rm o}$  & 4s\\,$^3$S\\hse           &M & Bashkin    & R\\\\ 85.41                          & (A,B,P)& \\ion{Mg}{vii}  & 2p$^2$\\,$^1$D\\hses       & 3d\\hse\\,$^1$F$^{\\rm o}$ &M & Kelly      & N\\\\ 86.82                          & A,B,P  & \\ion{Ne}{vii}  & 2p$^2$\\,$^1$D\\hses       & 4d\\,$^1$F$^{\\rm o}$     &M & Kelly      & N\\\\ 86.84, .85, 87.02              & A,B,P  & \\ion{Mg}{viii} & 2p$^2$\\,$^2$D\\hses       & 3s\\hse\\,$^2$P$^{\\rm o}$ &  & {\\sc Nist} & R, N\\\\ 87.46                          & A      & \\ion{Ne}{vii}  & 2s$^2$\\,$^1$S\\hsed       & 3s\\,$^1$P$^{\\rm o}$     &G & {\\sc Nist} & N\\\\ 87.72                          & A      & \\ion{Mg}{vii}  & 2p$^2$\\,$^1$D\\hses       & 3d\\hse\\,$^1$D$^{\\rm o}$ &M & Kelly      & N\\\\ 88.08, 88.12                   & A,B,P  & \\ion{Ne}{viii} & 2s\\hse\\,$^2$S\\hse        & 3p\\,$^2$P$^{\\rm o}$     &M & {\\sc Nist} & R\\\\ 88.68                          & (A),B,P& \\ion{Mg}{vii}  & 2p$^2$\\,$^1$S\\hsed       & 3d\\hse\\,$^1$P$^{\\rm o}$ &M & Kelly      & N\\\\ 89.64, .65                     & A,(P)  & \\ion{Mg}{vi}   & 2p$^3$\\,$^2$P$^{\\rm o}$  & 4s\\hse\\,$^2$P           &M & Kelly      & N\\\\ 91.56                          & P      & \\ion{Ne}{vii}  & 2p\\hsw\\,$^1$P$^{\\rm o}$  & 4s\\,$^1$S\\hse           &  & Kelly      & R\\\\ 92.13, .32                     & A,B,P  & \\ion{Mg}{viii} & 2p$^2$\\,$^2$S\\hsed       & 3s\\hse\\,$^2$P$^{\\rm o}$ &  & {\\sc Nist} & R, N\\\\ 92.85                          & P      & \\ion{Ne}{vii}  & 2p$^2$\\,$^1$S\\hsed       & 4d\\,$^1$P$^{\\rm o}$     &M & Kelly      & R\\\\ (93.89), 94.07, .10, (.27)     & A,B,P  & \\ion{Mg}{viii} & 2p\\hse\\,$^2$P\\hse        & 3s\\hse\\,$^2$P$^{\\rm o}$ &  & {\\sc Nist} & N\\\\ 94.04, (.17, .24)              & A,B,P  & \\ion{Mg}{vii}  & 2p$^3$\\,$^5$S$^{\\rm o}$  & 3s\\hse\\,$^5$P\\hse       &M & Kelly      & N\\\\ 94.26, .27, .30, .31, .36, .39 & B      & \\ion{Ne}{vii}  & 2p\\hse\\,$^3$P$^{\\rm o}$  & 3p\\,$^3$P\\hse           &M & Bashkin    & N\\\\ 95.03, .04                     & B      & \\ion{Mg}{vii}  & 2p$^3$\\,$^3$D$^{\\rm o}$  & 3s'\\,$^3$D              &  & Kelly      & N\\\\ 95.26, .38, .42, .49, .56, .65 & (A,B,P)& \\ion{Mg}{vii}  & 2p$^3$\\,$^3$P\\hse        & 3s\\hse\\,$^3$P$^{\\rm o}$ &G & Kelly      & N\\\\ (95.38, .42, .48)              & (A,B,P)& \\ion{Mg}{vi}   & 2p$^3$\\,$^4$S$^{\\rm o}$  & 3d\\hse\\,$^4$P           &G & Kelly      & R\\\\ 95.75, .81, .89, .90, .91, 96.0&A,B,P   & \\ion{Ne}{vii}  & 2p\\hse\\,$^3$P$^{\\rm o}$  & 3p\\,$^3$D\\hse           &M & Bashkin    & N\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*}  \\addtocounter{table}{-1} \\begin{table*} \\caption{continued} \\begin{tabular}{lllr@{\\,\\,--\\,\\,}llll} \\hline \\hline \\noalign{\\smallskip} $\\lambda$/\\AA\\ (\\hh\\ model)     & seen in  & Ion & \\multicolumn{2}{c}{Transition} & & Source & Remark \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 96.08, .09                     & (A,B,P)& \\ion{Mg}{vi}   & 2p$^3$\\,$^2$P$^{\\rm o}$  & 3d''\\,$^2$D             &M & Kelly      & N\\\\ 97.50                          & A,B,P  & \\ion{Ne}{vii}  & 2s$^2$\\,$^1$S\\hse        & 3p\\,$^1$P$^{\\rm o}$     &G & Kelly      & R\\\\ 98.11, .26                     & A,B,P  & \\ion{Ne}{viii} & 2p\\hse\\,$^2$P$^{\\rm o}$  & 3d\\,$^2$D\\hse           &  & {\\sc Nist} & R\\\\ 98.50, .51                     & B      & \\ion{Mg}{vi}   & 2p$^3$\\,$^2$P$^{\\rm o}$  & 3d'\\,$^2$S              &M & Kelly      & N\\\\ 99.69                          & B      & \\ion{O}{vi}    & 2s\\hses                  & \\hsx  6p                &G & Kelly      & N\\\\ 100.70, .90                    & A      & \\ion{Mg}{vi}   & 2p$^3$\\,$^2$D$^{\\rm o}$  & 3d\\,$^2$F               &M & Kelly      & N\\\\ 101.49, .55                    & B      & \\ion{Mg}{vi}   & 2p$^3$\\,$^2$D$^{\\rm o}$  & 3d\\,$^2$P               &M & Kelly      & N\\\\ 102.91, 103.08                 & A,B,P  & \\ion{Ne}{viii} & 2p\\hse\\,$^2$P$^{\\rm o}$  & 3s\\,$^2$S\\hse           &  & {\\sc Nist} & R\\\\ 103.09                         & (A,B,P)& \\ion{Ne}{vii}  & 2p\\hsw\\,$^1$P$^{\\rm o}$  & 3p\\,$^1$D\\hse           &  & Kelly      & N\\\\ 104.81                         & B,P    & \\ion{O}{vi}    & 2s\\hses                  & \\hsx  5p                &G & Kelly      & R\\\\ 105.17                         & A,(B)  & \\ion{Mg}{vii}  & 2p$^3$\\,$^1$D$^{\\rm o}$  & 3s'\\,$^1$D\\hse          &  & Kelly      & N\\\\ 106.03, .08, .19               & P      & \\ion{Ne}{vii}  & 2p\\hse\\,$^3$P$^{\\rm o}$  & 3d\\,$^3$D\\hse           &M & Kelly      & R, N\\\\ (111.10, .16), .26             & A,B,P  & \\ion{Ne}{vi}   & 2p\\hse\\,$^2$P$^{\\rm o}$  & 3p\\,$^2$D\\hse           &G & Kelly      & N\\\\ 111.15                         & (A),B,P& \\ion{Ne}{vii}  & 2p$^2$\\,$^1$D\\hses       & 3d\\,$^1$P$^{\\rm o}$     &M & Kelly      & N\\\\ 111.55, .75, .86               & B,(A,P)& \\ion{Mg}{vi}   & 2p$^3$\\,$^4$S$^{\\rm o}$  & 3s\\hse\\,$^4$P           &G & Kelly      & R\\\\ (115.33), .39, (.52)           & A,B,P  & \\ion{Ne}{vii}  & 2p\\hse\\,$^3$P$^{\\rm o}$  & 3s\\,$^3$S\\hse           &M & Kelly      & R\\\\ 115.82, .83                    & B      & \\ion{O}{vi}    & 2s\\hses                  & \\hsx  4p                &G & Kelly      & R\\\\ 115.96                         & B      & \\ion{Ne}{vii}  & 2p$^2$\\,$^1$D\\hses       & 3d\\,$^1$D$^{\\rm o}$     &M & Kelly      & N\\\\ (116.35), .42                  & B      & \\ion{O}{vi}    & 2p                       & \\hsx  5d                &  & Kelly      & N\\\\ 116.69                         & B      & \\ion{Ne}{vii}  & 2p\\hsw\\,$^1$P$^{\\rm o}$  & 3d\\,$^1$D\\hse           &  & Kelly      & R\\\\ 116.97, 117.22                 & A      & \\ion{Mg}{vi}   & 2p$^3$\\,$^2$P$^{\\rm o}$  & 3s'\\,$^2$D              &M & Kelly      & N\\\\ (117.33), .40                  & B      & \\ion{O}{vi}    & 2p                       & \\hsx  5s                &  & Kelly      & N\\\\ (117.43), .66, (.78)           & P      & \\ion{Mg}{vii}  & 2p$^3$\\,$^3$S$^{\\rm o}$  & 3s\\hse\\,$^3$P\\hse       &  & Kelly      & N\\\\ (117.52), .64, (.81)           & P      & \\ion{Mg}{vii}  & 2p$^3$\\,$^3$P$^{\\rm o}$  & 3p\\hse\\,$^3$P\\hse       &  & Kelly      & N\\\\ 120.20, .27, .33, .35, .42, .48& P      & \\ion{Ne}{vii}  & 2p$^2$\\,$^3$P\\hse        & 3s\\,$^3$P$^{\\rm o}$     &  & Kelly      & N\\\\ 122.49, .69                    & B,P    & \\ion{Ne}{vi}   & 2p\\hse\\,$^2$P$^{\\rm o}$  & 3d\\,$^2$D\\hse           &G & Kelly      & R, N\\\\ 123.59                         & P      & \\ion{Mg}{vi}   & 2p$^4$\\,$^2$D\\hse        & 3s$^{\\rm iv}$\\,$^2$D$^{\\rm o}$& &Kelly & N\\\\ 127.67                         & B,P    & \\ion{Ne}{vii}  & 2p\\hsw\\,$^1$P$^{\\rm o}$  & 3s\\,$^1$S\\hse           &  & Kelly      & R\\\\ 129.78, .87                    & A,B,P  & \\ion{O}{vi}    & 2p                       & \\hsx  4d                &  & Kelly      & R\\\\ 130.31, .64                    & B      & \\ion{Mg}{vi}   & 2p$^4$\\,$^2$P\\hse        & 3s$^{\\rm v}$\\,$^2$P$^{\\rm o}$& &Kelly  & N\\\\ 130.94, 131.09, .30            & A,B,P  & \\ion{Mg}{vii}  & 2p$^3$\\,$^3$S$^{\\rm o}$  & 3p\\hse\\,$^3$P\\hse       &  & Kelly      & N\\\\ 132.22, .31                    & A,B    & \\ion{O}{vi}    & 2p                       & \\hsx  4s                &  & Kelly      & N\\\\ 150.09, .12                    & B,P    & \\ion{O}{vi}    & 2s\\hses                  & \\hsx  3p                &G & Kelly      & R\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We have performed a detailed comparison of {\\it Chandra} soft X-ray spectra from the photosphere of the hottest known white dwarf, \\hh, with the corona spectra of \\ala, \\alb, and \\pp. With the help of a detailed model spectrum for \\hh\\ we have found that a large number of lines from multiplets of O, Ne, and Mg are present in both the photospheric absorption line spectrum and the coronal emission line spectra. In the coronal spectra we have newly identified lines from about 40 multiplets of \\ion{O}{vi}, \\ion{Ne}{vi-vii}, and \\ion{Mg}{vi-viii}. Some of these lines are blends with previously known lines, which are in use for diagnostic purposes, hence, their contribution to the line flux must be considered in detailed spectral analyses.  A more complete version of this paper will be published in Astronomy \\& Astrophysics."
    },
    "0410/astro-ph0410192.txt": {
        "abstract": "We present an analysis of the properties of the hot interstellar medium (ISM) in the merging pair of galaxies known as The Antennae (NGC 4038/39), performed using the deep, coadded $\\sim$411 ks {\\it Chandra} ACIS-S data set. These deep X-ray observations and {\\it Chandra}'s high angular resolution allow us to investigate the properties of the hot ISM with unprecedented spatial and spectral resolution. Through a spatially resolved spectral analysis, we find a variety of temperatures (from 0.2 to 0.7 keV) and $N_H$ (from Galactic to $2\\times10^{21}$ cm$^{-2}$). Metal abundances for Ne, Mg, Si, and Fe vary dramatically throughout the ISM from sub-solar values ($\\sim$0.2) up to several times solar. %Measures for the hot-gas mass ($\\sim$10$^7 M_\\sun$), cooling times %($10^7$--$10^8$ yr), %and pressure are derived.  In the two nuclei the hot-gas pressure is significantly %higher than the CO pressure, implying that shock waves may be driven into the %CO clouds. %Comparison of the abundances with the average stellar yields predicted by %theoretical %models of SN explosions points to SNe ot Type II as the main %contributor of metals to the hot ISM. %There is no evidence of any correlation between radio-optical star formation %indicators and the metal abundances. Although uncertainties in the average %density cannot exclude that mixing may have played some important role, %the time required to produce the observed metal masses ($\\la3$ Myr) %suggests that the correlations %are unlikely to be destroyed by the presence of efficient mixing. More likely %a significant fraction of SN~II ejecta may be in a %cool phase, in grains, or escaping in the wind. ",
        "introduction": "The Antennae, dynamically modeled by Toomre \\& Toomre (1972) and Barnes (1988), are the nearest pair of colliding galaxies involved in a major merger (D = 19~Mpc for H$_o = 75~\\rm km~s^{-1}~Mpc^{-1}$). Hence, this system provides a unique opportunity for getting the most detailed look possible at the consequences of a galaxy merger as evidenced by induced star formation and the conditions in the ISM. Each of the two colliding disks shows rings of giant H~II regions and bright stellar knots with luminosities up to $M_V \\sim -16$ (Rubin, Ford, \\& D'Odorico 1970), which are resolved with  the  {\\it Hubble Space Telescope (HST)} into typically about a dozen young star clusters (Whitmore \\& Schweizer 1995; Whitmore et al.\\ 1999). These knots coincide with the peaks of H$\\alpha$, 2.2$\\mu$, and 6-cm radio-continuum emission (Amram et al.\\ 1992; Stanford et al.\\ 1990; Neff \\& Ulvestad 2000), indicating an intensity of star formation in each single region exceeding that observed in 30 Doradus.  CO aperture synthesis maps reveal major concentrations of molecular gas, including a $\\sim$2.7$\\times 10^9 M_{\\odot}$ concentration in the region where the two disks overlap (Stanford et al.\\ 1990; Wilson et al.\\ 2000). This concentration in the overlap region is coincident with the bright emission peaks seen in {\\it SCUBA} 450$\\mu$ and 850$\\mu$ maps and {\\it ISOPHOT} 60$\\mu$ and 100$\\mu$ observations (Haas et al.\\ 2000), which show the simultaneous presence of warm dust ($T\\approx30$ K), usually observed in starbursts of luminous IR galaxies, and cold dust ($T<20$ K) more typical for quiet galaxies and extremely dense clouds. A K-band study derives ages for the star clusters ranging from 4 to 13~Myr, and yields high values of extinction ($A_V \\sim 0.7 - 4.3$~mag; Mengel et al.\\ 2000, 2005). Many of the youngest star clusters lie in the CO-rich overlap region, where even higher extinction is present ($A_V \\la 100$ mag), implying that they must lie in front of most of the gas (Wilson et al.\\ 2003). Some of them are found as far as 2~kpc from detectable CO regions, leading to the conclusion that either they formed from clouds less massive than $5\\times10^6$ $M_\\odot$ or they have already destroyed their parent molecular cloud, or perhaps both.  The presence of an abundant hot ISM in The Antennae was originally suggested by the first {\\it Einstein} observations of this system (Fabbiano \\& Trinchieri 1983), and has since been confirmed with several major X-ray telescopes ({\\it ROSAT}: Read, Ponman, \\& Wolstencroft 1995, Fabbiano, Schweizer, \\& Mackie 1997; and {\\it ASCA}: Sansom et al.\\ 1996). The {\\it ROSAT} HRI image was used in a recent multiwavelength study of The Antennae by Zhang, Fall, \\& Whitmore (2001), which suggests feedback between star clusters and the interstellar medium (ISM). The first {\\it Chandra} ACIS (Weisskopf, O' Dell \\& van Speybroeck 1996) observation of The Antennae in 1999 (Fabbiano, Zezas \\& Murray 2001; Fabbiano et al.\\ 2003, hereafter F03) gave us for the first time a detailed look at this hot ISM, revealing a complex, diffuse and soft emission component responsible for about half of the detected X-ray photons from the two merging galaxies. The  spatial resolution of {\\it Chandra} is at least 10 times superior to that of any other X-ray observatory, allowing us to resolve the emission on physical scales of $\\sim$75~pc for D$=$19~Mpc, and to detect and subtract individual point-like sources (most likely X-ray binaries; see Zezas et al.\\ 2002). At the same time, ACIS allows us to study for the first time the X-ray spectral properties of these emission regions, providing additional important constraints on their nature.  In the first detailed {\\it Chandra} study of the hot ISM, F03, using the 1999 data, presented a ``multi-color'' X-ray image that suggests both extensive absorption by the dust in this system, peaking in the overlap region, as well as variations in the temperature of different emission regions of the hot ISM. These conclusions were confirmed by spectral fits to the data extracted from different regions, which suggested that at least two thermal components with temperatures of $\\sim$0.3 and $\\sim$0.7~keV are present in the luminous regions of the diffuse emission. The parameters derived for the hot ISM suggested that in the nuclear regions of the two colliding galaxies the mass of the hot ISM is $\\sim$10$^5$--10$^6 M_{\\odot}$, much smaller than that in molecular gas (Wilson et al.\\ 2000), but that the thermal pressures of the two phases of the ISM are comparable, suggesting equilibrium. Comparisons with H$\\alpha$, radio-countinuum (Neff \\& Ulvestad 2000), and HI data (Hibbard et al.\\ 2001) suggested a displacement of the diffuse X-ray emission on the northen side of the system, consistent with the effect of ram pressure from the cold ISM (H~I) raining onto the hot gas in the northern star-forming arm. F03 did not make any conclusive statement on the metal abundance of the hot ISM. Metz et al. (2004) analyzed the 1999 data together with the second and third observations, concentrating on regions of high star-formation rate, and found that super star clusters play a significant role in heating the ISM, but that the amount of hot interstellar gas is not directly proportional to the cluster mass.  While this first {\\it Chandra} data set demonstrated the richness of the ISM in The Antennae, the number of detected photons was such that the detailed small-scale morphology and spectral properties could not be fully explored. We now have completed a deep monitoring observing campaign of The Antennae with {\\it Chandra} ACIS, that has produced a very detailed and rich data set. Our first look at these data revealed an intricate diffuse emission, with clear signatures of strong line emission in some regions, pointing to high metal abundances of the hot ISM  (Fabbiano et al.\\ 2004). In the present follow-up paper, we report a detailed analysis of this diffuse emission. The paper is organized as follows. In \\S2 we describe the {\\it Chandra} ACIS-S observations of The Antennae and the procedures followed to screen and prepare the data for analysis. The creation of a broad-band mapped-color image of the diffuse emission in The Antennae is described in \\S3, while the preparation of a mapped-color line-strength image---useful for detecting differences in the spatial distribution of the metals---is described in \\S4. In \\S5 we present the spectral analysis performed after subdividing the distribution of diffuse emission into 21 regions, following the indications from the mapped-color images. We then discuss the adopted spectral models and the method used for computing the errors, present the results of the spectral fits, and compare our findings with previous results obtained by Sansom et al.\\ (1996) and F03. %The hot gas parameters, like pressure and density, derived from the spectral %fits %are given in \\S6. In \\S7, we discuss our results, focusing especially on the %element abundances and on the constraints that our analysis sets on the %physical properties of the hot ISM. We summarize the results of our analysis in \\S6. In a companion paper (Baldi et al., submitted) we discuss the implications of these results for our understanding of the physical conditions and metal enrichment in the hot ISM. ",
        "conclusions": "We have performed a detailed study of the X-ray properties of the diffuse emission of The Antennae (NGC~4038/39), analyzing the entire 411~ks exposure obtained as part of our monitoring observing campaign of this system with {\\it Chandra} ACIS. Confirming the results of F03, which were based on the first of the seven observations used in the present study, we report a spatially and spectrally complex hot ISM. Thanks to our deep data, we also detect clear, spatially variable emission lines of Ne, Mg, and Si, in addition to the Fe-L blend, which indicate super-solar metal abundances in a few regions of the hot ISM.  In summary:  (1) We derive an X-ray broad-band color image of the hot ISM, as well as a color map displaying the apparent strength of line emission from O+Fe+Ne, Mg-XI, and Si-XIII. Both maps demonstrate complexity and have been used to guide our selection of 21 separate spectral extraction regions for detailed analysis.  (2) Spectral models combining optically thin thermal-plasma and power-law emission can be fitted to our data with the presence, in almost all regions, of a single thermal component with temperatures ranging from 0.2 to 0.7~keV. Significant variations of $N_{\\rm H}$ throughout the ISM are detected: column densities range from values consistent with the Galactic foreground absorption to higher values of $\\sim${}$2\\times10^{21}$ cm$^{-2}$, as in the southern nucleus and in Region 7, the actively star-forming Overlap Region (Mirabel et al.\\ 1998; Wilson et al.\\ 2000; Zhang et al.\\ 2001). A power-law component is present in half of the regions, showing values of $\\Gamma$ consistent with emission from unresolved X-ray binaries. The study of this power-law emission will be the subject of a forthcoming paper. %Summing the hot-gas masses determined for each region yields a %total measured mass of $\\sim${}$1.2\\times 10^7$ M$_\\sun$ for the hot ISM. %This hot gas has cooling times lying in the $10^7$--$10^8$ yr range, %in agreement with the results of F03.  (3) Fitting the Fe-L, Ne-IX, Mg-XI and Si-XIII emission, we find significant metal enrichment of the hot ISM. Metal abundances are generally consistent with solar, but reach extremes of subsolar and $\\sim$20--30 solar in a few regions. In supporting evidence, the $EW$ of the emission lines of Mg and Si correlates well with the abundances measured. Previous reports (from ASCA) of subsolar abundances in the hot ISM of the Antennae can be explained with an \"averaging\" effect due to the inclusion of the entire emission (diffuse from different regions and point sources) in a single spectrum. %The total hot gas masses derived for Fe, Ne, Mg, and Si are %$3.9\\times10^4$ M$_\\sun$, $5.4\\times10^4$ M$_\\sun$, $2.4\\times10^4$ M$_\\sun$, %and $2.4\\times10^4$ M$_\\sun$, respectively.   The implication of these results for the physical state of the hot ISM and SN enrichment scenarios are discussed in a companion paper (Baldi et al., submitted)."
    },
    "0410/astro-ph0410471_arXiv.txt": {
        "abstract": "We develop a new analytic solution for the steady-state structure of a thin accretion disc under the influence of a magnetic field that is anchored to the central star. The solution takes a form similar to that of Shakura and Sunyaev and tends to their solution as the magnetic moment of the star tends to zero. As well as the Kramer's law case, we obtain a solution for a general opacity. The effects of varying the mass transfer rate, spin period and magnetic field of the star as well as the opacity model applied to the disc are explored for a range of objects. The solution depends on the position of the magnetic truncation radius. We propose a new approach for the identification of the truncation radius and present an analytic expression for its position. ",
        "introduction": "\\label{sec:int}  There is a well known one dimensional solution for the structure of a thin accretion disc in the steady state. This solution, known as the Shakura-Sunyaev disc solution \\citep{sha73}, consists of a series of seven equations. The full solution contains expressions for disc surface density, scaleheight, density, temperature, opacity and viscosity. These expressions depend upon the Shakura-Sunyaev alpha parameter, a dimensionless parametrization of viscosity. In addition to the alpha parameter the solution is expressed as a function of radius, mass transfer rate, stellar mass and stellar radius.  We wish to reformulate the Shakura-Sunyaev disc solution to incorporate the effect of a torque from a stellar magnetic field. A solution for the structure of an accretion disc in a magnetic field is desirable because accretion discs are frequently found about magnetic stars. Such magnetic fields can allow the transfer of angular momentum between the disc and star and can therefore significantly affect the structure of the disc as well as the spin evolution of the star \\citep[\\eg][]{bra98,arm96}. These effects can be important in young stars \\citep*[\\eg][]{kon90,arm99}, cataclysmic variables \\citep[\\eg][]{sch02} and X-ray binaries \\citep[\\eg][]{pri72,rom04}. Accretion discs can exhibit outbursts due to the thermal viscous disc instability, and the interoutburst time, known as the recurrence time, can be increased by several orders of magnitude by truncating the inner disc \\citep*[\\eg][]{mat03}. Accretion disc truncation by various mechanisms has also been invoked to explain the ultraviolet (UV) lag in dwarf novae \\citep[\\eg][]{mey94,kin97}. In this paper we treat magnetically enhanced accretion, but not the magnetic propeller effect.  A solution for the structure of a disc in a magnetic field is given by \\citet{bra98}. By neglecting some mass transfer terms, they are able to find a detailed solution for the structure of the inner disc where a specific opacity model is assumed. In their solution magnetic pressure is considered in addition to magnetic torque, which improves the accuracy of the model for the inner disc regions. However, the azimuthal magnetic field is parametrized with respect to the shear between the disc and the stellar field using essentially the same form as is used in this paper. Since the effect of mass transfer into the disc is partially neglected, the solution is most valid in the very inner parts of the disc, where the magnetic field dominates the disc structure completely. A complementary solution is presented in this paper which, though less accurate in the centre of the disc, is continuous for the entire disc and can be applied to a range of opacity models. Furthermore the solution derived below tends to the Shakura-Sunyaev solution in the case where the stellar magnetic field vanishes, and also at large radii.  The paper begins with a review of the necessary assumptions for the new solution. Then an approximate formulation for the magnetic torque is derived. This is applied, in conjunction with a new estimate of the magnetic truncation radius, to the derivation of a disc model. This model is generalised in order that it is compatible with any opacity which takes the form of a power law in density and temperature. We then explore how variables such as the mass transfer rate, magnetic field and stellar spin influence the disc structure. ",
        "conclusions": "\\label{sec:discussion} A new self consistent analytic solution has been developed for a thin accretion disc, under the influence of a central magnetic field. Since this model represents a steady-state disc it is not applicable to all accretion discs, and cannot be used for the magnetic propeller. The model was obtained in a similar manner to the non-magnetic Shakura-Sunyaev model, and indeed the result tends to that non-magnetic case when the magnetic parameter $\\beta$ tends to zero. The two models also converge at large radii. The analytic model has also been expanded to include the effects of a range of opacities in addition to Kramer's law. The solution has been applied in the cases of young stellar objects, cataclysmic variables and X-ray binaries.  The new model produces a solution in which the inner disc is frequently truncated further out than the stellar boundary layer. This is caused by magnetically induced accretion. The results depend upon how the position of the truncation radius $R_{\\rm{t}}$ varies as a function of the parameters of disc and star. It has been shown that the truncation radius is independent of the opacity prescription used. The magnetic interaction model is uncertain, so that neither the exact values of the magnetic parameters $\\beta$ and $\\gamma$ nor their dependence on the field strength of the star are so far well known. Further work is therefore required to better constrain the relationship between the surface field and the truncation radius. In addition, a comparison of the position of the truncation radius between the one dimensional solution and three dimensional techniques, such as smoothed particle hydrodynamics, could be used to ensure that the one dimensional inner boundary condition acts in the same manner as the three dimensional analogue.  The critical radius at which thermal viscous outbursts begin varies almost linearly with radius \\citep*{can88}. Therefore a truncated disc can store more mass before it goes into outburst. As a result magnetic fields can exert a strong influence on outburst cycles in accretion discs by forcing outbursts to begin further out in the disc. This is discussed in more detail by \\citet{mat03}. It would be instructive to perform full outburst simulations in which disc annuli are permitted to switch between the hot and cold state according to local triggers, and on a thermal time-scale. Similar work has been carried out to model dwarf nova outbursts in smoothed particle hydrodynamics \\citep*[\\eg][]{tru04}. The steady-state viscosity prescription should not be applied to a time evolution model because it becomes inaccurate at low densities. In addition a negative value of $\\nu$, which can occur for $R < R_{\\rm{t}}$, will cause equation (\\ref{sigmaevol}) to have sinusoidal solutions. More realistic simulations could be performed by relating viscosity to surface density or by using a more physically motivated viscosity such as that generated by the magneto-rotational instability \\citep[\\eg][]{bal91}.  Additional future work will involve the generalisation of the new solution to encompass torques from other sources such as planets embedded in the disc. This would have important implications for the study of planet formation and migration in discs. The effect of the magnetic torque on the central star's spin could also be investigated in detail. The long term effect of such spin evolution may be found to have an significant effect on the structure and behaviour of the disc."
    },
    "0410/astro-ph0410647_arXiv.txt": {
        "abstract": "We investigate the practice of assigning high spin temperatures to damped Lyman-$\\alpha$ absorption systems (DLAs) not detected in \\HI~21-cm absorption. In particular, \\citet{kc02} have attributed the mix of 21-cm detections and non-detections in low redshift ($z_{\\rm abs}\\leq2.04$) DLAs to a mix of spin temperatures, while the non-detections at high redshift were attributed to high spin temperatures. Below $z_{\\rm abs}=0.9$, where some of the DLA host galaxy morphologies are known, we find that 21-cm absorption is normally detected towards large radio sources when the absorber is known to be associated with a large intermediate (spiral) galaxy. Furthermore, at these redshifts, only one of the six 21-cm non-detections has an optical identification and these DLAs tend to lie along the sight-lines to the largest background radio continuum sources. For these and many of the high redshift DLAs occulting large radio continua, we therefore expect covering factors of less than the assumed/estimated value of unity. This would have the effect of introducing a range of spin temperatures considerably narrower than the current range of $\\Delta T_{\\rm s}\\gapp9000$~K, while still supporting the hypothesis that the high redshift DLA sample comprises a larger proportion of compact galaxies than the low redshift sample. ",
        "introduction": "\\label{sec:intro}  Despite their relative paucity, the highest column density absorption systems, damped Lyman-$\\alpha$ systems (where $N_{\\rm HI}\\ge2\\times10^{20}$ \\scm), are an important component of the high redshift ($z \\sim 2-6$) Universe since they account for most of the neutral gas available for star formation \\citep[e.g.][]{lwt+91}. The \\HI ~21-cm hyperfine transition can provide an alternative and complementary view of DLAs. Presuming that the 21-cm and Lyman-$\\alpha$ absorption arise in the same cloud complexes \\citep{dl90}\\footnote{Where the 21-cm and Lyman-$\\alpha$ absorption trace the cold and total neutral hydrogen, respectively.}, the column density [\\scm] of the absorbing gas in a homogeneous cloud is related to the velocity integrated optical depth, where $\\tau\\equiv-\\ln\\left(1-\\frac{\\sigma}{f\\,S}\\right)$, of the 21-cm line via \\citep{wb75}: \\begin{equation} N_{\\rm HI}=1.823\\times10^{18}\\,T_{\\rm spin}\\int\\!\\tau\\,dv\\,, \\label{enew} \\end{equation} where $T_{\\rm s}$ [K] is the spin temperature, $\\sigma$ is the depth of the line (or r.m.s. noise in the case of a non-detection) and $S$ and $f$ the flux density and covering factor of the background continuum source, respectively.  As summarised by \\citet{wol80}, 21-cm absorption line studies can face some difficulties compared with optical work: (1) in blind surveys the bandwidth of radio receiver/backend combinations only allows small redshift intervals to be searched for \\HI ~absorption; (2) terrestrial interference can be severe below 1.4\\,GHz; (3) if the spin temperature of the absorbing gas is high or if the background source size exceeds that of the absorption cloud(s), then the apparent optical depth can be considerably overestimated. Despite these technical difficulties, 21-cm studies of DLAs have provided information about the physical conditions of absorbers, including the kinematics, temperature and velocity distribution of the gas.  One unresolved question concerns the typical size and structure of DLAs, with models ranging from large, rapidly rotating proto-disks (e.g. \\citealt{pw97}) to small, merging sub-galactic systems (e.g. \\citealt*{hsr98}). Direct imaging by various groups \\citep[e.g.][]{lbbd97,rnt+03,cl03} of the low-$z$ DLA host-galaxies reveals a mix of spirals, dwarf and low surface brightness (LSB) galaxies. This `mixed morphology' picture of DLA host-galaxies appears to be reflected in the nearby Universe in a recent 21-cm $z=0$ emission study \\citep*{rws03}.  From a survey of 10 sources, in conjunction with the available literature, \\citet[][ hereafter \\citetalias{kc02}]{kc02} find both detections and non-detections of 21-cm absorption in DLAs at $z<2$, whereas the high-$z$ results are almost exclusively non-detections (Fig. \\ref{kanekar}, top). \\begin{figure} \\vspace{8.8cm} \\special{psfile=ME1073rv-fig1.ps hoffset=-20 voffset=290 hscale=75 vscale=75 angle=270} \\caption{Top: Spin temperature against absorption redshift as given in \\protect\\citetalias{kc02}. The solid symbols represent the 21-cm detections and the shapes represent the type of galaxy with which the DLA is associated: circle--unknown type, star--spiral, square--dwarf, triangle--LSB (obtained from \\protect\\citetalias{kc02} and references therein). This key is also applied to Figs. \\ref{m2} to \\ref{Toverf}. The arrows show the spin temperature lower limits and all of these bar one (0454+039 at $z_{\\rm abs}=0.8596$) have unknown host identifications. Bottom: The covering factors applied to derive the spin temperatures. The shapes are coded as above with the solid and open symbols representing the covering factors applied for the 21-cm detections and non-detections, respectively.} \\label{kanekar} \\end{figure} Additionally, many of the $z<1$ DLAs have host-galaxy identifications from optical imaging, given in Table \\ref{t2}. By combining the 21-cm optical depth with the total \\HI\\ column density from optical spectroscopy covering the damped Lyman-$\\alpha$ line, \\citetalias{kc02} find low spin temperatures ($T_{\\rm s}\\approx 200$ K), which are typical of the Milky Way and local spirals \\citep{dl90}, in all cases where the absorber was identified as a spiral (Fig. \\ref{kanekar}). Furthermore, all DLAs at $z < 1$ with temperatures $T_{\\rm s} \\geq 1000$~K were found to be associated with dwarf or LSBs and \\citetalias{kc02} interpreted the non-detections at high-$z$ ($>2$) as DLAs with high $T_{\\rm s}$. They therefore advocate a picture where the DLA host-galaxy population is dominated by the warmer dwarfs and LSBs at high-$z$ and evolves to include a higher proportion of spirals at low-$z$, thus satisfying one expectation of hierarchical galaxy formation scenarios. However, in many of these studies, particularly at high redshift, the covering factor, $f$, of the background continuum is often assumed to be unity (Fig. \\ref{kanekar}, bottom -- see also Table \\ref{t2}.). Since $f$ plays an equal r\\^{o}le as $T_{\\rm s}$ in relating the column density to the observed optical depth profile, it is important to consider the robustness of this assumption, which is the aim of this article. ",
        "conclusions": "Upon a review of the literature, we suggest that less-than-complete covering factors may be at least partly responsible for the non-detection of 21-cm absorption in many DLAs.  Although it is as important as the spin temperature in determining the column density of the absorber, the covering factor is at best estimated from insufficient data, and, more often than not, simply assumed to be unity. This assumption is prevalent in the high redshift cases, where high spin temperatures are claimed to be dominant, despite their being no DLA host identifications at $z_{\\rm abs}\\gapp0.9$.  Although all detections of 21-cm absorption occur at $z_{\\rm abs}\\leq2.04$, nearly half of those searched at these redshifts remain undetected, thus giving an apparent mix of both low and high spin temperatures for this sample. However, of these non-detections, $\\approx70$ per cent of the absorbers occult large background radio sources ($\\gapp1''$), thus increasing the likelihood of $f<1$. Furthermore, 21-cm absorption tends to be detected: \\begin{enumerate} \\item Towards large background sources ($>1''$) only when the DLA is known to be associated with a large galaxy. \\item Towards compact sources ($<0.1''$) when the DLA is associated with an LSB or dwarf galaxy. \\end{enumerate} Naturally, we would also expect detections where the absorber is associated with a large spiral occulting a compact background continuum, thereby providing large coverage. However, this is not observed. Of the 6 detections towards small background sources, 3 of the hosts are unidentified, thus not ruling out the possibility of an approximately equal mix of large and compact DLA hosts.  From the observed bolometric background, {Wolfe}, {Gawiser} \\& {Prochaska} (2003) propose that the 21-cm absorbing CNM is a major component of DLAs. They suggest that this is comprised of several lower column density clouds which may exhibit a large total column density but low covering factor when aligned along the line-of-sight or, conversely, low individual cloud column densities and a large covering factor, where overlap along the line-of-sight is minimal. Therefore, when attributing the non-detection of 21-cm absorption in DLAs to the effects of the covering factor (or spin temperature), we should bear in mind that the total \\HI ~column density may not provide a reliable estimate of the column density of the CNM along the line-of-sight to the background continuum. Further complications arise from the fact that insufficient detail is known about the flux contributions from the various components of the background continuum illuminating the DLA.  The high resolution maps are at different frequencies than that of redshifted 21-cm and, if the components are extended, spectral indices may vary across them, making estimates of the effective continuum size at the appropriate frequency difficult. It is clear that high resolution observations at the appropriate frequencies are required in order to fully address this issue. Combining these with the absorption profiles at various locations over the background continuum (as in \\citealt{lbs00}) will give far better estimates of the covering factor.   Despite the above caveats, our analysis suggests that the covering factor may play an important r\\^{o}le, thereby significantly lowering the spin temperature estimates of the DLAs for which the covering factors have been assumed (i.e. most of the non-detections) and perhaps it would be more prudent to use $T_{\\rm spin}/f$ versus absorption redshift (Fig. \\ref{Toverf}) rather than just $T_{\\rm spin}$, for an \\begin{figure} \\vspace{6.2cm} \\special{psfile=ME1073rv-fig5.ps  hoffset=-20 voffset=255 hscale=75 vscale=75 angle=270} \\caption{Spin temperature/covering factor ratio against absorption redshift for the optically thin ($\\tau\\lapp0.3$) 21-cm searches. Again the FWHM of the non-detections is assumed to be 20 \\kms (Section \\ref{20kms}).} \\label{Toverf} \\end{figure} unknown value of $f$. The figure does in fact show that the values of $T_{\\rm spin}/f$ span similar ranges for both the detections and non-detections with the spirals (stars) being more concentrated towards the low $T_{\\rm spin}/f$ (and $z_{\\rm abs}$) regime\\footnote{Note however that the spread along the ordinate is considerably wider than in Fig. \\ref{kanekar} (top): Over 1.2 dex in $T_{\\rm spin}/f$ cf. 0.3 in $T_{\\rm spin}$. Arguably this is mainly influenced by 1229--021 at [0.39, 3.32] but also applies when we, as \\citetalias{kc02}, include the associated system 1413+135 at [0.25, 3.28].} (although, as stated previously, the majority of host identifications are unknown). This is apparently consistent with the interpretation of low spin temperatures in these objects (\\citetalias{kc02}), but also with high coverage by spirals, a possibility strongly suggested by the work presented here.  The lowering of covering factor estimates would introduce more uniformity to the spin temperatures of DLAs, while still supporting the hypothesis that large galaxies dominate the low redshift population \\citep{bmce00} with dwarf galaxies constituting more of the high redshift population \\citep{lf03}. That is, unlike \\citetalias{kc02} who suggest that this is evident through increased spin temperatures at high redshift, we hypothesize that the absorption cross-sections of compact galaxies at these redshifts simply lack the extent to effectively cover the background radio continua."
    },
    "0410/astro-ph0410121_arXiv.txt": {
        "abstract": "We have used the EFAR sample of galaxies to investigate the light distributions of early-type galaxies. We decompose the 2D light distribution of the galaxies in a flattened spheroidal component with a \\citet{Ser68} radial light profile and an inclined disc component with an exponential light profile. We show that if we assume that all galaxies can have a spheroid {\\it and} disc component, that then the brightest, bulge dominated elliptical galaxies have a fairly broad distribution in the S\\'ersic profile shape parameter \\nb, with a median of about 3.7 and a sigma of $\\sim$0.9. Other galaxies have smaller \\nb\\ values. This means that spheroids are in general less concentrated than the de Vaucouleurs \\Rq-law profile, which has \\nb=4.  While the result of our light decomposition is robust, we cannot prove without kinematic information that these components are spheroids and discs, in the usual sense of pressure- and rotation-supported stellar systems. However, we show that the distribution of disc inclination angles is consistent with random orientation if we take our selection effects into account. If we assume that the detected spheroids and discs are indeed separate components, we can draw the following",
        "introduction": "Studying the light distribution of galaxies in order to determine and compare their structural properties has a long tradition (e.g., \\citealt{deV48,Fis64,Fre70}). Traditionally, the light distributions of elliptical galaxies have been fitted by the \\Rq-law profile of \\citet{deV48}, while spiral galaxies have been fitted with exponential light profiles for their discs and often, due to their similarities to early-type galaxies, \\Rq-law bulges (e.g., \\citealt{deV59,Kor77}). In the recent decade we have seen some shift from this paradigm. A generalization of the \\Rq\\ and exponential light profile laws described first by \\citep{Ser68}, resulting in an \\Rn\\ law, has more often been fitted to early-type galaxies to account for observed deviations from the \\Rq profiles (e.g., \\citealt{Cao93, GraCol97}). Based on this development and based on theoretical predictions for some forms of bulge formation (e.g., \\citealt{Com90,PfeNor90}), more general forms of light profiles have been fitted to the bulge light of spiral galaxies as well \\citep{And95,deJ96II,Mac03}. With the realization that many early-type galaxies have disc-like components (e.g., \\citealt{Ben89,RixWhi90,JorFra94}), bulge/disc decompositions of elliptical galaxies have been performed, but in general with \\Rq\\ bulges and exponential discs (e.g., \\citealt{EfarIV}) except for a few notable exceptions on small samples of galaxies \\citep{DOn01, Bal04, Gut04}. Here we perform 2D bulge/disc decompositions on the large EFAR sample of early-type galaxies (\\citealt{EfarI}, hereafter Paper I) using an \\Rn\\ spheroid and exponential disc light profiles, determine relations between their structural parameters and their frequency of occurrence, and compare those to the properties of disc dominated galaxies.    In an earlier paper in this series (\\citealt{EfarIII}, Paper III hereafter) we presented the luminosity profiles for the 776 galaxies observed in the EFAR project to which we fitted seeing convolved \\Rq-law bulge plus exponential disc models.  In Paper III, 31\\% of the sample proved to be spiral or barred spiral galaxies and 69\\% of the sample were classified as early-type galaxies.  Of those early-types 18\\% were well fit by an $R^{1/4}$ alone, the great majority (70\\%) were very much better fit by the combination of an exponential disc plus \\Rq\\ bulge than by a bulge alone.  The remaining 12\\% of the early-type galaxies were classified as cD, without reference to their profile type, as they were the brightest in their cluster and had half-total-light radii larger than 10\\,kpc.     In Saglia et al. (\\citeyear{EfarIV}, hereafter Paper IV) we showed that the combination of an \\Rq-law spheroid plus an exponential disc is degenerate to an \\Rn-law spheroid, unless one has data for a very large range in surface brightness and radius.  Given that only 12\\% of the EFAR galaxies were best fitted with an \\Rq\\ law alone, and that the remainder of the galaxies were better fitted with an \\Rq\\ bulge and an exponential disc, one might wonder about the validity of the \\Rq\\ law in general. To address this issue, we have refitted all EFAR galaxies again with an \\Rn-law spheroid and an exponential disc.  If early-type galaxies are dominated by \\Rn\\ spheroids, most fits should converge to largely spheroid dominated systems with small discs.  If the \\Rq-law spheroids plus disc picture is correct, most fits should converge towards spheroids with S\\'ersic parameter $n$ close to 4.  Obviously, the truth could be somewhere in between with spheroids with S\\'ersic parameter $n$ covering a large range and with a large range in bulge-to-disc ratios.  We have to realize that even when we find that most early-type systems have light components that are well described by exponential light profiles, this in itself will never prove the existence of exponential discs. First of all, there is no reason why the exponential discs observed in most spiral galaxies should be present in a similar way in bulge dominated systems, especially realizing that exponential discs in spiral galaxies are still poorly understood. Secondly, even if light is present in exponential distribution, it might not be in a disc configuration flattened by rotation as discs of spiral galaxies are normally understood. One can probably make many stable configurations in bulge-like potentials that have small exponential like stellar distributions on top of them, without the need of any net rotation. The best confirmation of exponential light distributions in early type galaxies being rotationally flattened discs as in spiral galaxies has to come from measurements of their kinematics.  Even when kinematics is available, it still might be difficult to reach an unambiguous result. \\citet{ScoBen95} made dynamical models for a small sample of early-type galaxies based on their photometric bulge/disc decompositions and showed that the observed asymmetries in the long-slit stellar absorption line profiles were consistent with the dynamical model predictions.  However, with their data and modelling technique they could not prove this was an unique kinematic interpretation. Better kinematic constraints on discy components can be obtained with full 2D velocity profile measurements, such as obtained with the SAURON project \\citep{Fal04,Ems04}. However, so far these studies have been limited to the high surface brightness central regions of galaxies.  In recent years a number of studies were performed where the galaxy samples were large enough and the selection bias understood well enough that bivariate space density distributions of structural parameters could be calculated. \\citet{Cro01} presented the bivariate (luminosity, surface brightness distribution) for the 2dFGRS \\citep{Col01}. \\citet{Bla03} showed many different bivariate space density distributions derived from the Sloan Digital Sky Survey (SDSS; \\citealt{Yor00}), and presented some distributions in bins of colour and concentration index to split their sample in different galaxy types. \\citet{Kau03} derived mass-to-light ratios for the SDSS galaxies and presented stellar mass bivariate distributions. \\citet{She03} also investigated the bivariate distribution of structural parameters of different subsets of SDSS galaxies. However, even though calculated, they did not include the absolute space densities in their analysis.  While based on large sample to provide good statistics, in none of these studies bulge/disc decompositions were performed to derive space density distributions for the individual components. Such an analysis was performed on the discs of a sample of about 1000 Sb-Sdm galaxies by \\citet{deJLac00}. However, no such analysis on early-type galaxies has appeared so far.  A fundamental question for galaxy formation is the origin of the Hubble sequence. Why do some galaxies have discs and others not?  If in fact almost all galaxies have a disc component but of varying mass, luminosity, and size, the question of the origin of the Hubble sequence takes on a different perspective; in semi-analytic models of galaxy formation early-type galaxies go through phases where they have substantial discs which are disrupted and re-built during the cycle of hierarchical merging.  Determining galaxy disc and bulge parameters and their bivariate distributions will help us to test and constrain such hierarchical semi-analytic models of galaxy formation.  In this paper we will use a 2D fitting algorithm to measure the bulge and disc properties in the EFAR sample of early-type galaxies.  We estimate the frequency, luminosity and size of discs and bulges and explore the properties of the discs in the context of S0 and spiral galaxy discs.  In the next section we describe our method of 2D model fitting and parameter error estimation and compare our new results with the results of Paper III.  In section 3 we characterize the properties of the discs and their relation to the host bulge component.  Our conclusions and the implications of this work for models of galaxy evolution are summarized in section 4.  We use a cosmology of $\\Omega_m$=0.3, $\\Omega_\\Lambda$=0.7, and $H_0$=70 km s$^{-1}$ Mpc$^{-1}$ throughout. ",
        "conclusions": "  \\begin{enumerate}  \\item The scale sizes of the bulge and disc components are correlated. This may be partly due to the exclusion of spiral galaxies from the sample, excluding objects with small bulges and large discs. However, it should be noted that spiral galaxies show a similar trend \\citep{deJ96III,Cou96,Mac03}.  \\item The bulge-to-disc \\bt\\ ratios, the S\\'ersic \\nb\\ bulge shape parameters, and the bulge effective radii show positive correlations with each other. We have no model that tells us which is cause and what is effect.  \\item The median \\nb\\ value is 3.24 for all galaxies in our sample, is 3.66 for all non-spiral, bulge dominated galaxies in our sample, and is 4 for non-spiral galaxies with \\bt$>$0.7. Given that most spiral galaxies have bulge \\nb\\ values much lower than these values, this means that the standard \\Rq\\ de Vaucouleurs (i.e.\\ \\nb=4) fitted to spheroids may not be the most appropriate choice when doing bulge/disc decompositions on a random set of galaxies. In cases where one does not have the signal-to-noise or the spatial resolution to perform full S\\'ersic bulge plus exponential disc fit, it will be more correct to use fixed value of \\nb$\\sim$3.5 for the brightest galaxies. Lower values should be used for less luminous galaxies.  \\item The (luminosity, scale size) bivariate space density distribution of bright elliptical galaxies is well described by the analytic parametrization presented by \\citet{deJLac00}.  \\item Comparing the total galaxy (luminosity, scale size) bivariate distributions of EFAR non-spiral galaxies with SDSS disc-like galaxies shows that the brightest early-type galaxies could in principle be formed by merging a few large, current day spiral galaxies.  Low luminosity early-type galaxies have much smaller scale sizes than spiral galaxies of the same luminosity, and hence they can surely not be created from simple merging of two current day spiral galaxies.  \\item When comparing the EFAR bulge and disc (luminosity, scale size) bivariate distributions to those of the SDSS disc-like galaxies, we showed that the bulges of early-type galaxies are typically a factor 1.5 to 2.5 smaller and the discs a factor 2 larger than current day disc galaxies at a given luminosity.  \\end{enumerate}  We speculate that the spheroidal components are the result of merging of similar sized small proto-galaxies galaxies at high redshifts, while the disc-like component may be the result of smaller galaxies falling in later along the filaments of the cosmic web. While attractive from the galaxy structure point of view, this model requires careful tuning of the formation process of the proto-galaxies at high redshift in order to reproduce the mass-metallicity relation and the $\\alpha$-element over-abundace observed at lower redshifts."
    },
    "0410/astro-ph0410317_arXiv.txt": {
        "abstract": "We discuss three different observational diagnostics related to disks around hot stars: absorption line determinations of rotational velocities of Be stars; polarization diagnostics of circumstellar disks; and X-ray line diagnostics of one specific magnetized hot star, $\\theta^1$ Ori C.  Some common themes that emerge from these studies include (a) the benefits of having a specific physical model as a framework for interpreting diagnostic data; (b) the importance of combining several different types of observational diagnostics of the same objects; and (c) that while there is often the need to reinterpret traditional diagnostics in light of new theoretical advances, there are many new and powerful diagnostics that are, or will soon be, available for the study of disks around hot stars. ",
        "introduction": "Disks are a spatial and dynamic phenomenon, which would, ideally, be studied by direct imaging and {\\it in situ} measurements to learn about their spatial extent and physical properties (temperature, density, ionization structure, scale height, velocity fields).  Needless to say, direct imaging and {\\it in situ} measurements will not be practical any time in the near future, although indirect imaging, via interferometry, has begun to be a reality. Thus, we must use the traditional tools of stellar astrophysics: photometry, spectroscopy, and polarimetry in order to learn about the important physical properties of the disks around hot stars.  While employing these diagnostics, it is important to bear in mind several points, most if not all of which flow from our emphasis on deriving information about the specific {\\it physical} properties of these disks:  \\begin{itemize} \\item It is important to not put the cart before the horse and make observations simply because new instruments become available.  One must first determine what information is needed, preferably to test competing hypotheses, and then find the right diagnostic or set of diagnostics that can address the specific question.  \\item Interpreting the results of diagnostics is more straightforward when there is an underlying physical model, but it is also important to keep in mind that the physical model might be wrong or at least too simplistic.  Theorists are most useful when their theories can make clear, quantitative predictions that can be tested by observations.  \\item It is important to look at the sensitivity of diagnostics to relevant parameters and to critically evaluate how much diagnostic discrimination is possible from a given observation.  \\end{itemize}  There are important new diagnostics coming on line in the next few years, and our community should plan on exploiting them effectively.  These include sensitive polarization measurements of Zeeman splitting and new interferometers, among others.  But we should also bear in mind that some older and more traditional diagnostics can still be used profitably, and that the interpretation of old measurements can also change with new insights from theorists.  In the rest of this paper, we report on three different pieces of work involving diagnostics of hot star disks.  In \\S 2, Richard Townsend evaluates the methods for inferring stellar rotational velocities from observed line profiles.  In \\S 3, Karen Bjorkman discusses spectropolarimetric diagnostics of hot star disks.  In \\S 4, Marc Gagne combines X-ray spectroscopy with several other diagnostics to infer the physical properties of, and causal mechanisms at work in, the young magnetized O star, $\\theta^1$ Ori C.  We offer some concluding remarks in \\S 5. ",
        "conclusions": "The field of diagnostics is an extraordinarily broad and diverse one, and in this paper we have touched on only three rather specific diagnostics and their associated applications.  However, while there is both a certain amount of common ground as well as some interesting differences among these three examples, several underlying lessons emerge. First of all, researchers must be aware of their biases and always be leery of the assumptions that underly the models that they rely on to interpret data.  Employing models to interpret data is of crucial importance, though, as determining physical parameters of interest is the ultimate motivation for employing diagnostics in the first place.  The section on absorption line profiles and Be star rotational velocities shows that even traditional diagnostics must be tied to models in order to provide meaningful physical interpretations, and that the assumptions behind models should be reevaluated periodically as theories become more sophisticated.  Also, observers must not lose sight of which specific properties of the observational data are sensitive to the physical parameters of interest.  The section on polarization diagnostics demonstrates how the electromagnetic radiation we measure contains a huge amount of information that can be extracted if the diagnostics are interpreted in light of physical models.  And, like so many diagnostic techniques, there are flies in the ointment (in this case, the removal of interstellar polarization) that can be difficult to deal with, but which must be accounted for if meaningful conclusions are going to be drawn from the data.  This case is especially satisfying because it is in fact one of the inherently interesting properties of the class of objects under study -- namely the rapid dispersal and then formation of Be disks -- that enables us to make an accurate assessment of ISP contamination. Finally, the section on $\\theta^1$ Ori C presents a nearly unique case that nonetheless illustrates some of the basic guidelines for employing observational diagnostics.  In this final section, we see again how important it is to have an underlying physical model for the interpretation of diagnostics.  We also see how combining data at several wavelengths is vital.  And finally we see how important quantitative measurements of specific physical properties (e.g.\\ X-ray line widths and line ratios) are for making meaningful progress in understanding astrophysical systems of interest."
    },
    "0410/astro-ph0410067_arXiv.txt": {
        "abstract": "An essential feature of substellar dwarfs compared with the Sun and stars is the formation of dust in their photospheres (not in the outer envelope). It appears that the observed data could be understood if the dust exists in a form of thin cloud deep in the photosphere rather than in the cooler surface region. Recent observations also show that the dust column densities in the observable photosphere are quite different for the same effective temperature, gravity, and metallicity, but the reason for such a sporadic variation is unknown. Moreover, the effect of dust cloud is difficult to discriminate from those by other basic parameters such as the effective temperature which also has significant effect on the dust column density. For this reason, the spectra of dusty dwarfs were in fact mis-interpreted by ourselves and will be reanalyzed in this contribution. Also, even the  spectral classification is not free from such a difficulty, as is evidenced by an odd ``brightening'' of $M_{\\rm bol}$ plotted against the L and T types. ",
        "introduction": "It is only a decade ago that a genuine brown dwarf was discovered (\\cite{nak95}) and our experience in interpreting the spectra of brown dwarfs is still meager compared with the longer experience in the Sun and stars. A new feature in ultracool dwarfs such as brown dwarfs is the presence of dust in their photospheres, the possible importance of which has been recognized at an early time (\\cite{tsu96}). But how to treat the dust formed in the photospheric environment was not known. If a simple thermodynamical equilibrium is assumed, dust certainly forms but how to sustain the dust in the photosphere was unknown and  dust also tends to be over-produced compared to the known observations in general. Actually, dust in the observable photosphere should be controlled by such processes as nucleation, growth, segregation, precipitation, evaporation etc., and we proposed a simple model referred to as the unified cloudy model (UCM) to take into account these processes semi-empirically (\\cite{tsu02}).  In the UCM, we assumed that dust forms at the condensation temperature $ T_{\\rm cond}$,  but dust soon grows to be too large at a slightly lower temperature which we referred to as the critical temperature $ T_{\\rm cr}$ and segregates from the gaseous mixtures. Thus dust forms a homogeneous cloud in the region where $  T_{\\rm cr} \\la T \\la T_{\\rm cond}$. In this model, $ T_{\\rm cond}$ is essentially determined by the thermodynamical data but $ T_{\\rm cr} $ is not predictable at present. At first, we assumed that $ T_{\\rm cr} $ remains the same throughout L and T dwarfs for simplicity. Recent observations, however, revealed that $ T_{\\rm cr}$, which is a measure of the thickness of the cloud (a larger deviation of $ T_{\\rm cr} $ from $ T_{\\rm cond}$ implies a thicker cloud), shows a sporadic variation even at the fixed effective temperature (Sect.2). Our UCM is already formulated to allow for the change of $ T_{\\rm cr}$  and its application to the new situation is straightforward. Compared with our previous interpretation of the spectra of L and T dwarfs based on a constant value of $ T_{\\rm cr}$, our revised analysis allowing for the variation of $ T_{\\rm cr}$ in UCMs results in a drastic change in our understanding of L and T dwarfs. ",
        "conclusions": "\\label{sec:tit}  In the interpretation of the spectra of dust-contaminated object such as brown dwarfs, a parameter that specifies the thickness of the dust cloud, $T_{\\rm cr}$ in our formulation or $f_{\\rm sed}$ in that of \\cite*{mar02}, plays an important role. At present, it seems to be difficult to predict the value of $T_{\\rm cr}$ based on basic physics, especially because it is sporadic rather than related to the other basic parameters (Sect.2). The variation of $T_{\\rm cr}$ is especially large at $T_{\\rm eff} \\approx 1400$\\,K where the second convective zone appears in addition to the one deep in the photosphere (\\cite{tsu02}). It may be possible that $T_{\\rm cr}$ is related to the convective activities, but details are yet to be explored. Thus, in addition to the four parameters generally required to specify the stellar photosphere, i.e.,  chemical composition, $T_{\\rm eff}$, log\\,$g$,  and micro-turbulent velocity, fifth parameter $T_{\\rm cr}$ is needed for the characterization  of dusty dwarfs.  It is to be noted that the turbulent velocity is still determined empirically since its discovery more than half a century ago (\\cite{str34}) and not yet predictable based on basic physics for individual objects. At present, we must leave $T_{\\rm cr}$ as an empirical parameter something like the turbulent velocity, but there are difficulties inherent to dust. For example, dust, unlike atoms and molecules, shows almost no clear spectral signature and it is difficult to estimate the dust column density directly. While dust plays significant role in defining the spectral characteristics, spectral classification had to be done on the spectral features originating from atoms and molecules (e.g. \\cite{kir99}, \\cite{mar99}, \\cite{bur02}, \\cite{geb02}), and such a difficulty should be fatal in the spectral classification of dusty objects.  Despite some formidable problems, recent progress in observation of such faint objects as brown dwarfs, not only in spectroscopy but also in photometry and astrometry, is quite marvelous, and a more consistent interpretation can be achieved by considering all these data collectively."
    },
    "0410/astro-ph0410251_arXiv.txt": {
        "abstract": "We present the first spatially resolved spectrum of scattered light from the \\twhya\\ protoplanetary disk. This nearly face-on disk is optically thick, surrounding a classical T~Tauri star in the nearby 10~Myr old \\tw\\ association. The spectrum was taken with the \\hst-STIS CCD, providing resolution $R \\sim 360$ over the wavelength range 5250 -- 10300~\\AA. Spatially resolved spectroscopy of circumstellar disks is difficult due to the high contrast ratio between the bright star and faint disk. Our novel observations provide optical spectra of scattered light from the disk between 40~AU and 155~AU from the star. The scattered light has the same color as the star (gray scattering) at all radii, except the innermost region. This likely indicates that the scattering dust grains are larger than about 1~\\um\\ all the way out to large radii. From the spectroscopic data, we also obtained radial profiles of the integrated disk brightness at two position angles, over almost the same region as previously observed in \\hst-WFPC2 and NICMOS coronagraphic images (35~AU to 173~AU from the star). The profiles have the same shape as the earlier ones, but show a small azimuthal asymmetry in the disk not previously noted. Our STIS broad-band coronagraphic images of \\tw\\ confirm the reality of this asymmetry, and show that the disk surface brightness interior to 140~AU has a sinusoidal dependence on azimuthal angle. The maximum brightness occurs at a position angle of $233.^{\\circ}6 \\pm 5.^{\\circ}7$ East of North. This might be caused by the combination of forward-scattering and an increase in inclination in the inner region of the disk, suggesting that the \\tw\\ disk has a warp like that seen in the \\bpic\\ debris disk. ",
        "introduction": "\\label{sec:intro}  \\objectname[HIP 53911]{TW Hydrae} is a classical T~Tauri star, indicating that gas and dust from a circumstellar (CS) disk are accreting onto the young star (spectral type K7 Ve, distance $= 56.4^{+8.1}_{-6.2}$~pc). This disk has now been imaged at visible, near-IR, and millimeter wavelengths (\\citet{Krist:2000}, \\citet{Weinberger:2002}, \\citet{Wilner:2000}). The disk inclination is $7^\\circ$ from face-on \\citep{Qi:2004} and no azimuthal asymmetries have been previously reported. A large fraction of the stellar light is reprocessed by the CS dust to far-IR wavelengths ($L_{\\mathrm{IR}}/L_{\\star} = 0.25$), indicating that the disk is optically thick \\citep{Adams:1987}. These characteristics suggest that the \\tw\\ disk is relatively unevolved for its age ($\\sim 10$~Myr; \\citet{Webb:1999}).  However, while $\\sim$1~\\um\\ grains in the surface layers of the disk can produce the mid-IR spectrum of the disk, very large (mm to cm sized) grains are needed to explain the mm-wavelength spectral energy distribution \\citep{Weinberger:2002}. The relative lack of near-IR excess flux from the disk indicates that the inner 4~AU has been partially cleared of material, possibly by formation of a giant planet \\citep{Calvet:2002}. The presence of crystalline silicates in the disk, which are seen in Solar System meteorites and comets but not in the interstellar medium, also suggests that the growth of planetary material has begun \\citep{Weinberger:2002, Uchida:2004}.  Previous broad-band photometry indicated that the disk scattering was wavelength-independent and that the grain albedo might be large \\citep{Weinberger:2002}. This suggests that the disk is composed mostly of icy grains larger than $\\sim$1~\\um. However, this conclusion is tentative, since it was based on broad-band photometry of the scattered light, rather than actual spectra. Also, this photometry provided no information about changes in the scattering with location in the disk.  As part of our efforts to understand how the planetary formation process depends on location in the disk, we obtained the first spatially resolved spectrum of scattered light from a protoplanetary disk. Our goals were to study the structure of the \\tw\\ dust disk, and the size and composition of dust grains as a function of radius from the central star. In addition, we present broad-band coronagraphic imaging of \\tw\\ (Section~\\ref{sec:coron}). ",
        "conclusions": "\\label{sec:conc}  These observations demonstrate the usefulness of high-contrast coronagraphic spectroscopy of scattered light from CS disks. We here summarize our principal observational results from this study of the \\tw\\ protoplanetary disk.  \\begin{enumerate} \\item Interior to 140~AU, the disk surface brightness shows a sinusoidal dependence on azimuthal angle. The maximum brightness is at PA = $233.^{\\circ}6 \\pm 5.^{\\circ}7$. The origin of this feature requires further study, but we suggest it is caused by an increase in inclination in the inner part of the disk, i.e.\\ a warp. \\item The disk spectra definitively show that the scattering is wavelength-independent in the optical between 55~AU to 155~AU from the star. We argue this indicates that the scattering grains are larger than about 1~\\um\\ out to large radii. \\item The color of the scattered light from the innermost portion of the disk (34.5 -- 46.0~AU) appears bluer than the outer disk. This result is tentative, due to possible systematic errors in PSF subtraction close to the fiducial bar. However, taken at face value, the change in color might indicate a decrease in grain size at small radii. \\end{enumerate}"
    },
    "0410/astro-ph0410298_arXiv.txt": {
        "abstract": "We analyze a 3D Smoothed Particle Hydrodynamics simulation of an evolving and later collapsing pre-stellar core. Using a 3D Continuum Radiative Transfer program, we generate images at 7, 15, 175 $\\mu$m, and 1.3 mm for different evolutionary times and viewing angles. We discuss the observability of the properties of pre-stellar cores for the different wavelengths. For examples of non-symmetric fragments, it is shown that, misleadingly, the density profiles derived from a 1D analysis of the corresponding images are consistent with 1D core evolution models. We conclude that 1D modeling based on column density interpretation of images does not produce reliable structural information and that multi-dimensional modeling is required. ",
        "introduction": "Molecular cloud cores are thought to be the direct progenitors of stars. However, their initial properties and early evolution are still poorly understood. Current observations therefore aim to find and study these cores in more detail. Density structure, velocity field, and temperature distribution of the gas and dust are key parameters for the physical interpretation of the long lifetimes of these objects. The continuum radiation spectra of deeply embedded sources contain only ambiguous information about the density and temperature distribution of the dust. Column densities can be inferred from the analysis of continuum images in the mm-range \\citep[e.g.][]{War94} and the mid-infrared (MIR) \\citep[e.g.][]{Bac00}. The emission of molecules within the core contains information about the structure, velocity field \\citep[][and references therein]{Taf04}, and the turbulence \\citep[][and references therein]{Oss01}. The currently discussed supporting mechanisms against gravitational collapse are magnetic fields and turbulence. There are a number of simulations treating the full 3D structure of collapsing pre-stellar cores \\citep[e.g.][]{BBB03, BKV03, BurBod93, BurBod00, KruFisKle03, KleHeiMac00, HeiMacKle01}. In some of these papers, the resulting column density along a line of sight was compared with observationally obtained column densities. The question which structures of the distribution can be seen at which wavelength, however, can only be answered by producing images of the core for given dust properties. In turn, the majority of models that have been applied to observational data are mostly based on spherical symmetry \\citep[e.g.][]{And96}. They rely on spherically symmetric models of isolated star formation, which describe core formation, gravitational collapse, and proto-stellar accretion. Commonly, averaged radial density profiles are derived and compared with power-law density distributions \\citep[e.g.][]{And96}, or Bonnor-Ebert spheres \\citep[e.g.][]{Alv04}. The major difficulties using 3D models in observations are i) the information loss due to the projection effect, ii) the complex and unique structure of each individual pre-stellar cores, and iii) the numerical effort of multi-dimensional radiative transfer.  In this letter, we investigate how an evolving cloud core simulated by a 3D Smoothed Particle Hydrodynamics (SPH) code would appear at different wavelengths, which structures are visible, and which density distributions would be inferred using common 1D models. The results from the SPH simulation are described in Sect.~2, along with a discussion of underlying assumptions. We present the images from the 3D Continuum Radiative Transfer (CRT) modeling and discuss the observability of different structures and physical effects. The comparison to density structures obtained from applying 1D models to the images is given in Sect.~3, and the findings are summarized and discussed in Sect.~4. ",
        "conclusions": "We have presented 3D simulations assuming that initially low mass condensations pass through a stage of turbulence dominated condensation where they accumulate mass and merge together to form extended pre-stellar core like objects. The typical density structures in the cores are non-spherical throughout their evolution. The asymmetry is driven by the turbulent motion and causes complex structures from the very beginning. This complexity is partially seen in images that have been calculated from the densities obtained in the cloud core simulation. However, projection effects can lead to a severe misinterpretation of images. We showed that a 1D analysis of the vicinity of the density maxima would suggest a density profiles in agreement with 1D core collapse models. The underlying density structure, however, is intrinsically three-dimensional and deviates strongly from the obtained 1D model distribution.  As the column density also enters the optical thin emission in the mm-range (aside from the Planck function), we expect the same projection ambiguities to occur when interpretating mm-maps of dense molecular cloud regions. This aspect will be discussed in a forth-coming paper.  We conclude that 1D modeling based on column density interpretation of images does not produce reliable structural information. For flattened structures appearing in later stages of the core evolution, a 2D modeling might be applicable, but for the general case, multi-dimensional continuum and line radiative transfer modeling is required to derive consistent density and temperature distributions of the gas and dust in pre-stellar cores."
    },
    "0410/astro-ph0410720_arXiv.txt": {
        "abstract": "\\vspace{-0.4cm} \\baselineskip 12pt \\parskip 0pt  In this paper we use the predicted spatial development of helical structures along an expanding jet to model observed structures and motions in the 3C\\,120 jet.  New results of VLBI imaging of the parsec-scale radio jet in 3C\\,120 at 5 GHz are examined along with older long term monitoring results at 5 GHz and older results obtained at 22 and 43 GHz.  The high frequency observations provide detailed information on motions and structure from 0.5 to 10~mas from the core and the lower frequency observations from 1 to 30~mas from the core. Proper motions of helical components associated with the pattern and of other components that move through the pattern provide estimates of flow and helical pattern speeds.  Theoretical modeling of the motion and appearance of the helical pattern allows determination of sound speeds as a function of the jet viewing angle.  The jet sound speed declines although probably not as fast as adiabatically.  At a 12\\arcdeg\\ viewing angle the most likely scenario involves a decline in jet sound speed from $c/3 < a_j < c/\\sqrt 3$ at $\\sim 0.5$~mas from the core to $0.1 c < a_j < 0.25c$ at $\\sim 25$~mas from the core accompanied by some acceleration in the jet flow from Lorentz factor $\\gamma \\lesssim 5$ to $\\gamma \\gtrsim 7$.  The sound speed in the cocoon medium around the jet is less well determined but is less than the sound speed in the jet probably by a factor of 1.5 - 5.  A largest possible viewing angle of 15\\arcdeg\\ implies a jet sound speed at the upper limit of these estimates and somewhat higher flow Lorentz factors. However, jet morphology argues against viewing angles larger than 12\\arcdeg. At smaller viewing angles the jet sound speed is lower and at a 6\\arcdeg\\ viewing angle the jet sound speed is about a factor 2 less but the flow Lorentz factor is comparable. The decline in radio intensity is on the order of what would be associated with isothermal jet expansion.  Knot interknot intensity variations are greater than would be expected from adiabatic compressions associated with the helical twist and we infer the presence of a shock along the leading edge of the helical twist in addition to shock or density structures flowing through the helical pattern. Our results imply that the macroscopic heating of the expanding jet fluid is less than the microscopic energization of the synchrotron radiating relativistic electrons. ",
        "introduction": "\\baselineskip 12pt \\parskip 2pt  The radio source 3C\\,120 is dominated by a variable core and a prominent one-sided jet extending from the core on subparsec scales to about a hundred kiloparsecs \\citep{W87}.  The galaxy has a redshift $z=0.033$ \\citep{B80}, and with $H_o = h~ 100$~km~s$^{-1}$~Mpc$^{-1}$ and $h/0.7 \\equiv h_{70} = 1$ there are $\\sim 0.65$~pc~mas$^{-1}$.  Recent high frequency VLBI observations at 86, 43 and 22 GHz \\citep{G98,G99,G00,G01,M02} have produced high resolution images of the innermost features and motions in the jet.   The 43 and 22 GHz observations, show that the inner jet structure, within 10~mas from the core, is complex and exhibits a subluminal to superluminal transition on very small scales. Global VLBI observations at 5~GHz and 10.7~GHz \\citep{W01} find superluminal motions consistent with those observed at 22 and 43~GHz.  The 3C\\,120 jet shows a variety of knots and side-to-side structures, and it has been suggested that the jet structure of 3C\\,120 inside 10~mas may be described by components interacting with the external medium in their motion along an underlying helical twist \\citep{G98,G00,G01}. VLBI images at 1.7~GHz, see Figure 6 in \\citet{W01}, also show structure suggestive of a helical twist beyond 25~mas but with considerably longer wavelength.  Helical structure in relativistic jets can arise as a result of ordered variation in the flow direction at the central engine, e.g., precession, and/or as a result of random perturbations to the jet flow such as the jet cloud interaction discussed in \\citet{G00}.  Initial random perturbations can trigger or propagate as pinch, helical, or higher order normal modes of jet distortion.  Any normal mode structure will depend on the initial excited wave frequency or frequencies and initial amplitudes, and on the subsequent propagation and growth or damping of these wave frequencies along the jet, e.g., axisymmetric perturbations as in \\citet{H98} and \\citet{A01} or some combination including asymmetric perturbations as in \\citet{A03} and \\citet{HH03}. In these simulations linear and/or non-linear perturbations were applied.  In particular, an injection event can coexist and interact with a precession induced helical twist \\citep{A03}. In general, non-linearly triggered normal modes are damped until reaching a ``linear'' regime and both non-linear triggers and ``linear'' normal mode structures can coexist on the jets. Within the ``linear'' regime normal mode structures can be satisfactorily modeled by the linearized relativistic hydrodynamic (RHD) equations \\citep{H98,A01,HH03}.   The behavior of normal modes and, in particular, helical twist propagation and growth or damping can sensitively depend on jet speed via the Lorentz factor, on the sound speeds in the jet and surrounding material, and on the rate of jet expansion \\citep{H03}.  That the normal wave modes predicted theoretically operate on AGN jets has been suggested by \\citet{LZ01} who successfully fitted the twisted emission threads observed in the 3C\\,273 jet by a combination of helical and elliptical surface and internal normal modes.  Other similar fitting has been performed by \\citet{LHE03} in the context of M\\,87.  Helical twists have been invoked to explain other observed jet structure and, for example, it has been shown that superluminal motions and accelerations along a curved trajectory can be produced by helical jet models, e.g., the radio source 3C\\,345 \\citep{H87,S95}.  While some fitting of helically twisted structure to observed jet morphology has been performed, to date no detailed self-consistent models have been constructed.  Only in the past few years have time dependent relativistic hydrodynamical codes become readily available \\citep{DH94,MMI94,KNM96,FK96,A99a}.  More recently this has allowed fully 3D relativistic jet simulations to be performed with resolution sufficient to compare structure with theoretical predictions and/or with features observed in AGN jets \\citep{A99b,A00,H01,HMD02,HH03,A03}.  Ideally we would like to model observed jet structures numerically and include non-linear effects. In general, computational constraints make it very difficult to simulate and model the superluminal jets that must lie near to the line of sight.  Satisfactory numerical modeling would require very high grid resolution across the jet to reduce the effects of numerical viscosity ($\\sim 100$ computational zones) \\citet{P04} and additional zones outside the jet, lengthy computational grids in order to account for projection, and very large storage capabilities in order to account for light travel time effects, e.g., \\citet{G97,A03}, that make it impractical to conduct numerical modeling involving many trials.  This difficulty can be overcome using theoretical models based on the linearized fluid equations.  In this paper we consider high frequency VLBI images and proper motions from \\citet{G98,G99,G01} that cover the inner 10~mas of the jet, new 5~GHz images that cover the inner jet out to 30~mas, and proper motions from \\citet{W01} that cover this portion of the jet.   With these observations, we begin to discern what trajectory components follow, whether motions are constant for a given feature and between features, and whether there are similarities between component speeds, trajectories, and structure at parsec and tens of parsec scales.  We use the observed structure and proper motions to constrain the macroscopic properties of the jet and surrounding medium.  In \\S 2 the observed relevant observations are summarized.  In \\S 3 we show how to model relativistically moving helical twists along with the accompanying line of sight appearance and model the innermost 10~mas portion of the jet as a single frequency helical twist on an isothermally expanding jet.  In \\S 4 we extend the modeling to larger distances with dual harmonic twist frequencies and also consider adiabatic expansion.  We conclude in \\S 5 with a summary and discussion of the implications of our intensity and dynamical modeling.  \\vspace{-0.7cm} ",
        "conclusions": "We have assumed that prominent components at $z_{ob} \\sim 2 - 30$~mas are associated with a projected helical twist.  These ``helical'' components indicate the wavelength and motion of the helical twist. Nearly constant spacing, $\\lambda^{ob} \\sim 3$~mas, and similar motion, $\\beta_w^{ob} \\sim 4$, of prominent components along with a ``saturated'' component intensity structure between $2 - 10$~mas suggests that the motion of these components could be associated with a saturated helical wave in the high frequency regime.  Between $10 - 15$~mas a discontinuity in the regular pattern of components occurs and between $15 - 30$~mas the component spacing indicates a potential helical pattern with a twist wavelength $\\lambda^{ob} \\sim 8$~mas, or about twice the wavelength inside 10~mas.  The increase in component spacing is accompanied by an apparent increase in the motion of components to $\\beta_w^{ob} \\sim 6$.   Non-helical components outline a curved path in intensity contours between helical components and their motion is assumed indicative of the flow speed $\\beta_f^{ob} \\sim 6 - 6.5$.  In general wave growth, wavelengths, and wave motions are consistent with jet expansion that lies somewhere between adiabatic and isothermal but closer to adiabatic.  Saturation can occur naturally on the expanding jet as a consequence of jet expansion and declining growth rate of helical waves at frequencies above resonance. Modeling the observed component spacing and motion as a helical twist from 2 - 30~mas requires two harmonic twist frequencies.  The higher frequency has a period of $\\lesssim 2$~years.  A typical intrinsic wavelength of the superluminally moving helical twist produced by the two frequencies is six times the jet radius.  Declining jet sound speed accompanied by modest jet acceleration provides the appropriate conditions for growth of multiple frequencies. Still it is necessary to assume a damping mechanism for the high frequency wave that allows the low frequency wave to grow to saturation.  Modeling of the helical twist at the $\\theta = 12\\arcdeg$ viewing angle that has been extensively considered here, suggests a decline in jet sound speed from $c/3 < a_j < c/\\sqrt 3$ at $\\sim 0.5$~mas from the core to $0.1 c < a_j < 0.25c$ at $\\sim 25$~mas from the core accompanied by some acceleration in the jet flow from Lorentz factor $\\gamma \\lesssim 5$ to $\\gamma \\gtrsim 7$.  At $z_{ob} \\sim 15$~mas $\\beta_s \\equiv a_j/c > 0.13$ is obtained in the limit $\\sin\\theta \\rightarrow 0$.  The initial sound speed estimate at $\\sim 0.5$~mas from the core increases to $a_j \\sim c/\\sqrt 3$ at a maximum possible viewing angle of $\\theta \\sim 15\\arcdeg$.  At the larger viewing angle, higher typical Lorentz factor is required but additional jet acceleration also occurs.  Thus, the initial Lorentz factor would not be too much higher than for the 12\\arcdeg\\ viewing angle. However, viewing angles  larger than 12\\arcdeg\\ can be ruled out.  At larger viewing angles relativistic aberration produces northern instead of southern ``helical'' knots when $\\beta_{w}^{ob} > 5$ for $z_{ob} > 10$~mas. At smaller viewing angle the initial sound speed is less and less flow acceleration would occur.  At a viewing angle of 6\\arcdeg\\ the jet sound speed might be about half of the estimates at 12\\arcdeg, and the typical Lorentz factor would be comparable but with less variation along the jet.  For $z_{ob} > 2$~mas the flow Lorentz factor is in the range $\\gamma \\sim 6 - 8$ for viewing angles $\\theta \\sim 4 - 14\\arcdeg$.  Subluminal component motion at $z_{ob} < 1$~mas provides an estimate of the sound speed in the cocoon medium outside the jet.  The subluminal component motion requires an external (cocoon) sound speed $a_x > 0.1c$ at $z_{ob} \\sim 0.5$~mas that is relatively independent of the viewing angle. The value depends on the Lorentz factor at this location.  A lower Lorentz factor requires a higher external sound speed but less than the jet sound speed by a factor 1.5 - 5 depending on viewing angle.  We cannot be more accurate as only the innermost $z_{ob} < 1$~mas of the jet provides any significant information on slow moving pattern speeds.  This estimate was made by assuming that the innermost subluminal components were associated with the helical twist.  However, similar results will be obtained if the components are modeled as {\\it trailing} pinch structures triggered by passage of a shock down the jet \\citep{A01,G01}.  The combination of jet expansion, jet acceleration and sound speed decline implied by our results automatically guarantees that components associated with normal modes will show rapid acceleration and increased spacing at $z_{ob} < 2$~mas.  More modest increase in speeds and spacing occurs at larger core distance.  We cannot determine how the external cocoon sound speed might vary with core distance from the observations considered here.  Basically any determination of the change in external sound speed requires information on low frequency or resonant wave motion far from the core and here we see only high frequency waves (far above resonance) far from the core.  Note that the present values for the external sound speed and an assumed adiabatic decline are consistent with a hot slow wind around a high speed hotter jet.  Such a picture may be reasonable at the parsec to tens of parsec scales that we consider here.  A wind with outflow speeds of a few tenths of lightspeed influences normal mode dynamics and increases the transverse velocity fluctuation produced by a saturated pressure fluctuation, e.g., \\citet{HH03}.  This effect would enhance differential Doppler boosting, possibly to significant levels.  In earlier work \\citet{W01} showed that a flow helical pitch comparable to the intrinsic helical pitch implied by the observations can result in too much knot interknot intensity variation as a result of differential Doppler boosting.  In this case differential Doppler boosting depended strongly on the viewing angle and thus knot interknot intensity variation could be used to constrain the viewing angle. Unfortunately, the present modeling of intensities as adiabatic compressions does not deliver sufficient helical projection plus differential Doppler boosting to produce the observed knot interknot intensity variations at any allowed viewing angle.  Here our first order accurate flow field helical pitch is strongly decoupled and much less than the intrinsic helical pitch of the pressure and density enhancement.  Differential Doppler boosting could be enhanced by a significant wind around the jet.  Alternatively the apparent decoupling between the macroscopic fluid properties implied by an accelerating cooling jet and the microscopic radiating particle properties implied by an isothermal decline in the intensity can be taken to indicate the presence of a shock along the leading edge of the helix.  The presence of a shock would more nearly align the flow field parallel to the shock face and enhance differential Doppler effects.  The investigation of these possible effects is beyond the scope of this paper.  Finally our dynamical modeling implies that macroscopic viscosity and heating are relatively low along this jet, at least when compared to the energization requirements of the radiating electrons implied by our intensity modeling.  A low macroscopic viscosity is consistent with the requirement that a relatively sharp velocity discontinuity between jet and cocoon must be present in order for high frequency short intrinsic wavelength helical patterns to exist. The requirement of sharp velocity  discontinuity indicates a relatively low rate of mass entrainment.  Broad velocity shear, indicative of higher macroscopic viscosity and mass entrainment, would eliminate coherent structures based on short normal mode wavelengths, e.g., \\citet{B91}.  Note however, that a shear layer increasing in thickness proportionally to the jet radius may be one reason for a growing lower frequency wave supplanting a higher frequency wave.  The higher frequency shorter intrinsic wavelength would be damped as the jet expands and the shear layer grows in thickness while the lower frequency longer intrinsic wavelength could continue to grow."
    },
    "0410/astro-ph0410516_arXiv.txt": {
        "abstract": "Stellar-mass black holes have been detected by radial-velocity observations in star/black hole binaries. These allow only the determination of the mass function. Tracking the astrometric orbits of the visible components of the star/black hole binaries would allow the full determination of the black hole masses, which would be of great astrophysical interest. We investigate the possibilities to do this with Gaia. A very promising object seems to be Cyg X-1. The donor star is an O9.7Iab supergiant having a mass of 25 $M_\\odot$, and the mass of the black hole is estimated to be at least 13 $M_\\odot$ masses. The period of 5.6 days implies a semi-major axis of 1.01 $\\mu$pc of the binary orbit. At the distance of Cyg X-1 (2.5kpc) this translates into a semi-major axis of about 28 $\\mu$arcsec for the visible component, which can be detected by Gaia. We discuss further candidates. ",
        "introduction": "Candidates for stellar--mass black holes have preferentially been found as X-ray sources in binary systems by radial velocity measurements of the optically visible companion star. Such measurements allow, however, only the determination of the mass functions \\begin{equation} \\label{eq1} f(M) = {Pa^3_{RV} \\over 2 \\pi G} = {M^3_{BH}~sin~i^3 \\over (M_{BH} + M_*)^2} \\,, \\end{equation} where $P$ denotes the orbital period, $a_{RV}$ is the semi amplitude of the companion star's radial velocity curve and $G$ is the constant of gravity. $M_{BH}$ and $M_*$ are the masses of the black hole and the companion star, respectively. $i$ denotes the inclination of the orbital plane to the line of sight in the sense that $i = 90^\\circ$, if the orbit is seen edge-on. The inclination angle is not known from the radial velocity curve. If the latter is combined with the light curve and a measurement of the projected rotational velocity broadening of the spectral lines of the companion star, $i$~can sometimes be estimated indirectly (Orosz 2003). A direct, much more accurate determination of the mass of the black hole is possible, if the orbit of the companion star is observed also astrometrically. A precise measurement of the black hole mass would help to constrain the radius of the last stable orbit around the black hole. If the orbital period of material on such an orbit can be identified by intensity fluctuations of the system, this would be a physical experiment of great significance, because it would prove the existence of event horizons and allow the measurement of the spin parameter of the black hole. ",
        "conclusions": ""
    },
    "0410/astro-ph0410385_arXiv.txt": {
        "abstract": "{ We report the detection of narrow \\fetfive\\ and \\fetsix\\ X-ray absorption lines at $6.68 \\pm 0.04$~keV and $6.97 \\pm 0.05$~keV in the persistent emission of the dipping low-mass X-ray binary \\src\\ during a 2003 January XMM-Newton observation. These features are superposed on a broad emission feature centered on $6.6\\,^{+0.1}_{-0.2}$~keV. During dipping intervals the equivalent width of the \\fetfive\\ feature increases while that of the \\fetsix\\ feature decreases, consistent with the presence of less strongly ionized material in the line-of-sight.  As observed previously, the changes in the 1.0--10~keV spectrum during dips are inconsistent with a simple increase in absorption by cool material.  However, the changes in both the absorption features {\\it and} the continuum can be modeled by variations in the properties of an ionized absorber. No partial covering of any component of the spectrum, and hence no extended corona, are required.  From persistent to deep dipping the photo-ionization parameter, $\\xi$, expressed in \\xiunit, decreases from \\logxi\\ of $3.9 \\pm 0.1$ to \\logxi\\ of $3.13 \\pm 0.07$, while the equivalent hydrogen column density of the ionized absorber increases from ($3.8 \\pm 0.4$)~\\ttnh\\ to ($37 \\pm 2$)~\\ttnh.  Since highly-ionized absorption features are seen from many other dip sources, this mechanism may also explain the overall changes in X-ray spectrum observed during dipping intervals from these systems. ",
        "introduction": "\\label{sec:intro}   \\src\\ is a faint (3--10 mCrab) low-mass X-ray binary (LMXB) which exhibits X-ray bursts and 2.94~hour periodic intensity dips. The source was first detected by {\\it Uhuru} and {\\it Ariel V} \\citep{forman78apj,warwick81mnras} and the dips and bursts were discovered using EXOSAT \\citep{1323:vdk85ssr,1323:parmar89apj}.  From the burst properties, the distance of \\src\\ was constrained to be between 10--20~kpc \\citep{1323:parmar89apj}. During the dips, which typically last for 30\\% of the orbital cycle, the 1--8 keV intensity varies irregularly with a minimum of $\\sim$50\\% of the average value outside of the dips. These dips are believed to be due to obscuration in the thickened, azimuthally structured, outer regions of an accretion disk \\citep{1916:white82apjl}.  The presence of periodic dips indicates that the source is viewed relatively close to edge-on, at an inclination angle, $i$, of 60--80$\\degmark$ \\citep{frank87aa}, where $i$ is defined as the angle between the line-of-sight and the rotation axis of the accretion disk.  The BeppoSAX 1.0--150~keV spectrum of the \\src\\ persistent emission can be modeled by a cutoff power-law with a photon index, \\phind, of $1.48 \\pm 0.01$ and a folding energy, \\ecut, of $44.1 \\, ^{+5.1} _{-4.4}$~keV together with a blackbody with a temperature of $1.77 \\pm 0.25$~keV which contributes $\\sim$15\\% of the 2--10~keV flux \\citep{1323:balucinska99aa}. From observations with the RXTE Proportional Counter Array, \\citet{1323:barnard01aa} report the presence of a $6.43 \\pm 0.21$~keV Fe emission feature with an equivalent width, $EW$, of $110 \\pm 55$~eV.  Quasi-periodic oscillations at $\\sim$1~Hz were discovered in the persistent emission, dips, and bursts by \\citet{1323:jonker99apjl}.  Radial intensity profiles from the imaging Medium-Energy Concentrator Spectrometer on BeppoSAX indicate the presence of a substantial dust halo around the source \\citep{1323:barnard01aa}.  Such an intense halo is not unexpected since \\src\\ is located close to the galactic plane ($b = 0.5\\degmark$) and its emission is absorbed by interstellar material in the line-of-sight equivalent to $4 \\times 10^{22}$~H~atoms~cm$^{-2}$. This high absorption probably also explains why the IR counterpart is undetected at optical wavelengths \\citep{1323:smale95aj}.  The 1--10~keV spectra of most of the dip sources, including \\src, become harder during dipping.  However, these changes are inconsistent with a simple increase in photo-electric absorption by cool material, as an excess of low-energy photons is usually present. Two approaches have been used to model this spectral evolution. Initially, in the ``absorbed plus unabsorbed'' approach \\citep[e.g.,][]{0748:parmar86apj} the persistent (non-dipping) spectral shape was used to model spectra from dipping intervals. It was included with, and without, additional absorption. The spectral evolution during dipping was accounted for by a large increase in the column density of the absorbed component, and a decrease of the normalization of the unabsorbed component. The latter decrease was attributed to electron scattering in the absorber.  More recently, in the ``progressive covering'', or ``complex continuum'' approach \\citep[e.g.,][]{1916:church97apj}, the X-ray emission is assumed to originate from a point-like blackbody, or disk-blackbody component, together with a power-law component from an extended corona.  This approach models the spectral changes during dipping intervals by the partial and progressive covering of the extended component by an opaque neutral absorber.  The absorption of the point-like component is allowed to vary independently from that of the extended component. This approach has been successfully applied to a number of dipping LMXBs including \\src\\ \\citep[][]{1323:balucinska99aa,1323:barnard01aa}.  The improved sensitivity and spectral resolution of {\\it Chandra} and XMM-Newton is allowing narrow absorption features from highly ionized Fe and other metals to be observed from a growing number of X-ray binaries.  In particular, \\fetfive\\ (He-like) or \\fetsix\\ (H-like) 1s-2p resonant absorption lines near 7~keV were reported from the micro-quasars \\gro\\ \\citep{1655:ueda98apj,1655:yamaoka01pasj}, \\grs\\ \\citep{1915:kotani00apj,1915:lee02apj} and \\seventeen\\ \\citep{1743:miller04apj}, and from the neutron star systems \\cir\\ \\citep{cirx1:brandt96mnras,cirx1:brandt00apjl,cirx1:schulz02apj}, \\gx\\ \\citep{gx13:ueda01apjl,gx13:sidoli02aa,gx13:ueda04apj}, \\mxb\\ \\citep{1658:sidoli01aa}, \\bigdip\\ \\citep{1624:parmar02aa}, \\twelve\\ \\citep{1254:boirin03aa} and \\nineteen\\ \\citep{1916:boirin04aa}.  The lines from \\cir, \\gx\\ and \\seventeen\\ are blue-shifted, indicating that the highly ionized plasma is an outflow in these sources. The systems that exhibit the \\fetfive\\ and \\fetsix\\ features are mainly dipping sources \\citep[see Table~5 of ][]{1916:boirin04aa}.  The lack of any orbital phase dependence of the features (except during dips) suggests that the absorbing plasma is located in a thin cylindrical geometry around the compact object.  Such highly ionized plasma is probably a common feature of accreting binaries, but is preferentially detected in the dipping sources, presumably due to being viewed at inclination angles within $\\sim$$30\\degmark$ of the orbital plane.  Here, we report the detection of narrow \\fetfive\\ and \\fetsix\\ 1s-2p X-ray absorption lines from the dipping LMXB \\src. These features are independently reported by \\citet{1323:church05mnras}. We show that the lines are associated with an ionized absorber which has different properties during persistent and dipping intervals.  We demonstrate that these differences can account for the broad-band spectral changes observed between persistent and dipping intervals. ",
        "conclusions": "We report the detection of narrow \\fetfive\\ (He-like) and \\fetsix\\ (H-like) 1s-2p resonant X-ray absorption lines at 6.68$ \\pm 0.04$ and $6.97 \\pm 0.05$~keV in the persistent emission of \\src.  These features are superposed on a broad emission feature centered on $6.6\\,^{+0.1}_{-0.2}$~keV. During dipping intervals the $EW$ of the \\fetfive\\ feature increases while that of the \\fetsix\\ feature decreases implying the presence of less strongly ionized material in the line-of-sight during dips. The overall 1.0--10~keV spectrum becomes harder during the dips, but this change is inconsistent with a simple increase in absorption by cool material. We demonstrate that both the changes in the continuum {\\it and} absorption lines observed between persistent and dipping intervals can be modeled mainly by changes in the properties of an ionized absorber which has a higher column density and lower ionization parameter during dipping intervals than during persistent intervals.   \\fetfive\\ or \\fetsix\\ absorption lines near 7~keV were reported from the micro-quasars \\gro\\ \\citep{1655:ueda98apj,1655:yamaoka01pasj}, \\grs\\ \\citep{1915:kotani00apj,1915:lee02apj} and \\seventeen\\ \\citep[][the classification as a micro-quasar is preliminary]{1743:miller04apj}, and from the neutron star systems \\cir\\ \\citep{cirx1:brandt96mnras,cirx1:brandt00apjl,cirx1:schulz02apj}, \\gx\\ \\citep{gx13:ueda01apjl,gx13:sidoli02aa,gx13:ueda04apj}, \\mxb\\ \\citep{1658:sidoli01aa}, \\bigdip\\ \\citep{1624:parmar02aa}, \\twelve\\ \\citep{1254:boirin03aa}, \\nineteen\\ \\citep{1916:boirin04aa} and now \\src.  All these systems except \\grs\\ and \\gx\\ show dipping activity indicating that their inclination is 60--80$\\degmark$. Note that the report of dipping activity in \\seventeen\\ is preliminary \\citep{1743:miller04apj}. An inclination of 70$\\degmark$ is attributed to \\grs\\ assuming that the superluminal jets are perpendicular to the accretion disk \\citep{1915:mirabel94nature}. The inclination of \\gx\\ is unknown, but its 24~day periodic energy-dependent X-ray modulation \\citep{gx13:corbet03apj} might indicate a relatively high inclination. Therefore, most of the systems exhibiting \\fetfive\\ or \\fetsix\\ absorption lines are viewed relatively close to edge-on.  This and the lack of any orbital phase dependence of the features \\citep[e.g.,][]{1655:yamaoka01pasj,1254:boirin03aa} except during dips, suggest that the highly-ionized absorber is located in a thin cylindrical geometry around the compact object.  Such highly-ionized disk atmosphere or wind is probably a common feature of accreting binaries, but is preferentially detected in the dipping sources, presumably due to being viewed at inclination angles within $\\sim$$30\\degmark$ of the orbital plane.   {\\it Chandra} HETGS observations revealed blue-shifts of $\\sim$500~\\kms\\ for the lines from \\cir, \\gx\\ and \\seventeen, interpreted as evidence for outflows in these systems \\citep{cirx1:schulz02apj,gx13:ueda04apj,1743:miller04apj}.  The absorption lines from the other binaries do not appear to be strongly blue-shifted. However, most of the results come from ASCA Solid-state Imaging Spectrometer and XMM-Newton EPIC which has a factor $\\sim$4 poorer energy resolution than the HETGS at Fe-K. Therefore, the geometry and dynamics of the highly-ionized plasma remain an open question. Whether the plasma is in- or out-flowing and its velocity might depend on parameters such as the luminosity and the distance between the absorber and the ionizing source.  The upper limit to any blue-shift during persistent emission from \\src\\ is 1020~km~s$^{-1}$.  The emission line detected at $6.6\\,^{+0.1}_{-0.2}$~keV in the persistent spectrum of \\src\\ is very broad, with a $FWHM$ of $2.0\\,^{+0.6}_{-0.4}$~keV.  While relativistic broadening cannot be excluded, the conditions for such an effect to occur are not expected to be met in neutron star systems like \\src\\ as easily as in black hole systems where relativistically broadened lines are usually observed. Furthermore, the inclination of $87.5\\degmark\\,^{+2.5}_{-1.3}$ derived from the Laor model is inconsistent with the inclination of $<$$80\\degmark$ inferred from the absence of eclipses in \\src, assuming a mass of 1.4~\\msun\\ and 1~\\msun\\ for the neutron star and the secondary star, respectively. Alternatively, a secondary mass of $\\approxlt$0.3~\\msun\\ would be required to make the two inclinations consistent with each other. Other mechanisms such as Compton scattering, rotational velocity broadening and emission from a range of ionization states may contribute to the broadening in \\src.  We caution that the detection and width of such an emission Gaussian depends on the continuum modeling which is poorly constrained due to the lack of data above 10~keV. However, an emission line was reported at $6.43 \\pm 0.21$~keV in the 2--20~keV RXTE Proportional Counter Array spectrum of \\src\\ \\citep[the line width \\sig\\ was fixed to 0.4~keV,][]{1323:barnard01aa}.  Furthermore, broad Fe lines have been detected using XMM-Newton from other neutron star systems that show highly-ionized absorption features such as \\mxb\\ \\citep[$FWHM = 1.4 \\, ^{+0.3}_{-0.4}$~keV,][]{1658:sidoli01aa} and \\gx\\ \\citep[$FWHM = 1.9 \\pm 0.5$~keV,][]{gx13:sidoli02aa}.   Changes in the \\fetsix\\ and \\fetfive\\ absorption features between persistent and dipping intervals have been reported from \\cir\\ \\citep{cirx1:schulz02apj}, \\seventeen\\ \\citep{1743:miller04apj}, and from the classical dipping LMXBs \\bigdip\\ \\citep{1624:parmar02aa} and \\nineteen\\ \\citep{1916:boirin04aa}.  From persistent to dipping intervals, the strength of the He-like feature increases while that of the H-like feature decreases, indicating that the material in the line-of-sight is less ionized during dipping. In the case of \\mxb, whilst \\citet{1658:sidoli01aa} report that there is no obvious orbital dependence of the $EW$s of the Fe features, their Fig.~4 shows that the ratio of \\fetfive/\\fetsix\\ $EW$s is at a maximum during dips (orbital phase $\\sim$0.8), consistent with the presence of less-ionized material. \\citet{1254:boirin03aa} did not examine the dip seen from \\twelve\\ due to simultaneous enhanced background counting rate related to solar activity. Despite the lower signal to noise ratio implied during the dip, a re-analysis of the 2001 XMM-Newton observation using the photo-ionized absorber model presented in this paper (assuming the same ionizing continuum) indicates that \\logxi\\ is $6.1 \\pm 0.4 $, $5.1\\,^{+0.4}_{-0.3} $ and $4.9\\,^{+0.7}_{-0.4}$, during persistent, shallow and deep dipping intervals, respectively, clearly demonstrating that the absorber is less ionized during dipping.  In the case of \\src, this absorbing material has an ionization parameter which decreases from a \\logxi\\ of $3.9 \\pm 0.1$ during persistent intervals to $3.13 \\pm 0.07$ during deep dipping. We note that the absolute values of $\\xi$ might suffer from a systematic uncertainty due to mis-evaluation of the ionizing continuum (see Appendix~A). However, this systematic effect does not affect the relative differences in $\\xi$ observed between persistent and dipping intervals.  Thus, it is likely that the absorbing material responsible for the dipping activity in X-ray binaries is in general less-ionized than during persistent intervals.  In the case of \\src, we demonstrate not only that the ionization parameter is lower during dipping, but also that the equivalent hydrogen column density is significantly higher than during persistent intervals. The line-of-sight absorber may be considered to have changed from being ``hot'' (almost fully ionized) to ``warm''. During persistent emission, most of the abundant metals except for Fe are almost fully ionized and the only prominent absorption features are due to \\fetfive\\ and \\fetsix\\ (Fig.~\\ref{fig:transmission}; upper right panel). As the level of dipping increases, absorption from a much wider range of ions is evident because of the higher column density and the lower ionization parameter. The continuum transmission becomes highly affected due to the continuum opacity of the ions and to the presence of many more deep narrow features which blend together (Fig.~\\ref{fig:transmission}, middle and lower right panels). This results in an apparent change in the continuum shape.  The spectral changes during dips from LMXBs are often modeled using the ``progressive covering'', or ``complex continuum'' approach \\citep[e.g.,][]{1916:church97apj,1323:balucinska99aa,1323:barnard01aa}. There, the X-ray emission is assumed to originate from a point-like blackbody, or disk-blackbody component, together with a power-law component from an extended corona.  This approach models the spectral changes during dipping intervals by the partial and progressive covering of the extended component by an opaque neutral absorber. In the case of \\src\\ we have demonstrated that changes in the properties of an ionized absorber provide an alternative explanation for the overall spectral changes during dips.  The changes in the continuum and in the narrow lines are modeled self-consistently in a simple way. No partial covering of any component of the spectrum is required. Consequently, the model does not exclude that the X-ray emission originates from point-like components, and an extended corona is not required.  Since similar changes in the Fe line ratios (and therefore also in ionized absorbers) are observed from many other dip sources, it is likely that this mechanism is important in modifying the continuum shape during dips from these other sources as well. However, answering whether this mechanism is sufficient to model the complex changes seen from the other dip sources, or whether ``progressive covering'' is also required, will have to await further investigation and particularly the high spectral resolution observations of dipping LMXBs expected from Astro-E2.   While this paper was under revision, an analysis of the same XMM-Newton observation of \\src\\ was reported by \\citet{1323:church05mnras}. There, it is argued that our model based on an ionized absorber can be ruled out because ``the increase of electron column density in dipping would cause a decrease in X-ray continuum intensity at every energy due to Thomson scattering by a factor $\\exp[-N_{\\rm e}\\sigma_{\\rm T}]$'', and in particular a too large decrease of the 20--50~keV flux compared to what was observed with BeppoSAX. However, the electron scattering optical depth, $\\tau_{\\rm e}$ (see Sect.~\\ref{sec:narrow}), is energy dependent, and its approximation by the energy independent Thomson term, $N_{\\rm e}\\sigma_{\\rm T}$, does not hold at high energy. In fact, using the numbers quoted in our Table~\\ref{tab:weightedaverages}, we calculate that our model (in which the energy dependent Klein-Nishina formula is used to compute $\\tau_{\\rm e}$) predicts a 20--50~keV flux decrease of $\\sim$25\\% during dips, consistent with the value of less than $10 \\pm 10$\\% measured with BeppoSAX \\citep{1323:balucinska99aa}."
    },
    "0410/astro-ph0410666_arXiv.txt": {
        "abstract": "The mass-transfer scenario links chemical peculiarities with stellar duplicity for an increasing number of stellar classes (classical and dwarf barium stars, subgiant and giant CH stars, S stars without technetium, yellow symbiotic stars, WIRRING stars, Abell-35-like nuclei of planetary nebulae...). Despite these successes, the mass-transfer scenario still faces several problems: What is the mass-transfer mode? Why orbital elements of dwarf barium stars do not fully match those of the classical barium stars? What is the origin of the few non-binary stars among dwarf barium stars? ",
        "introduction": "The impact of the binary nature of a star on its chemical composition seems to have been recognized first in the context of Am and Ap stars in the years 1960--1968 \\citep{Bonsack-61,Fowler-65,Conti-65,VandenHeuvel-67a,VandenHeuvel-67b, VandenHeuvel-68a,VandenHeuvel-68b,VandenHeuvel-68c}. \\citeauthor{VandenHeuvel-68a}\\footnote{At the time, E. Van den Heuvel was still affiliated with the Astronomical Institute of Brussels Free University!} (\\citeyear{VandenHeuvel-68a}) writes: {\\it[many properties of] Ap and Am stars (...) can all be explained if these stars were originally the less massive members in spectroscopic binaries in which the primary during its post-main-sequence evolution filled its Roche lobe, transferred a large fraction of its mass towards the secondary and finally evolved into a white dwarf}. Current views about the role of duplicity among Am and Ap stars will be shortly reviewed in Sects.~\\ref{Sect:bariumdwarfs} and \\ref{Sect:filiation}  Two decades later, the {\\it mass-transfer scenario} popped up again, this time in relation with barium stars and following the discovery by \\citet{McClure-80} of a large frequency of spectroscopic binaries among the family of barium stars. According to McClure, {\\it it is not unreasonable to conclude that BaII stars are all binaries with low-mass secondaries consisting of degenerate objects. It is possible that these systems are such that mass has been lost from a more massive evolving star and deposited onto the present primary, the secondary having now evolved to the white dwarf stage.  The carbon and s-process elements\\footnote{For a review about the s-process of nucleosynthesis and its major astrophysical site, the asymptotic giant branch (AGB) stars, the reader is referred to \\protect\\citet{Lattanzio-Wood-03}} in this case could be dumped onto the present primary-star atmosphere. (...)  The CH stars, which are probably the Population II equivalent of BaII stars, may also, by implication, be binaries}.  These are the first definitions of the {\\it mass-transfer scenario} that appeared in the literature. From the very beginning, this scenario was thus supposed to pollute not only giant companions but also dwarf ones. ",
        "conclusions": "The impact of stellar duplicity on cosmic abundances has been investigated   for more than 35 years now. The mass-transfer scenario, producing polluted stars with abundance patterns otherwise incompatible with their evolutionary stage, is widely accepted as the cause of the chemical peculiarities exhibited by classical barium giants, S(no-Tc) stars, CH stars, yellow symbiotics, WIRRING stars, Abell-35-like binary nuclei of planetary nebulae... Yet this is not the end of the story, as the mass-transfer scenario still faces some difficulties:  \\begin{itemize} \\item Precisely how the mass transfer occurs is not at all understood. Physical insight and constraints come essentially from the comparison of orbital elements of  pre- and post-mass-transfer stellar families. Such a comparison reveals in particular that (i) dynamical effects occuring at periastron are probably decisive for the circularization and/or spiral-in of the binary system; (ii) the total lack of short-period, circular systems among pre-mass-transfer M giants (whereas such systems are numerous among K giants) may hint at the key role of wind mass loss from the mass donor.   \\item Orbital elements do not contradict the filiation post-AGB -- dwarf barium stars. The very specific abundance pattern of post-AGB stars seems to  result from chemical fractionation processes occuring {\\it after} the bulk of the AGB envelope has been transferred to the companion star (the barium star in the making). Abundance analyses of post-AGB companions are needed to confirm that their abundances are consistent with those of barium stars.  \\item The barium dwarf -- barium giant filiation is generally  well supported  by the comparison of their orbital elements, with the exception, however,  of a small excess of short-period, eccentric orbits among dwarf barium (and post-AGB) stars. These systems might disappear through spiral-in when they reach the red giant branch. Alternatively, these barium dwarfs could have a main sequence companion (rather than the expected white dwarf), as is the case, incidentally, for the Tc-poor S star HD~191589 which falls in the same region of the period -- eccentricity diagram. But their chemical peculiarities cannot be attributed then to the mass-transfer scenario. These systems would thus join the small set of Tc-poor S stars with a main-sequence companion (like HD~191589) and of {\\it non-binary} dwarf barium stars and subgiant CH stars, which require as well an alternative to the mass-transfer scenario.   \\end{itemize}"
    },
    "0410/astro-ph0410450_arXiv.txt": {
        "abstract": "We show that it is possible to explain the physical origin of  jet precession in active galactic nuclei  (AGNs) through the misalignment between the rotation axes of the accretion disk  and of the  Kerr black hole.  We apply this scenario to  quasars, Seyfert galaxies and also to the Galactic Center black hole Sgr\\,A$^{\\ast}$, for which signatures of % either jet or disk precession have been found. The formalism adopted is parameterized by the ratio of the precession period to the black hole mass and can be used to put constraints to the physical properties of the accretion disk as well as to the black hole spin in those systems. ",
        "introduction": "It is a well-accepted idea that the center of an AGN harbors a supermassive black hole surrounded by  an accretion disk, the dynamics of which eventually produces  a bipolar jet. In some AGNs, there were found evidence of disk precession. In  Seyfert  galaxies this precession was detected, in direct form, from the variability of double-peaked Balmer lines   (e.g.,  \\citealt{sto03}). In blazars, where the  accretion disk emission is overrun  by the non-thermal jet emission, precession can be inferred indirectly from the jet kinematics (ejection of superluminal knots  in distinct directions and with different apparent proper motions) or from periodic boosting of the continuum emission (e.g., \\citealt{abra00,caab04a}).  A question that could be arisen is: what is the  physical mechanism driving the disk precession?  In Seyfert  galaxies, precession has been attributed  to the relativistic advance of the pericenter in  an elliptical accretion disk (e.g., \\citealt{era95}),  to the precessing single-spiral arm in a  circular disk (e.g., \\citealt{sto03}), or to  a radiation-induced warp in the disk (e.g.,  \\citealt{prin96,sto97}).  For blazars, several authors have proposed a binary black hole system scenario  wherein the primary  accretion disk precesses due to torques induced by the secondary black hole, whose orbital plane does  not coincide with that of the accretion disk (\\citealt{sil88,leva96,katz97,rom00,sti03,caab04a,caab04b}).  Although supermassive binary black hole systems provide precession periods compatible with  observations, the mean separation of the black holes might be sometimes substantially short, so that  their timescale for coalescence might reach less than  1000 years \\citep{capr04}.  In this work, we discuss the possibility that the disk precession arises from torques induced by the  misalignment between the angular momenta of a Kerr black hole and that of the accretion disk, which can  be applied both to blazars and Seyfert galaxies. In $\\S$ 2, we derive the relation between precession  period, mass of the black hole and physical parameters of the accretion disk. In $\\S$ 3 we apply this  formalism to several objects for which precession period and black hole mass are known, and in $\\S$ 4  we present our conclusions. ",
        "conclusions": "In this work we analyzed the possibility of jet/disk precession  in AGNs    be driven by torques in an accretion disk produced by the misalignment     between the angular momenta of the black hole and that of the accretion disk.     We adopted a formalism based on \\citet{lime02}, and obtained an equation  for the ratio of the precession period to the black hole mass, parametrized by the Kerr metric and the disk structure, which  is mostly determined by  physical processes, such as viscosity and radiation. This expression can be  used to extract the physical properties of the     accretion disk, as well as of the black hole spin in astrophysical     systems that present signatures of disk/jet precession.  Two different disk models were used in the calculations: power law and  % exponential surface density profiles. We applied these models to eight  objects for which precession periods and black hole  masses are published. We  showed that it is always possible to obtain a set of parameters that provides a jet/disk precession period compatible with those derived from high-resolution radio maps and continuum variability (e.g., \\citealt{lime02,caab04b}), or from the double-peak emission-line variability (e.g., \\citealt{new97,sto03}). However, not all the possible combinations of $a_\\ast$, $s/\\sigma$ and $R_\\mathrm{out}$ render results compatible with the observation, putting constraints  on the disks  properties if the precession is due to the proposed mechanism."
    },
    "0410/astro-ph0410499_arXiv.txt": {
        "abstract": "A transiting extrasolar planet sequentially blocks off the light coming from the different parts of the disk of the host star in a time dependent manner. Because of the spin of the star, this produces an asymmetric distortion in the line profiles of the stellar spectrum, leading to an apparent anomaly in the radial velocity curves, known as the Rossiter -- McLaughlin effect.  Here, we derive approximate but accurate analytic formulae for the anomaly in the radial velocity curves, taking into account the stellar limb darkening. The formulae are particularly useful in extracting information on the projected angle between the planetary orbit axis and the stellar spin axis, $\\lambda$, and the projected stellar spin velocity, $V\\sin I_s$.  We create mock samples for the radial curves for the transiting extrasolar system HD 209458 and demonstrate that constraints on the spin parameters $(V\\sin I_s,~\\lambda)$ can be significantly improved by combining our analytic template formulae and the precision velocity curves from high-resolution spectroscopic observations with 8--10 m class telescopes. Thus, future observational exploration of transiting systems using the Rossiter -- McLaughlin effect will be one of the most important probes for a better understanding of the origin of extrasolar planetary systems, especially the origin of their angular momentum. ",
        "introduction": "With more than 130 extrasolar planets discovered so far, major scientific purposes in this field are rapidly moving from mere detection to characterization of the planetary systems, i.e., statistics of planetary masses, orbital periods, eccentricities, stellar metallicities, and so on.  Most of these systems have been discovered through the periodic change of radial velocities of the central stars. One of them, HD 209458b, was first discovered spectroscopically and was soon later found to exhibit a transit signature in front of the stellar disk for a duration of $\\sim 2$ hr in its orbital period of 3.5 days \\citep{henry00,Charbonneau00}. More recently, a few additional extrasolar planets have been discovered photometrically from a transiting signature in their light curves and later confirmed spectroscopically \\citep[e.g.,][]{Udalski02c,Bouchy,Alonso}.  Indeed, such transiting planets provide important information for the extrasolar planetary systems that is otherwise unavailable: planetary size, atmospheric composition, and the degree of (mis)alignment of the planetary orbit axis and the stellar spin axis.  Among others, \\citet{Queloz00} showed that the planetary orbit and the stellar rotation of the HD 209458 system share the same direction; the planet sequentially blocks off the light from the approaching and then from the receding parts of the stellar surface. This produces a distortion in the line profiles of the stellar spectrum during the transit in a time-dependent manner, leading to an anomaly of the radial velocity curves, previously known as the Rossiter -- McLaughlin (RM) effect in eclipsing binary stars \\citep{Rossiter24,McLaughlin24,Kopal99}. The signature of the RM effect was first reported by \\citet{Schlesinger1909} and the effect was later isolated and measured by \\citet{Rossiter24} and \\citet{McLaughlin24}. The first attempt at a theoretical study of the RM effect was made by \\citet{Petrie38} and the analysis was subsequently  extended by \\citet{Kopal42,Kopal45} to incorporate the effects of limb and gravity darkening and rotational and tidal distortion. In the context of extrasolar planetary systems, \\citet{Queloz00} numerically computed the expected amplitude of the radial velocity anomaly caused by the RM effect. They first put constraints on the stellar spin angular velocity and its direction angle with respect to the planetary orbit by comparing the observed velocity anomalies. Incidentally, \\citet{Snellen04} proposed an interesting suggestion to use the RM effect as an efficient diagnostic of the atmosphere of a transiting planet.  Obviously, this methodology provides unique and fundamental clues to understanding the formation process of extrasolar planetary systems. Planets are supposed to form in the proto planetary disk surrounding the proto star \\citep[e.g.,][]{Pollack96}. Thus, the stellar spin and the planetary orbital axes are expected to be aligned. In turn, any constraints on their (mis)alignment degree are useful clues to the origin of the angular momentum of planets and its subsequent evolution during possible migration of the planets into the close-in orbits \\citep[e.g.,][]{Lin96}. In order to improve the reliability and precision of such results, analytic templates are of great value.  In this paper, we show analytic formulae for radial velocities of transiting extrasolar planets.  Those formulae can be used as standard templates in constraining a set of parameters for numerous transiting planets that will be detected in the near future with {\\it Corot} and {\\it Kepler}.  The rest of the paper is organized as follows. Section \\ref{sec:radvel_non-transit} introduces a variety of parameters that characterize the dynamics of a planetary system and summarizes the radial velocity curve, neglecting the transit effect (or, equivalently, for a non rotating star). Section \\ref{sec:radvel_transit} shows a general theoretical framework of the RM effect for a star during the planetary transit. Sections \\ref{sec:without_darkening} and \\ref{sec:limb_darkening} derive analytic expressions for radial velocity curves; in \\S \\ref{sec:without_darkening}, we consider an idealistic case of the stellar intensity model without limb darkening and derive the exact analytic expressions. Notation used in the derivation of are summarized in Table 1. On this basis, \\S \\ref{sec:limb_darkening} presents approximate formulae taking into account the stellar limb-darkening effect.  We apply these analytic templates to a transiting system, HD 209458, in \\S \\ref{sec:HD209458} and examine the sensitivity to the parameters of the system.  Finally, \\S \\ref{sec:conclusion} is devoted to the main conclusions and discussion. ",
        "conclusions": "\\label{sec:conclusion}  We have discussed a methodology to estimate the stellar spin angular velocity and its direction angle with respect to the planetary orbit for transiting extrasolar planetary systems using the RM effect previously known in eclipsing binary stars \\citep{Rossiter24,McLaughlin24,Kopal99}. In particular we have derived analytic expressions of the radial velocity anomaly, $\\Delta v_s$, which are sufficiently accurate for extrasolar planetary systems. If the stellar limb darkening is neglected, the expression is exact. We have extended the result to the case with limb darkening and obtained approximate but accurate analytic formulae.  For a typical value of $\\gamma=R_p/R_s\\sim0.1$, the formulae reduce to a simple form (eqs.[\\ref{eq:dV_s_limbdarkening_A}], [\\ref{eq:approx_W1}], [\\ref{eq:approx_W2}], [\\ref{eq:dV_s_limbdarkening_B}], [\\ref{eq:approx_W3}], and [\\ref{eq:approx_W4}]): \\begin{eqnarray} \\Delta v_{\\rm s}=  \\Omega_s X_p \\sin I_s \\,\\, \\frac{ \\gamma^2\\,\\{1-\\epsilon(1-W_2)\\}} { 1-\\gamma^2-\\epsilon\\left\\{\\frac{1}{3}-\\gamma^2\\right\\}} \\end{eqnarray} during the complete transit phase and \\begin{eqnarray} \\Delta v_{\\rm s} &=&  \\Omega_s X_p \\sin I_s \\,\\,\\frac{\\displaystyle (1-\\epsilon)\\left\\{-z_0\\zeta + \\gamma^2\\cos^{-1}(\\zeta/\\gamma)\\right\\} + \\frac{\\epsilon}{1+\\eta_p}\\,\\, W_4} {\\pi\\left(1-\\frac{1}{3}\\epsilon\\right)- (1-\\epsilon)\\left\\{\\sin^{-1}z_0-(1+\\eta_p)z_0 +\\gamma^2\\cos^{-1}(\\zeta/\\gamma)\\right\\}} \\end{eqnarray} during the egress/ingress phases, where \\begin{eqnarray} W_2 & = & \\frac{(R_s^2-X_p^2-Z_p^2)^{1/2}}{R_s}, \\\\ W_4 & = & \\frac{\\pi}{2}\\,\\,\\gamma^{3/2}(2-\\gamma)^{1/2}(\\gamma-\\zeta)\\,\\,x_c \\,\\,\\frac{g(x_c\\,;\\,\\eta_p,\\,\\gamma)} {g(1-\\gamma;\\,-\\gamma,\\,\\gamma)}, \\end{eqnarray} where $g(x;a,b)$ is defined in equation (\\ref{appen:func_g}).  The definition and the meaning of the variables in the above expressions are summarized in Table \\ref{tab:notation}.  The numerical accuracy of the above formulae was checked using a specific example of the transiting extrasolar planetary system, HD 209458, and we found that they are accurate within a few percent.  Our analytic formulae for the radial velocity anomaly are useful in several ways. One can estimate the planetary parameters much more efficiently and easily, since one does not have to resort to computationally demanding numerical modeling. Furthermore, the resulting uncertainties of the fitted parameters and their correlations are easily evaluated.  To illustrate these advantages specifically, we performed a parameter estimation applying the formulae against mock data for HD 209458. We showed that with precision data obtainable by 8--10 m class telescopes such as Subaru HDS, our formulae improve the efficiency and robustness of estimating the spin parameters, $V\\sin I_s$ and $\\lambda$. Furthermore, the combined data analysis with asteroseismology \\citep[e.g.,][]{Gizon2003} and/or the correlation between the mean level of emission and the rotation period \\citep[e.g.,][]{Noyes1985} may break the degeneracy between $V$ and $I_s$ in extrasolar planetary systems.  Among the recently discovered transiting extrasolar planetary systems, i.e., TrES-1 by the Trans-Atlantic Exoplanet Survey \\citep[][]{Alonso} and OGLE-TR 10, 56, 111, 113, 132 by the Optically Gravitational Lens Event survey \\citep[e.g.,][]{Udalski02a,Udalski02b,Udalski02c,Udalski03,Konacki03,Bouchy,Pont04}, TrES-1 has similar orbital period and mass to those of HD 209458b, but its radius is smaller. Thus, it is an interesting target to determine the spin parameters via the RM effect; if its planetary orbit and the stellar rotation share the same direction as discovered for the HD 209458 system, it would be an important confirmation of the current view of planet formation out of the protoplanetary disk surrounding the protostar. If not, the result would be more exciting and even challenge the standard view, depending on the value of the misalignment angle $\\lambda$.   We also note that the future satellites {\\it Corot} and {\\it Kepler} will detect numerous transiting planetary systems, most of which will be important targets for the RM effect in 8 - 10 m class ground-based telescopes.  We hope that our analytic formulae presented here will be a useful template in estimating parameters for those stellar and planetary systems.  Finally, it is interesting to note that the RM effect may even be used as yet another new detection method of transiting planetary systems.  For the HD 209458 system, the stellar radial velocity amplitude due to the Kepler motion is around 100 m s${}^{-1}$. Since the stellar rotation velocity is around 4 km s${}^{-1}$, the maximum radial velocity anomaly due to the RM effect is $\\sim 40 (\\gamma/0.1)^2$m s${}^{-1}$ and thus is very close to the former. On the other hand, the latter could be significantly larger if the host star rotates faster, and/or the host star (the planet) has a smaller (larger) radius.  In extreme cases, the radial velocity curve, with a periodicity of a few days for instance, is barely detectable, but the velocity anomaly is quite visible for a few hours of the transiting phase. Thus, the conventional radial velocity curve analysis might have missed some of the potentially interesting spectroscopic signature of transiting planets.  A search for such specific signatures deserves an attempt against existing or future spectroscopic samples that do not show any clear conventional feature of radial velocity periodicity.  In conclusion, we have demonstrated that the radial velocity anomaly due to the RM effect provides a reliable estimation of spin parameters. Combining data with the analytic formulae for radial velocity shift $\\Delta v_s$, this methodology becomes a powerful tool in extracting information on the formation and the evolution of extrasolar planetary systems, especially the origin of their angular momentum. Although it is unlikely, we may even speculate that a future RM observation may discover an extrasolar planetary system in which the stellar spin and the planetary orbital axes are antiparallel or orthogonal.  This would have a great impact on the planetary formation scenario, which would have to invoke an additional effect from possible other planets in the system during the migration or the capture of a free-floating planet. While it is premature to discuss such extreme possibilities at this point, the observational exploration of transiting systems using the RM effect is one of the most important probes for a better understanding of the origin of extrasolar planets."
    },
    "0410/nucl-th0410091_arXiv.txt": {
        "abstract": "The first lecture provides an introduction to the physics of color superconductivity in cold dense quark matter. The main color superconducting phases are briefly described and their properties are listed. The second lecture covers recent developments in studies of color superconducting phases in neutral and $\\beta$-equilibrated matter. The properties of gapless color superconducting phases are discussed. ",
        "introduction": "\\subsection{Dense baryonic matter}  Almost all matter around us is made of several dozen chemical elements and their isotopes. Each element (atom) has a compact nucleus, made of protons and neutrons, and a cloud of electrons surrounding the nucleus. The proton and the neutron are baryons, i.e., ``heavy'' particles. They are about two thousand times heavier than the electron. (The actual values of the masses are $m_{\\rm p} \\approx 938.3\\,\\mbox{MeV}/c^2$, $m_{\\rm n} \\approx 939.6\\,\\mbox{MeV}/c^2$, and $m_{\\rm e} \\approx 0.511\\,\\mbox{MeV}/c^2$, see Ref.~\\cite{HDG}.) Thus, the mass of an atom comes mostly from its heavy nucleus. A typical size of a nucleus is about five orders of magnitude smaller than a typical atomic size. It is clear, therefore, that the density of matter inside a nucleus is much larger than the density of matter made of chemical elements. In fact, typical densities inside nuclei are of order $10^{14}\\,\\mbox{g}/ \\mbox{cm}^{3}$. To get an impression how large this is, it suffices to say that a tablespoon of a material with this density would weigh about a billion tonnes. Creating even a tiny macroscopic sample of such matter in laboratory is impossible. On the other hand, baryonic matter at even higher densities exists in the Universe in the interior of compact stars. Therefore, it is of fundamental importance to study such very dense matter.  Protons and neutrons, as well as all mesons and baryons that may appear in dense matter, are strongly interacting particles. Their properties and the properties of dense baryonic matter are, in principle, described by the microscopic theory of strong interactions, Quantum Chromodynamics (QCD). The only problem is that QCD is a rather complicated non-Abelian gauge theory. The Lagrangian density of the theory is written in terms of quark fields which carry color charges. Because of the property of confinement, however, quarks cannot exist in vacuum as free particles. Instead, they combine with other quarks and antiquarks to form color neutral hadrons (i.e., mesons and baryons) whose description is not easy in the framework of QCD.  It is known that baryons are not point-like particles. They have a typical size of about $1\\,\\mbox{fm}=10^{-13}\\,\\mbox{cm}$. In sufficiently dense matter, therefore, baryons can be forced to stay so close to one another that they would overlap. At such densities, constituent quarks are shared by neighboring baryons and, with increasing the density further, quarks should eventually become mobile over large distances. This means that quarks become deconfined. When this happens, it does not make sense to talk about hadronic matter any more. It is transformed into quark matter. The properties of such matter is the topic of these lectures.  It was suggested long time ago that quark matter may exist inside central regions of compact stars \\cite{quark-star}. By making use of the property of asymptotic freedom in QCD \\cite{GWP}, it was argued that quarks interact weakly, and that realistic calculations taking full account of strong interactions are possible for sufficiently dense matter \\cite{ColPer}. The argument of Ref.~\\cite{ColPer} consisted of the two main points: (i) the long-range QCD interactions are screened in dense medium causing no infrared problems, and (ii) at short distances, the interaction is weak enough to allow the use of the perturbation theory. As we shall see below, the real situation in dense quark matter is slightly more subtle.  \\subsubsection{Weakly interacting quark matter}  By assuming that very dense matter is indeed made of weakly interacting quarks, one could try to understand the thermodynamic properties of the corresponding ground state by first completely neglecting the interaction between quarks. In order to construct the ground state, it is important to keep in mind that quarks are fermions, i.e., particles with a half-integer spin, $s=1/2$. Therefore, quarks should obey the Pauli exclusion principle which says that no two identical quarks can simultaneously occupy the same quantum state.  In the ground state of zero temperature non-interacting quark matter, quarks occupy all available quantum states with the lowest possible energies. This is described formally by the following quark distribution function: \\begin{equation} f_F(\\mathbf{k}) =\\theta\\left(\\mu-E_\\mathbf{k}\\right), \\quad \\mbox{at} \\quad T=0, \\label{f_F_T=0} \\end{equation} where $\\mu$ is the quark chemical potential, and $E_\\mathbf{k}\\equiv \\sqrt{k^2+m^2}$ is the energy of a free quark (with mass $m$) in the quantum state with the momentum $\\mathbf{k}$ (by definition, $k\\equiv |\\mathbf{k}|$). As one can see, $f_F(\\mathbf{k})=1$ for the states with $k<k_F\\equiv \\sqrt{\\mu^2-m^2}$, indicating that all states with the momenta less than the Fermi momentum $k_F$ are occupied. The states with the momenta greater than the Fermi momentum $k_F$ are empty, i.e., $f_F(\\mathbf{k})=0$ for $k>k_F$.  In the chiral limit, i.e., when the quark mass is vanishing, the pressure of zero temperature free quark matter is given by \\cite{Kapusta} \\begin{equation} P^{(0)} = 2 N_f N_c \\int \\frac{d^3\\mathbf{k}}{(2\\pi)^3} \\left(\\mu-E_\\mathbf{k}\\right) \\theta\\left(\\mu-E_\\mathbf{k}\\right) -B =\\frac{N_f N_c}{12 \\pi^2} \\mu^4 - B, \\label{pressure_0} \\end{equation} where the overall factor $2 N_f N_c$ counts the total number of degenerate quark states, namely 2 spin states (i.e., $s=\\pm 1/2$), $N_f$ flavor states (e.g., up, down and strange), and $N_c$ color states (e.g., red, green and blue). The extra term $B$ in the expression for the pressure (\\ref{pressure_0}), called the bag constant, was added by hand. This term effectively assigns a nonzero contribution to the vacuum pressure and, in this way, provides the simplest modelling of the quark confinement in QCD \\cite{bag-model}. [Note that the vacuum pressure is higher than the pressure of quark matter when $\\mu<\\left(4\\pi^2 B/N_f\\right)^{1/4}$.] Here, we ignore the requirements of the charge neutrality and the $\\beta$ equilibrium in quark matter. These will be addressed in detail in the second lecture. The energy density of quark matter reads \\begin{equation} \\epsilon^{(0)} \\equiv \\mu \\frac{\\partial P^{(0)}}{\\partial \\mu} - P^{(0)} =\\frac{N_f N_c}{4 \\pi^2} \\mu^4 +B. \\label{e_0} \\end{equation} Thus, one arrives at the following equation of state: \\begin{equation} P^{(0)}=\\frac{1}{3} \\left(\\epsilon^{(0)} -4B\\right). \\label{EoS_0} \\end{equation} This equation of state could be easily generalized to the case of nonzero quark masses and/or to the case of non-equal chemical potentials for different quarks. Also, it could be further improved by adding the lowest order corrections due to the interaction. For example, to leading order in coupling, the interaction results in the following correction to the pressure \\cite{FreMcL}: \\begin{equation} \\delta P^{(0)} =-\\frac{\\alpha_s N_f (N_c^2-1)}{16 \\pi^3} \\mu^4 +O\\left(\\alpha_s^2\\right), \\label{pressure_cor} \\end{equation} where $\\alpha_s\\equiv g^2/4\\pi$ is the value of the running coupling constant of strong interactions defined at the scale of the quark chemical potential. The next to leading order corrections can also be calculated \\cite{FreMcL}.  It is generally believed that the equation of state in Eq.~(\\ref{pressure_0}), or its generalization, provides a good approximation for the description of weakly interacting quark matter. Below we argue, however, that some underlying assumptions regarding the ground state of such quark matter are not correct.  \\subsection{Cooper instability and color superconductivity}  It appears that the perturbative ground state of quark matter, characterized by the distribution function in Eq.~(\\ref{f_F_T=0}), is unstable when there is an attractive (even arbitrarily weak in magnitude\\,!) interaction between quarks. This is because of the famous Cooper instability \\cite{Cooper} that develops as a result of the formation of Cooper pairs $\\langle q_{\\mathbf{k}}\\, q_{-\\mathbf{k}}\\rangle$ made of quarks from around the highly degenerate Fermi surface, i.e., quarks with the absolute value of momenta $k\\simeq k_F$. Cooper pairs are bosons, and they occupy the same lowest energy quantum state at zero temperature, producing a Bose condensate. In the presence of such a condensate of Cooper pairs, the ground state of quark matter becomes a (color) superconductor. This is very similar to the ground state of the electron system in the Bardeen-Cooper-Schrieffer theory of low-temperature superconductivity \\cite{BCS}. The only real difference comes from the fact that quarks, unlike electrons, come in various flavors (e.g., up, down and strange) and carry non-Abelian color charges. To emphasize this difference, superconductivity in quark matter is called {\\it color superconductivity}. For recent reviews on color superconductivity, see Ref.~\\cite{reviews}.  It was known for a long time that dense quark matter should be a color superconductor \\cite{BarFra,Bail}. In many studies, however, this fact was commonly ignored. Only recently, the potential importance of this phenomenon was appreciated. To large extent, this was triggered by the observation in Ref.~\\cite{cs} that the value of the color superconducting gap could be as large as $100\\,\\mbox{MeV}$ at baryon densities existing in the central regions of compact stars, i.e., at densities which are a few times larger than the normal nuclear density, $n_0\\simeq 0.15 \\mbox{~fm}^{-3}$. This very natural estimate for the value of the gap in QCD, in which a typical energy scale itself is $200\\,\\mbox{MeV}$, opened a wide range of new theoretical possibilities, and the subject bursted with numerous studies. The main reason is that the presence of such a large energy gap in the quark spectrum may allow to extract clear signatures of color superconducting states of matter in observational data from compact stars.  As in low-temperature superconductors, one of the main consequences of color superconductivity in dense quark matter is the appearance of a nonzero energy gap in the one-particle spectrum, \\begin{equation} {\\cal E}_{\\mathbf{k}} = \\sqrt{\\left(E_{\\mathbf{k}}-\\mu\\right)^2+\\Delta^2}, \\label{E_qp} \\end{equation} where $\\Delta$ is the gap. The presence of the gap in the energy spectrum should affect transport properties (e.g., conductivities and viscosities) of quark matter. Thus, if quark matter exists in the interior of compact stars, this will be reflected, for example, in the cooling rates and in the rotational slowing down of such stars. Also, a nonzero gap modifies thermodynamic properties, e.g., the specific heat and the equation of state. In application to stars, this could modify theoretical predictions for the mass-radius relations, or even suggest the existence of a new family of compact stars.  Color superconductivity may affect directly as well as indirectly many other observed properties of stars. In some cases, for example, superconductivity may be accompanied by the baryon superfluidity and/or the electromagnetic Meissner effect. If matter is superfluid, rotational vortices would be formed in the stellar core, and they would carry a portion of the angular momentum of the star. Because of the Meissner effect, the star interior could become threaded with magnetic flux tubes. In either case, the star evolution may be affected.  In general, it is of great phenomenological interest to perform a systematic study of all possible effects of color superconductivity in compact stars. Before this can be done, however, one needs to know the structure of the QCD phase diagram and properties of various color superconducting phases in detail. Despite the recent progress in the field, such knowledge still remains patchy. While many different phases of quark matter have been proposed, there is no certainty that all possibilities have already been exhausted. This is especially so when additional requirements of charge neutrality and $\\beta$ equilibrium are imposed.  In the rest of this lecture, the main color superconducting phases, that involve different number of pairing quark flavors, will be reviewed briefly. The role of charge neutrality and $\\beta$ equilibrium in quark matter will be discussed in the second lecture.  \\subsection{Two-flavor color superconductivity ($N_f=2$)}  It is instructive to start this lecture from the description of the simplest color superconducting phase, namely the two-flavor color superconductor (2SC). This is a color superconducting phase with the spin-0 Cooper pairing in quark matter made of up and down quarks.  As we discussed in the previous subsection, color superconductivity comes as a result of the Cooper instability driven by an attractive interaction between quarks. To study this instability in detail, one needs to specify the origin and the strength of the interaction. At very large densities, when quark matter is weakly interacting, the dominant interaction between quarks is provided by the one-gluon exchange. Then, the Cooper pairing dynamics, including the screening effects due to dense medium as well as the effects due to higher order corrections, can be studied from first principles within the framework of QCD. Some key details of the analysis in this regime are given below.  We consider quark fields $\\psi_i^a$ in the fundamental representation of the SU(3)$_c$ color gauge group. These fields carry flavor ($i=1,2$) and color ($a=1,2,3$) indices. The QCD Lagrangian density reads \\begin{equation} {\\cal L}_{\\rm QCD} = \\bar\\psi_i^a\\left(i\\gamma^\\mu \\partial_\\mu +\\gamma^0\\mu -m^{(0)}_i\\right) \\psi_i^a +g A^{A}_{\\mu} \\bar\\psi_i^a \\gamma^\\mu T^{A}_{ab} \\psi_i^b -\\frac{1}{4}G^{A}_{\\mu\\nu}G^{A,\\mu\\nu}, \\label{QCD-action} \\end{equation} where $A^{A}_{\\mu}$ is the vector gauge field in the adjoint representation of SU(3)$_c$, $G^{A}_{\\mu\\nu}= \\partial_\\mu A^{A}_{\\nu} -\\partial_\\nu A^{A}_{\\mu}+g f^{ABC}A^{B}_{\\mu}A^{C}_{\\nu}$ is the field strength, and the generators of color transformations are defined as $T^{A}_{ab} \\equiv \\frac{1}{2}\\left(\\lambda^{A} \\right)_{ab}$ where $\\lambda^{A}$ are the Gell-Mann matrices. The current quark masses and the quark chemical potential are denoted by $m^{(0)}_i$ and $\\mu$, respectively. At high densities, one can neglect small current masses of quarks. Then, the QCD Lagrangian density (\\ref{QCD-action}) becomes invariant under SU(2)$_L\\times$SU(2)$_R$ global chiral transformations.  In QCD, the constituent masses of quarks $m_i$ are generated dynamically, and they can be very different from the current masses $m^{(0)}_i$, appearing in the Lagrangian density. At zero density, for example, typical values of the constituent quark masses are of order $\\frac{1}{3} m_{\\rm n}\\approx 313\\,\\mbox{MeV}$ even in the chiral limit. At high densities, on the other hand, the masses of the up and down quarks become small. This is because they are proportional to the value of the chiral condensate $\\langle \\bar\\psi_L \\psi_R\\rangle$ which melts in dense matter. Thus, it is often also justified to neglect the constituent quark masses in studies of color superconducting phases.  The interaction vertex in the Lagrangian density (\\ref{QCD-action}) has a non-trivial color structure given by the color generators $T^{A}_{ab}$. As a result, the quark-quark scattering amplitude in the one-gluon exchange approximation is proportional to the following color tensor: \\begin{equation} \\sum_{A=1}^{N_c^2-1} T^{A}_{a a^\\prime} T^{A}_{b^\\prime b} = -\\frac{N_c+1}{4N_c}\\left(\\delta_{a a^\\prime}\\delta_{b^\\prime b} -\\delta_{ab^\\prime}\\delta_{a^\\prime b}\\right) +\\frac{N_c-1}{4 N_c}\\left(\\delta_{a a^\\prime}\\delta_{b^\\prime b} +\\delta_{a b^\\prime}\\delta_{a^\\prime b}\\right). \\label{T-T} \\end{equation} The first antisymmetric term corresponds to the {\\em attractive} antitriplet channel, while the second symmetric term corresponds to the repulsive sextet channel, see Fig.~\\ref{fig-1-glu}. It is the first antisymmetric antitriplet channel that plays the crucial role in Cooper pairing. \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.4\\textwidth]{one-gluon1.eps} \\caption{The diagrammatic representation of the one-gluon exchange interaction between two quarks in QCD. The color structure of the corresponding amplitude contains an antisymmetric antitriplet and a symmetric sextet channel.} \\label{fig-1-glu} \\end{center} \\end{figure}  Here, it is appropriate to note that quark matter is unlikely to be truly weakly interacting at moderate densities existing in the central regions of compact stars, i.e., $n\\lesssim 10n_0$. In this case, the use of the microscopic theory of strong interactions is very limited, and one has to rely on various effective models of QCD. A very simple type of such a model is the Nambu-Jona-Lasinio (NJL) model with a local four-fermion interaction \\cite{NJL}. Originally, this model was introduced for the description of nucleons with dynamically generated masses. Nowadays, it is also commonly used for the description of quarks. For example, the NJL type models were used in Refs.~\\cite{cs,4fermi,4fermi-rev,huang_2sc} for the description of color superconducting phases. One of the simplest NJL models, that respects the SU(2)$_L\\times$SU(2)$_R$ global chiral symmetry (in the limit $m^{(0)}_i\\to 0$), is defined by the following Lagrangian density \\cite{SKP}: \\begin{eqnarray} {\\cal L}_{\\rm NJL} &=& \\bar\\psi_i^a\\left(i\\gamma^\\mu \\partial_\\mu +\\gamma^0\\mu -m^{(0)}_i\\right) \\psi_i^a + G_S\\left[(\\bar\\psi \\psi)^2 + (i\\bar\\psi \\gamma_5\\vec{\\tau}\\psi)^2\\right] \\nonumber\\\\ &+&G_D (i \\bar{\\psi}^C \\varepsilon \\epsilon^a \\gamma_5 \\psi ) (i \\bar{\\psi} \\varepsilon \\epsilon^a \\gamma_5 \\psi^C), \\label{NJL-action} \\end{eqnarray} where $\\psi^C=C \\bar{\\psi}^T$ is the charge-conjugate spinor and $C=i\\gamma^2\\gamma^0$ is the charge conjugation matrix. The matrix $C$ is defined so that $C\\gamma_{\\mu} C^{-1}=-\\gamma_{\\mu}^{T}$. Regarding the other notation, $\\vec{\\tau}=(\\tau^1,\\tau^2,\\tau^3)$ are the Pauli matrices in the flavor space, while $(\\varepsilon)^{ik} \\equiv \\varepsilon^{ik}$ and $(\\epsilon^a)^{bc} \\equiv \\epsilon^{abc}$ are the antisymmetric tensors in the flavor and in the color spaces, respectively. The dimensionful coupling constant $G_S=5.01\\,\\mbox{GeV}^{-2}$ and the momentum integration cutoff parameter $\\Lambda=0.65\\,\\mbox{GeV}$ (which appears only in loop calculations) are adjusted so that the values of the pion decay constant and the value of the chiral condensate take their standard values in vacuum QCD: $F_\\pi = 93$ MeV and $\\langle\\bar{u}u\\rangle=\\langle\\bar{d}d\\rangle =(-250\\,\\mbox{MeV})^3$ \\cite{SKP}. The strength of the coupling constant $G_D$ is taken to be proportional to the value of $G_S$ as follows: $G_D=\\eta G_S$ where $\\eta$ is a dimensionless parameter of order 1. Note that the value of $\\eta$ is positive which corresponds to having the antisymmetric diquark channel {\\em attractive}. This is motivated by the microscopic QCD interaction, as well as by the instanton motivated models \\cite{cs}.   \\subsubsection{Color and flavor structure of the condensate}  Now let us discuss the color and flavor structure of the condensate of Cooper pairs in the 2SC ground state. This can be determined from the nature of the diquark interaction and the Pauli exclusion principle.  The color structure of the one-gluon exchange interaction in Eq.~(\\ref{T-T}) shows that the color antisymmetric antitriplet channel is attractive. The antisymmetric color channel is also attractive in the instanton motivated models of Ref.~\\cite{cs}. This is reflected in the color structure of the condensate of spin-0 Cooper pairs, \\begin{equation} \\left\\langle \\left(\\bar{\\psi}^C\\right)_i^a \\gamma^5 \\psi_j^b \\right\\rangle \\sim \\varepsilon_{ij} \\epsilon^{abc} , \\label{2sc-cond} \\end{equation} which is antisymmetric in the color indices of the constituent quarks. As one can check (not shown explicitly), this spin-0 condensate is also antisymmetric in Dirac indices. Then, in accordance with the Pauli exclusion principle, it should be antisymmetric in flavor indices as well. The presence of the $\\gamma^5$ matrix on the left hand side of Eq.~(\\ref{2sc-cond}) makes the condensate a scalar. A similar condensate without $\\gamma^5$ is a pseudoscalar which is disfavored in QCD because of the instanton effects.  The color antitriplet condensate (\\ref{2sc-cond}) may have an arbitrary orientation in the color space. By making use of the global color transformations, this orientation can be changed as convenient. It is conventional to choose the condensate to point in the third (blue) color direction, $\\left\\langle \\left(\\bar{\\psi}^C\\right)_i^a \\gamma^5 \\psi_j^b \\right\\rangle \\sim \\varepsilon_{ij} \\epsilon^{ab3}$. In this case, the Cooper pairs in the 2SC phase are made of the red and green quarks only, and the blue quarks do not participate in the pairing. These unpaired blue quarks give rise to gapless quasiparticles in the low-energy effective theory.  The flavor antisymmetric structure in Eq.~(\\ref{2sc-cond}) corresponds to a singlet representation of the global SU(2)$_L\\times$SU(2)$_R$ chiral group. This means that the chiral symmetry is not broken in the 2SC ground state. In fact, there are no other global continuous symmetries which are broken in the 2SC phase. There exist, however, several approximate symmetries which are broken. One of them is the approximate U(1)$_A$ symmetry which is a good symmetry at high density when the instantons are screened \\cite{scr_inst}. Its breaking in the 2SC phase results in a pseudo-Nambu-Goldstone boson \\cite{low-e-2sc}. Additional four exotic pseudo-Nambu-Goldstone states may appear as a result of a less obvious approximate axial color symmetry discussed in Ref.~\\cite{pseudoNG}.  In the 2SC ground state, the vector-like SU(3)$_c$ color gauge group of QCD is broken down to the SU(2)$_c$ subgroup. Therefore, five out of total eight gluons of SU(3)$_c$ become massive due to the Anderson-Higgs mechanism. The other three gluons, which correspond to the unbroken SU(2)$_c$, do not interact with the gapless blue quasiparticles. They give rise to a pure SU(2)$_c$ gluodynamics. The red and green quasiparticles decouple from this low-energy SU(2)$_c$ gluodynamics because they are gapped \\cite{RSS}.  \\subsubsection{Quark propagator in the Nambu-Gorkov basis}  To simplify the study of diquark condensates in dense quark matter, instead of using the Dirac spinors for quark fields, it is more convenient to introduce the following eight-component Nambu-Gorkov spinors: \\begin{equation} \\Psi \\equiv \\left(\\begin{array}{l}\\psi \\\\ \\psi^C \\end{array}\\right), \\end{equation} where $\\psi^C$ is the charge-conjugate spinor, defined earlier. The structure of the inverse quark propagator in this basis reads \\begin{equation} S_0^{-1} = \\left( \\begin{array}{cc} \\left[ G_0^+ \\right]^{-1} & 0 \\\\ 0 & \\left[ G_0^- \\right]^{-1} \\end{array} \\right), \\end{equation} where \\begin{equation} \\left[ G_0^\\pm \\right]^{-1} = \\gamma^\\mu K_\\mu \\pm \\mu \\gamma_0 \\end{equation} are the inverse Dirac propagators for massless quarks ($G_0^{+}$) and charge-conjugate quarks ($G_0^{-}$).  In the ground state with Cooper pairing, see Eq.~(\\ref{2sc-cond}), the quark propagator in the Nambu-Gorkov basis also has nonzero off-diagonal components, \\begin{equation} S^{-1} = \\left( \\begin{array}{cc} \\left[ G_0^+ \\right]^{-1} & \\Delta^{-} \\\\ \\Delta^{+} & \\left[ G_0^- \\right]^{-1} \\end{array} \\right), \\label{inv_prop} \\end{equation} where $\\Delta^{-} = -i \\epsilon^{3}\\varepsilon\\gamma^5 \\Delta$ and $\\Delta^{+} \\equiv \\gamma^0 \\left(\\Delta^{-}\\right)^{\\dagger} \\gamma^0 = -i \\epsilon^{3}\\varepsilon\\gamma^5\\Delta^{*}$ are the matrices in the Dirac space, and $\\Delta$ is the diquark gap parameter. This is the simplest ansatz for the quark propagator in the color superconducting ground state. It can be made more general, for example, by adding a regular (diagonal) part of the quark self-energy. Such a self-energy plays an essential role in the dynamics of chiral symmetry breaking. In dense quark matter, on the other hand, it is not so important. Thus, for simplicity, it is neglected here (the effects of the diagonal part of the self-energy were studied in Ref.~\\cite{self-e}).  From the inverse propagator in Eq.~(\\ref{inv_prop}), one can derive the following expression for the quark propagator: \\begin{equation} S = \\left(\\begin{array}{cc} G^{+} & \\Xi^{-} \\\\ \\Xi^{+} & G^{-} \\end{array}\\right), \\label{quarkpropagator} \\end{equation} where \\begin{eqnarray} G^{\\pm} &=& \\left[\\left(G_0^{\\pm}\\right)^{-1} -\\Delta^{\\mp}G_0^{\\mp}\\Delta^{\\pm}\\right]^{-1}, \\\\ \\Xi^{\\pm} &=& -G_0^{\\mp} \\Delta^{\\pm} G^{\\pm} . \\end{eqnarray} This structure of the quark propagator carries basic information about the properties of the 2SC ground state. For instance, the location of its poles in the energy plane gives the quasiparticle dispersion relations.  The only quantity that remains unknown in the ansatz (\\ref{inv_prop}) is the value of the gap parameter $\\Delta$. This should be determined from a self-consistent gap equation that takes into account the pairing dynamics between quarks. This is discussed in more detail in the next three subsections.  \\subsubsection{One-gluon exchange interaction}  In dense QCD, in which quarks are weakly interacting, the one-gluon exchange is the dominant interaction between quarks. It appears, however, that the one-gluon interaction is partially screened by surrounding dense medium. Thus, before studying the pairing dynamics between quarks, we first address the modification of gluon properties due to dense medium.  The inverse propagator of the medium modified gluon can be written in the following form: \\begin{equation} \\left({\\cal D}^{-1}\\right)_{\\mu\\nu}^{AB} = \\left(D_0^{-1}\\right)_{\\mu\\nu}^{AB} +\\Pi_{\\mu\\nu}^{AB}, \\label{glu_prop} \\end{equation} where $\\Pi_{\\mu\\nu}^{AB}$ is the gluon self-energy, or the polarization tensor. The gluons with soft momenta, $p\\lesssim g_s\\mu$, play the key role in Cooper pairing of quarks. The main contribution to the corresponding soft gluon polarization tensor comes from the quark loop with hard internal momenta of order $\\mu$. This contribution is large compared to all other contributions in QCD (e.g., those from the ghost and the gluon loops) because it is proportional to the density of quark states at the Fermi surface (as well as to the coupling constant $\\alpha_s$), i.e., $\\Pi_{\\mu\\nu}^{AB}\\sim \\alpha_s\\mu^2$. The corresponding approximation, in which the sub-leading terms suppressed by powers of the quark chemical potential $\\mu$ are neglected, is called the hard dense loop (HDL) approximation.  The calculation of the polarization tensor in the HDL approximation was performed in Refs.~\\cite{Vija,Manuel} (in the case of dense QCD in $2+1$ dimensions, similar calculations were performed in Ref.~\\cite{qcd2+1}). Here, we present the final result and briefly comment on its main features.  The HDL polarization tensor can be written as follows: $\\Pi_{\\mu\\nu}^{AB}\\equiv \\delta^{AB}\\Pi_{\\mu\\nu}$, where the explicit form of the components of $\\Pi_{\\mu\\nu}$ are given by \\begin{eqnarray} \\Pi^{00}(p_0, \\mathbf{p}) & = & \\Pi_{l}(p_0, \\mathbf{p}) , \\label{Pi-00}\\\\ \\Pi^{0i}(p_0, \\mathbf{p}) & = & p_0 \\frac{p^i}{p^2} \\Pi_{l} (p_0, \\mathbf{p}) , \\label{Pi-0i}\\\\ \\Pi^{ij}(p_0, \\mathbf{p}) & = & \\left ( \\delta^{ij}- \\frac{p^i p^j}{p^2} \\right) \\Pi_{t} (p_0,\\mathbf{p})+ \\frac{p^i p^j} {p^2} \\frac{p_0^2}{p^2} \\Pi_{l} (p_0, \\mathbf{p}), \\label{Pi-ij} \\end{eqnarray} with \\begin{eqnarray} \\Pi_{l}(p_0,\\mathbf{p})&=&m_D^2\\left(\\frac{p_0}{2p} \\ln\\left|\\frac{p_0+p}{p_0-p}\\right|-1 -i\\pi\\frac{p_0}{2p}\\theta(-p^2) \\right),\\\\ \\Pi_{t}(p_0,\\mathbf{p})&=&\\frac{1}{2}m_D^2-\\frac{P^2}{2p^2} \\Pi_{l}(p_0,\\mathbf{p}). \\end{eqnarray} Here, we use the notation $p^2\\equiv \\mathbf{p}^2$, $P^2=p_0^2-p^2$, and $m_D^2\\equiv 2N_f\\alpha_{s}\\mu^2/\\pi$. It is remarkable that this result coincides with the polarization tensor derived in the framework of the {\\em classical} transport theory of dense Yang-Mills plasma \\cite{Heinz}. One can check that the HDL polarization tensor is transverse, i.e., \\begin{equation} P^{\\mu} \\Pi_{\\mu\\nu}(p_0, \\mathbf{p})=0. \\end{equation} It is convenient to write the tensor $\\Pi_{\\mu\\nu}$ in terms of two transverse projection operators, \\begin{equation} \\Pi_{\\mu\\nu}=-O^{(1)}_{\\mu\\nu} \\Pi_{t} +O^{(2)}_{\\mu\\nu} \\left(2\\Pi_{t}-m_D^2\\right). \\end{equation} where \\begin{eqnarray} O^{(1)}_{\\mu\\nu}&=& g_{\\mu\\nu}-u_{\\mu} u_{\\nu} +\\frac{\\mathbf{p}_{\\mu}\\mathbf{p}_{\\nu}}{p^{2}}, \\label{def-O1} \\\\ O^{(2)}_{\\mu\\nu}&=& u_{\\mu} u_{\\nu} -\\frac{\\mathbf{p}_{\\mu}\\mathbf{p}_{\\nu}}{p^{2}} -\\frac{P_{\\mu}P_{\\nu}}{P^{2}}, \\label{def-O2} \\end{eqnarray} and $u_{\\mu}=(1,0,0,0)$ is the time-like four-vector that defines the rest frame of dense quark medium. In order to have a complete set of projectors such that $\\sum_I O^{(I)}_{\\mu\\nu} =g_{\\mu\\nu}$, one has to include also the following longitudinal projector: \\begin{equation} O^{(3)}_{\\mu\\nu}= \\frac{P_{\\mu}P_{\\nu}}{P^{2}}. \\label{def-O3} \\end{equation} By making use of this set of three projectors, inverse gluon propagator in Eq.~(\\ref{glu_prop}) can be represented as follows: \\begin{equation} \\left({\\cal D}^{-1}\\right)_{\\mu\\nu}^{AB} =i\\delta^{AB} \\left[\\left(P^2-\\Pi_{t}\\right) O^{(1)}_{\\mu\\nu} +\\left(P^2+2\\Pi_{t}-m_D^2\\right) O^{(2)}_{\\mu\\nu} +\\frac{P^2}{\\lambda} O^{(3)}_{\\mu\\nu}\\right], \\label{inv-D} \\end{equation} where $\\lambda$ is the gauge fixing parameter. This expression can be easily inverted, leading to the following result for the gluon propagator: \\begin{equation} {\\cal D}_{\\mu\\nu}^{AB} =-i\\delta^{AB}\\left( \\frac{1}{P^2-\\Pi_{t}} O^{(1)}_{\\mu\\nu} +\\frac{1}{P^2+2\\Pi_{t}-m_D^2} O^{(2)}_{\\mu\\nu} +\\frac{\\lambda}{P^2} O^{(3)}_{\\mu\\nu}\\right). \\label{D} \\end{equation} The gluon modes which are transverse and longitudinal with respect to the {\\em three-momentum} $\\mathbf{p}$ are called the magnetic and electric modes, respectively. These two types of physical gluon modes are described by the following fields: \\begin{eqnarray} a^{\\rm (mag)}_{\\mu}&=& O^{(1)}_{\\mu\\nu}A^{\\nu}, \\label{def-mag} \\\\ a^{\\rm (el)}_{\\mu}&=&O^{(2)}_{\\mu\\nu}A^{\\nu} . \\label{def-el} \\end{eqnarray} From Eq.~(\\ref{D}) one can see that the propagators of these modes are naturally separated from each other, as well as from the unphysical longitudinal mode, described by the field $a^{\\rm (\\parallel)}_{\\mu} =O^{(3)}_{\\mu\\nu}A^{\\nu}$.  In the study of the Cooper pairing dynamics, the most relevant part of the one-gluon exchange interaction is the interaction with the space-like momenta such that $p_0\\ll p$. This should be clear from the kinematics of the quark scattering around the Fermi surface: the energy exchange between pairing quarks is typically much smaller than the change of their momenta. It is justified, therefore, to use the approximate form of $\\Pi_{t}$, obtained in the region $p_0\\ll p$, \\begin{equation} \\Pi_{t}\\simeq m_D^2\\left(i\\pi \\frac{p_0}{4p} \\theta(P^2)+\\frac{p_0^2}{p^2}+\\dots\\right). \\label{asym-2} \\end{equation} By substituting this asymptote into the gluon propagator (\\ref{D}), one can check that the magnetic modes produce a long-range (dynamically screened) interaction, while the electric mode works only at short distances, $r \\lesssim 1/m_D$.  In the next subsection, the HDL gluon propagator is used in the study of the Cooper pairing dynamics. As we shall see there, the long-range interaction mediated by the magnetic modes will play a particularly important role in the dynamics.  \\subsubsection{Gap equation in QCD}  Color superconductivity is an essentially non-perturbative phenomenon which cannot be addressed with perturbative techniques even if the theory is weakly interacting. The standard method for studying superconductivity is the method of the Schwinger-Dyson equations. In general, the complete set of Schwinger-Dyson equations contains an infinite number of coupled equations for Green's functions. There exist no systematic ways of solving such equations exactly. However, there exist various standard approximations that can be justified if the dynamics is controlled by a small parameter (e.g., a weak coupling constant, or a small $1/N$ parameter where $N$ is the number of flavors or colors in the model).  In the case of QCD at asymptotic densities, the main features of its dynamics are well described by the so-called improved rainbow approximation. In this approximation, one uses the bare quark-gluon vertices in the Schwinger-Dyson equation for the quark propagator. One also uses the gluon propagator with the screening effects taken into account in the HDL approximation. (Note that, in the simple rainbow approximation, the screening is usually neglected.) The graphical representation of the Schwinger-Dyson equation is shown in Fig.~\\ref{fig-sd-eq}. \\begin{figure}[ht] \\begin{center} \\includegraphics{SD11-equations.eps} \\caption{The graphical representation of the Schwinger-Dyson equation in the improved rainbow approximation. The quark propagator in the Nambu-Gorkov basis and the gluon propagator are denoted by the solid lines and the wavy lines, respectively.} \\label{fig-sd-eq} \\end{center} \\end{figure} The analytical form of this equation reads \\begin{equation} S^{-1}(K) = S_0^{-1}(K) +4\\pi\\alpha_s \\int\\frac{d^4 p}{(2\\pi)^4} \\Gamma^{A}_{\\mu} S(P) \\Gamma^{B}_{\\nu} {\\cal D}_{AB}^{\\mu\\nu}(K-P), \\label{sd-eq} \\end{equation} where the quark-gluon vertices in the Nambu-Gorkov basis are given by \\begin{equation} \\Gamma^{A}_{\\mu} = \\gamma_\\mu \\left( \\begin{array}{cc} T^A & 0 \\\\ 0 & -\\left(T^A\\right)^T \\end{array} \\right). \\end{equation} In the case of the simplest ansatz for the quark propagator, as in Eq.~(\\ref{inv_prop}), the Schwinger-Dyson equation (\\ref{sd-eq}) reduces to an equation for the gap parameter $\\Delta$. For this reason, it is also called the gap equation. After neglecting the dependence of the gap on the three-momentum and after performing the momentum integration, one arrives at the following approximate form of the gap equation in Euclidean space ($k_4\\equiv ik_0$) \\cite{weak-son,weak,weak-sp1}: \\begin{equation} \\Delta(k_4) \\simeq \\frac{\\alpha_s}{9\\pi} \\int \\frac{dp_4 \\Delta(p_4)} {\\sqrt{p_4^2+\\Delta^2}}\\ln\\frac{\\Lambda}{|k_4-p_4|}, \\label{gap_eq-qcd} \\end{equation} where $\\Lambda = 2 (4\\pi)^{3/2}\\mu \\alpha_s^{-5/2}$. If the quark wave function renormalization from the diagonal part of the self-energy were taken into account, the extra factor $\\exp\\left(-\\frac{4+\\pi^2}{8}\\right)$ would appear in the expression for $\\Lambda$ \\cite{self-e}. The appearance of the logarithm in the integrand of the gap equation (\\ref{gap_eq-qcd}) is an artifact of the long-range force in QCD, mediated by the magnetic gluon modes. An approximate solution to the gap equation reads \\begin{equation} \\Delta(0) \\simeq \\Lambda \\exp\\left(-\\frac{3\\pi^{3/2}}{2^{3/2} \\sqrt{\\alpha_s}}\\right). \\label{gap-sol} \\end{equation} The result is plotted in Fig.~\\ref{fig-gap-sol}. Here, the value of the coupling constant $\\alpha_s\\equiv \\alpha_s(\\mu)$ is taken at the scale of the quark chemical potential. \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.5\\textwidth]{gap_vs_mu1.eps} \\caption{The value of the gap from Eq.~(\\ref{gap-sol}) as a function of the quark chemical potential. The coupling constant is defined at the scale of the quark chemical potential. The solid (dashed) line gives the result with (without) taking the quark wave function renormalization into account.} \\label{fig-gap-sol} \\end{center} \\end{figure}  It appears that, at densities that exist inside stars, i.e., $n\\lesssim 10n_0$, the value of the quark chemical potential cannot be much larger than $500\\,\\mbox{MeV}$. The corresponding value of the QCD coupling constant $\\alpha_s\\equiv \\alpha_s(\\mu)$ is not small, and, in contrast to the situation in asymptotically dense QCD, the Cooper pairing dynamics is not weakly interacting. This suggests that the above analysis of the Schwinger-Dyson equation may not be suitable for the description of matter inside stars. In spite of this, the solution in Eq.~(\\ref{gap-sol}) is of fundamental importance in theory. This is one of very few non-perturbative solutions in QCD which is obtained from first principles, and which can be systematically improved by studying higher order corrections. Also, this solution shows that the phase diagram of QCD in the plane of temperature and chemical potential contains color superconducting phases at least at asymptotic densities.  \\subsubsection{Gap equation in the NJL model}  If deconfined quark matter exists in the central regions of compact stars, it is likely to be color superconducting even when it is not weakly interacting. In order to study such a strongly interacting regime of quark matter with a typical value of the quark chemical potential $\\mu\\simeq 500\\,\\mbox{MeV}$, one can use various effective models of QCD. The simplest type of such models is the NJL model with a local four-fermion interaction in Eq.~(\\ref{NJL-action}).  The NJL model in Eq.~(\\ref{NJL-action}) lacks gluons, and, as reflection of this, it possesses the global instead of the local (gauge) SU(3)$_{c}$ color symmetry. This model can be viewed as a result of integrating out heavy gluons from the QCD action. This is possible if the gluons get nonzero masses from non-perturbative effects. One arrives at an effective NJL model similar to that in Eq.~(\\ref{NJL-action}) when the gauge fixing in QCD, needed to perform the integration, is consistent with the global color symmetry.  The global color symmetry of the NJL model is broken in the 2SC phase. Then, because of the Goldstone theorem, five massless Nambu-Goldstone bosons should appear in the low-energy spectrum of such a theory. In QCD, however, there is no room for such Nambu-Goldstone bosons. The seeming contradiction is removed by noting that these Nambu-Goldstone bosons are not physical. Their appearance is an artifact of the gauge fixing. In particular, there exist a gauge choice in QCD, namely the unitary gauge, in which these bosons can be completely eliminated.  The gap equation in the NJL model in the mean field approximation looks as follows: \\begin{equation} \\Delta \\simeq \\frac{4 G_D}{\\pi^2} \\int_0^{\\Lambda} \\left(\\frac{\\Delta}{\\sqrt{(p-\\mu)^2+\\Delta^2}} +\\frac{\\Delta}{\\sqrt{(p+\\mu)^2+\\Delta^2}}\\right) p^2 d p . \\label{gap-eq-NJL} \\end{equation} This gap equation can be obtained from a Schwinger-Dyson equation similar to that in QCD in Eq.~(\\ref{sd-eq}) after the gluon long-range interaction is replaced by a local interaction proportional to $\\delta(K-P)$.  The approximate solution to the gap equation in Eq.~(\\ref{gap-eq-NJL}) reads \\begin{equation} \\Delta \\simeq 2\\sqrt{\\Lambda^2-\\mu^2} \\exp\\left(-\\frac{\\pi^2}{8G_D\\mu^2}+\\frac{\\Lambda^2-3\\mu^2}{2\\mu^2}\\right). \\end{equation} This is very similar to the Bardeen-Cooper-Schrieffer solution in the case of low temperature superconductivity in solid state physics \\cite{BCS}. As in the Bardeen-Cooper-Schrieffer theory, it has the same type non-analytic dependence on the coupling constant and the same type dependence on the density of quasiparticle states at the Fermi surface.  When the quark chemical potential $\\mu$ takes a value in the range between $400\\,\\mbox{MeV}$ and $500\\,\\mbox{MeV}$, and the strength of the diquark pairing is $G_D=\\eta G_S$ with $\\eta$ between $0.7$ and $1$, the value of the gap appears to be of order $100\\,\\mbox{MeV}$. In essence, this is the result that was obtained in Ref.~\\cite{cs}.  \\subsubsection{Properties of quark matter in the 2SC phase}  Now let us briefly summarize the main physical properties of the 2SC phase of dense quark matter. To large extent, these follow directly from the symmetry of the ground state, determined by the structure of the condensate in Eq.~(\\ref{2sc-cond}), and the magnitude of the diquark gap $\\Delta$.  As we argued earlier, the blue up and blue down quarks do not participate in Cooper pairing in the 2SC phase. They give rise to gapless quasiparticles in the low energy spectrum of the theory. The density of states of such quasiparticles is proportional to $\\mu^2$ and, therefore, is very large. This simple fact has important implications for the quark matter properties. At small temperatures, these gapless quasiparticles give dominant contributions to the specific heat, as well as to the electrical and heat conductivities of the 2SC phase. Also, the presence of the ungapped blue up and blue down quarks should result in a large neutrino emissivity due to the $\\beta$-processes, $d_b \\to u_b +e^{-}+\\bar{\\nu}_{e}$ and $u_b +e^{-}\\to d_b +\\nu_{e}$ \\cite{Iwamoto}.  The other four quark quasiparticles, originating from the red and green paired quarks, are gapped. Their dispersion relations look like that in Eq.~(\\ref{E_qp}). At small temperatures, $T\\ll \\Delta$, the contributions of these quasiparticles to all transport and many thermodynamic quantities are suppressed by the exponentially small factor $\\exp\\left(-\\Delta/T\\right)$. Gluons do not play a very important role either. They are bosons and their thermal number densities are small at small temperatures. In addition, five out of total eight gluons are gapped because of the color Meissner effect.  The presence of the unpaired blue quarks is also connected with the absence of baryon superfluidity in the 2SC phase. The generator of the baryon number conservation symmetry is defined as follows: \\begin{equation} \\tilde{B}=B-\\frac{2}{\\sqrt{3}}T_8 =\\mbox{diag}_{\\rm color}\\left(0,0,1\\right), \\label{U1_B} \\end{equation} where $B$ is the generator of the U(1)$_B$ symmetry in vacuum. This vacuum generator $B$ mixes with the color generator $T_8$ to produce the generator $\\tilde{B}$ of the $\\tilde{\\mbox{U}}(1)_B$ symmetry in medium. From Eq.~(\\ref{U1_B}) it is clear that only (anti-)blue quasiparticles carry a nonzero baryon number in the 2SC phase.  In the 2SC phase, there is also an unbroken $\\tilde{\\mbox{U}}(1)_{\\rm em}$ gauge symmetry. The corresponding generator reads \\begin{equation} \\tilde{Q}=Q-\\frac{1}{\\sqrt{3}}T_8 , \\label{U1_em_2sc} \\end{equation} where $Q=\\mbox{diag}_{\\rm flavor}(\\frac{2}{3},-\\frac{1}{3})$ is the generator of the U(1)$_{\\rm em}$ symmetry in vacuum. The gauge boson of $\\tilde{\\mbox{U}}(1)_{\\rm em}$ is nothing else but the medium photon. It should be clear, therefore, that the 2SC ground state is not subject to the electromagnetic Meissner effect. Thus, if there is an external magnetic field, it is not expelled from such a medium.  One should notice that the quark quasiparticles carry the following $\\tilde{Q}$-charges: \\begin{eqnarray} \\tilde{Q}(u_r)=\\tilde{Q}(u_g)=\\frac{1}{2}, && \\tilde{Q}(u_b)=1,\\\\ \\tilde{Q}(d_r)=\\tilde{Q}(d_g)=-\\frac{1}{2}, && \\tilde{Q}(d_b)=0. \\end{eqnarray} It is interesting that the blue up quark, which remains unpaired in the 2SC phase, is charged (this is not true for the blue down quark which is also unpaired). It means that this phase of matter is a $\\tilde{Q}$-conductor, and that the value of the electrical conductivity is very high.  The pressure of cold dense quark matter is dominated by the Pauli pressure. The partial contribution of each quark is $\\mu^4/12\\pi^2$, see Eq.~(\\ref{pressure_0}). The correction due to color superconductivity is parametrically suppressed by the factor $\\left(\\Delta/\\mu\\right)^2$. The reason is that the Cooper pairing affects only those quark states that are close to the Fermi surface. In contrast, it is the whole Fermi sphere that contributes to the Pauli pressure. At weak coupling, the contribution due to diquark pairing can be easily calculated. It is given by $\\left(\\mu\\Delta/2\\pi\\right)^2$ per each gapped quark quasiparticle \\cite{eff-pot}. In the 2SC phase, there are six (two flavors times three colors) quarks in total, and four of them give rise to quasiparticles with the gap $\\Delta$ in their energy spectra. Thus, the pressure of such matter is approximately given by \\cite{eff-pot,kappa} \\begin{equation} P_{\\rm (2SC)}\\simeq \\frac{N_c N_f \\mu^4}{12\\pi^2} -B + 4 \\left(\\frac{\\mu\\Delta}{2\\pi}\\right)^{2} =\\frac{\\mu^4}{2\\pi^2} -B + \\frac{\\mu^2\\Delta^2}{\\pi^2}. \\label{pres-2sc} \\end{equation} Note that the bag pressure was also added on the right hand side. If present, this term partially cancels the leading $\\mu^4$ term. Because of this, color superconductivity may have a large effect on the hadron-quark phase transition \\cite{LH02} and even on the properties of compact stars \\cite{AlRe03}. By making use of thermodynamic identities, one can derive the corresponding energy density, \\begin{equation} \\epsilon_{\\rm (2SC)}\\simeq \\frac{3\\mu^4}{2\\pi^2} +B + \\frac{\\mu^2\\Delta^2}{\\pi^2} \\left(1+\\frac{2\\mu}{\\Delta}\\frac{\\partial\\Delta }{\\partial\\mu}\\right). \\label{ener-2sc} \\end{equation} The expressions in Eqs.~(\\ref{pres-2sc}) and (\\ref{ener-2sc}) give a parametric representation of the equation of state of dense quark matter in the 2SC phase. Such an equation of state is the key input in the Tolman-Oppenheimer-Volkoff equations \\cite{Tol} which determine the interior structure of compact stars. In short, such equations describe hydrostatic equilibrium of matter inside a star when the pressure gradient in each radial layer of the star is balanced by the gravitational weight of the layer itself. Several constructions of compact stars with quark matter in their interior were presented in Refs.~\\cite{LH02,AlRe03,BB02,BBBNOS02,SHH,BGAYT,BNOS}.  \\subsection{Color-flavor locked phase ($N_f=3$)}  It may happen that dense baryonic matter is made not only of the lightest up and down quarks, but of strange quarks as well. In fact, because of a possible reduction in the free energy from converting non-strange quarks into strange quarks, one may even speculate that strange quark matter is the true ground state of baryonic matter \\cite{Bod}.  The constituent strange quark mass in vacuum QCD is estimated to be of order $500\\,\\mbox{MeV}$. Its current mass is about $100\\,\\mbox{MeV}$. In dense baryonic matter (say, with $\\mu\\simeq 500\\,\\mbox{MeV}$), the value of the strange quark mass should be in the range between $100\\,\\mbox{MeV}$ and $500\\,\\mbox{MeV}$. It is plausible then that strange quarks also participate in Cooper pairing.  Here, we consider an idealized version of three-flavor quark matter, in which all quarks are assumed to be massless. The more realistic case of a nonzero strange quark mass will be briefly discussed in the second lecture. In the massless case, the quark model possesses the global SU(3)$_L\\times$SU(3)$_R$ chiral symmetry and the global U(1)$_B$ symmetry connected with the baryon number conservation. This is in addition to SU(3)$_c$ color gauge symmetry. Note that the generator $Q=\\mbox{diag}_{\\rm flavor}(\\frac{2}{3},-\\frac{1}{3}, -\\frac{1}{3})$ of the U(1)$_{\\rm em}$ symmetry of electromagnetism is traceless, and therefore it coincides with one of the vector-like generators of the SU(3)$_L\\times$SU(3)$_R$ chiral group.  To large extent, the color and flavor structure of the diquark condensate of Cooper pairs in the three-flavor quark matter is again determined by the symmetry of the attractive diquark channel (i.e., color-antisymmetric antitriplet) and the Pauli exclusion principle. The spin-0 condensate corresponds to the following ground state expectation value \\cite{cfl}: \\begin{equation} \\left\\langle \\left(\\bar{\\psi}^C\\right)_i^a \\gamma^5 \\psi_j^b \\right\\rangle \\sim \\sum_{I,J=1}^{3} c^{I}_{J}\\varepsilon_{ijI} \\epsilon^{abJ} +\\cdots, \\label{cfl-cond} \\end{equation} which is antisymmetric in the color and flavor indices of the constituent quarks, cf. Eq.~(\\ref{2sc-cond}). The $3\\times 3$ matrix $c^{I}_{J}$ is determined from the global minimum of the free energy. It appears that $c^{I}_{J}=\\delta^{I}_{J}$. The ellipsis on the right hand side stand for a contribution which is symmetric in color and flavor. A small contribution of this type is always induced in the ground state, despite the fact that it corresponds to a repulsive diquark channel \\cite{cfl,weak-cfl}. This is not surprising after noting that the symmetric condensate $\\left\\langle \\left(\\bar{\\psi}^C\\right)_i^a \\gamma^5 \\psi_j^b \\right\\rangle \\sim \\delta_{i}^{a}\\delta_{j}^{b}+\\delta_{j}^{a}\\delta_{i}^{b}$ does not break any additional symmetries \\cite{cfl}.  In the ground state, determined by the condensate (\\ref{cfl-cond}), the chiral symmetry is broken down to its vector-like subgroup. However, the mechanism of this symmetry breaking is very unusual. To see this, let us rewrite the condensate in Eq.~(\\ref{cfl-cond}) as follows: \\begin{equation} \\left\\langle \\psi_{L,i}^{a,\\alpha} \\epsilon_{\\alpha\\beta} \\psi_{L,j}^{b,\\beta}\\right\\rangle =-\\left\\langle \\psi_{R,i}^{a,\\dot\\alpha} \\epsilon_{\\dot\\alpha\\dot\\beta} \\psi_{R,j}^{b,\\dot\\beta}\\right\\rangle \\sim \\sum_{I=1}^{3} \\varepsilon_{ijI} \\epsilon^{abI} +\\cdots, \\label{LL-RR-cond} \\end{equation} where $\\alpha,\\beta,\\dot\\alpha,\\dot\\beta=1,2$ are the spinor indices. The condensate of left-handed fields in Eq.~(\\ref{LL-RR-cond}) breaks the SU(3)$_c$ color symmetry and the SU(3)$_L$ chiral symmetry, but leaves the diagonal SU(3)$_{L+c}$ subgroup unbroken. Indeed, as one can check, this condensate remains invariant under the simultaneous flavor transformation $g_L$ and the compensating color transformation $g_c=(g_L)^{-1}$. Similarly, the condensate of right-handed fields in Eq.~(\\ref{LL-RR-cond}) leaves the SU(3)$_{R+c}$ subgroup unbroken.  When both condensates are present, the symmetry of the ground state is SU(3)$_{L+R+c}$. At the level of global symmetries, the original SU(3)$_{L}\\times$SU(3)$_{R}$ is broken down to the vector-like SU(3)$_{L+R}$, just like in vacuum. Unlike in vacuum, however, this breaking does {\\em not} result from any condensates mixing left- and right-handed fields. Instead, it results from two separate condensates, made of left-handed fields only and of right-handed fields only. The color-flavor orientations of the two condensates are ``locked'' to each other by color transformations. This mechanism is called locking, and the corresponding phase of matter is called color-flavor-locked (CFL) phase \\cite{cfl}.  To large extent, the gap equation in the three-flavor quark matter is the same as in the two-flavor case. The differences come only from a slightly more complicated color-flavor structure of the off-diagonal part of the inverse quark propagator (gap matrix) \\cite{cfl,weak-cfl}, \\begin{equation} \\Delta^{ij}_{ab}=i\\gamma^5\\left[ \\frac{1}{3}\\left(\\Delta_1+\\Delta_2\\right)\\delta^{i}_{a}\\delta^{j}_{b} -\\Delta_2\\delta^{i}_{b}\\delta^{j}_{a}\\right], \\label{delta_cfl} \\end{equation} where two parameters $\\Delta_1$ and $\\Delta_2$ determine the values of the gaps in the quasiparticles spectra. In the ground state that respects the SU(3)$_{L+R+c}$ symmetry, the original nine quark states give rise to a singlet and an octet of quasiparticles. The singlet has gap $\\Delta_1$, and the octet has gap $\\Delta_2$.  When the small symmetric diquark condensate is neglected, one finds that $\\Delta_1=2\\Delta_2$, i.e., the gap of the singlet is twice as large as the gap of the octet. In real case, however, this relation holds only approximately. The value of the octet gap $\\Delta_2$ in asymptotically dense QCD is given by an expression similar to that in Eq.~(\\ref{gap-sol}), but the overall coefficient is multiplied by an extra factor $2^{-1/3}\\left(2/3\\right)^{5/2}\\approx 0.288$ \\cite{weak-cfl,loops}.  \\subsubsection{Properties of quark matter in the CFL phase}  Let us briefly review physical properties of the CFL phase. In contrast to the 2SC phase, there are no gapless quark quasiparticles in the low energy spectrum of the CFL phase. This implies that, at small temperatures, $T\\ll \\Delta$, the contributions of quark quasiparticles to all transport and many thermodynamic quantities are suppressed by the exponentially small factor $\\exp\\left(-\\Delta/T\\right)$. Gluons do not play any important role either. All of them are gapped because of the color Meissner effect.  Unlike the 2SC phase, the CFL phase is superfluid. This is because the U(1)$_B$ baryon number symmetry is broken in the ground state. If such a phase appears in a core of a rotating star, it will be threaded with rotational vortices. This may be related to the existence of such phenomena as glitches. These are sudden changes of the rotational frequency observed in some pulsars. The glitches may be caused by occasional releasing of the angular momentum of rotational vortices, that remained pinned to the stellar crust \\cite{gl-super}.  Because of the Goldstone theorem, the breaking of the U(1)$_B$ baryon number symmetry should result in the appearance of a Nambu-Goldstone boson in the low-energy theory. In absence of gapless quark quasiparticles, this Nambu-Goldstone boson turns out to play an important role in many transport properties of cold CFL matter \\cite{kappa,opaque}.  The CFL phase, like the 2SC phase, is not an electromagnetic superconductor. Therefore, it does not expel a magnetic flux from its interior. This is the result of having an unbroken $\\tilde{\\mbox{U}}(1)_{\\rm em}$ gauge symmetry in the ground state. The corresponding generator reads \\begin{equation} \\tilde{Q}=Q-T_3-\\frac{1}{\\sqrt{3}}T_8 . \\label{U1_em_cfl} \\end{equation} The quark quasiparticles carry the following $\\tilde{Q}$-charges: \\begin{eqnarray} \\tilde{Q}(u_r)=0, && \\tilde{Q}(u_g)=\\tilde{Q}(u_b)=1,\\\\ \\tilde{Q}(d_r)=-1, && \\tilde{Q}(d_g)=\\tilde{Q}(d_b)=0,\\\\ \\tilde{Q}(s_r)=-1, && \\tilde{Q}(s_g)=\\tilde{Q}(s_b)=0. \\end{eqnarray} Since all quasiparticles are gapped, the corresponding phase of matter is a $\\tilde{Q}$-insulator. Its electrical conductivity at small temperatures is dominated by thermally excited electrons and positrons \\cite{opaque}. At sufficiently low temperatures, when the thermal density is low, the CFL phase becomes transparent to light \\cite{opaque,opaque2}.  In order to write down the expression for the pressure of three-flavor quark matter in the CFL phase, we take the Pauli pressure contributions of nine (three flavors times three colors) quarks and add the correction due to color superconductivity from one quasiparticle with gap $\\Delta_1$ and eight quasiparticles with gap $\\Delta_2$. Thus, we derive \\begin{eqnarray} P_{\\rm (CFL)} &\\simeq &\\frac{N_c N_f \\mu^4}{12\\pi^2} -B + \\left(\\frac{\\mu\\Delta_1}{2\\pi}\\right)^{2} + 8\\left(\\frac{\\mu\\Delta_2}{2\\pi}\\right)^{2} \\nonumber \\\\ &\\simeq& \\frac{3\\mu^4}{4\\pi^2} -B + 3\\frac{\\mu^2\\Delta^2}{\\pi^2}, \\label{pres-cfl} \\end{eqnarray} where we used the approximate relation between the singlet and octet gaps, $\\Delta_1=2\\Delta_2\\equiv 2\\Delta$. By making use of thermodynamic identities, one can derive the corresponding energy density, \\begin{equation} \\epsilon_{\\rm (CFL)}\\simeq \\frac{9\\mu^4}{4\\pi^2} +B + 3\\frac{\\mu^2\\Delta^2}{\\pi^2} \\left(1+\\frac{2\\mu}{\\Delta}\\frac{\\partial\\Delta }{\\partial\\mu}\\right). \\label{ener-cfl} \\end{equation} These two expressions give a parametric representation of the equation of state of dense quark matter in the CFL phase.  \\subsubsection{Low-energy effective action}  As in the case of chiral perturbation theory in vacuum QCD \\cite{GasLeut}, one could write down an effective low-energy theory in the CFL phase. From the symmetry breaking pattern, it is known that there are nine Nambu-Goldstone bosons and one pseudo-Nambu-Goldstone boson in the low-energy spectrum of the theory \\cite{HRZ,CasGat}. Eight of the Nambu-Goldstone bosons are similar to those in vacuum QCD: three pions ($\\pi^0$ and $\\pi^\\pm$), four kaons ($K^0$, $\\bar{K}^0$ and $K^\\pm$) and the eta-meson ($\\eta$). The additional Nambu-Goldstone boson ($\\phi$) comes from breaking the U(1)$_B$ baryon symmetry. Finally, the pseudo-Nambu-Goldstone boson ($\\eta^\\prime$) results from breaking the approximate axial U(1)$_A$ symmetry.  The low energy action for the Nambu-Goldstone bosons in the CFL phase was derived in Refs.~\\cite{HRZ,CasGat,SonSt}. In the chiral limit, the result reads \\begin{eqnarray} {\\cal L}_{eff} &=& \\frac{f_\\pi^2}{4} {\\rm Tr}\\left[ \\partial_0\\Sigma\\partial_0\\Sigma^\\dagger - v^2 \\partial_i\\Sigma\\partial_i\\Sigma^\\dagger \\right] + \\frac{1}{2} \\left[ (\\partial_0 \\phi)^{2} - v^2 (\\partial_i\\phi)^{2} \\right] \\nonumber \\\\ &+&\\frac{1}{2} \\left[ (\\partial_0 \\eta^{\\prime})^{2} - v^2 (\\partial_i\\eta^{\\prime})^{2} \\right], \\label{cov-der} \\end{eqnarray} where $f^2_\\pi=(21-8\\ln 2)(\\mu/6\\pi)^2$ and $v^2=1/3$ were calculated in asymptotically dense QCD in Refs.~\\cite{SonSt,fpi-more}. (A nonzero mass of the $\\eta^\\prime$-meson was neglected here \\cite{SSZh}.) By definition, $\\Sigma\\equiv \\exp \\left(i\\lambda^A \\pi^A/f_\\pi\\right)$ is a unitary matrix field which describes the octet of the Nambu-Goldstone bosons, transforming under the chiral $SU(3)_{L} \\times SU(3)_{R}$ group as follows: \\begin{equation} \\Sigma \\to U_{L} \\Sigma U_{R}^{\\dagger}, \\end{equation} where $(U_{L},U_{R}) \\in SU(3)_{L} \\times SU(3)_{R}$.  The extra terms in the low-energy action should be added when quark masses are nonzero \\cite{SonSt,mmm}. Since the quark masses break the chiral symmetry of dense QCD explicitly, these type of corrections produce non-vanishing masses for all Nambu-Goldstone bosons, except one. The single boson left massless is the Nambu-Goldstone boson related to the baryon number breaking. Of course, this symmetry is not affected by the quark masses.  Within the framework of the high density effective theory \\cite{loops,hdet}, it was shown that an additional effect of nonzero quark masses may appear when the strange quark mass exceeds a critical value $m_s\\sim m_u^{1/3}\\Delta^{2/3}$. In this case, the CFL phase is expected to undergo a phase transition to a phase with a meson (e.g., kaon or/and eta) condensate \\cite{BS,KR,Kry2004cw}. In the low-energy action of the corresponding phase, additional Nambu-Goldstone bosons with very unusual properties can appear \\cite{abnormal}.  \\subsection{Spin-1 color superconductivity ($N_f=1$)}  In the case of neutral matter in $\\beta$ equilibrium, as we shall discuss in the second lecture in more detail, it may happen that different quark flavors cannot create Cooper pairs because of a large mismatch between their Fermi momenta. In this case, one could consider the possibility of a much weaker spin-1 Cooper pairing \\cite{Bail,weak-sp1,sp1995,spin-1,spin-1-Meissner,andreas}. It is the Pauli principle that does not allow to construct spin-0 Cooper pairs from quarks of the same flavor. The color antisymmetric wave function of a pair can only be symmetric in spin indices. This corresponds to a spin-1 state.  The pairing in a spin-1 triplet state and in a color antisymmetric antitriplet state can lead to a rather complicated structure of the diquark condensate. The general structure of the gap matrix can be written in the following form \\cite{andreas}: \\begin{equation} \\Delta^{ab}=i\\Delta_0 \\sum_{c,i=1}^{3} \\epsilon^{abc} {\\cal C}_{ci} \\left[\\hat{k}^i\\cos\\theta +\\gamma^i_\\perp\\sin\\theta\\right], \\label{delta_1sc} \\end{equation} where $\\hat{\\mathbf{k}}\\equiv \\mathbf{k}/k$, $\\gamma^i_\\perp\\equiv \\gamma^i-\\hat{k}^i (\\mbox{\\boldmath $\\gamma$}\\cdot\\hat{\\mathbf{k}})$, and the explicit form of the $3\\times 3$ matrix ${\\cal C}_{ci}$ as well as the angular parameter $\\theta$ are determined by the minimization of the quark matter free energy. In the special cases with $\\theta=0$ and $\\theta=\\pi/2$, the corresponding gaps are called longitudinal and transverse, respectively.  By choosing various specific matrices ${\\cal C}$, one could construct many different spin-1 phases. The four most popular of them are the so-called A-phase, the color-spin-locked (CSL) phase, the polar and the planar phases. These are determined by the following matrices \\cite{andreas}: \\begin{eqnarray} {\\cal C}^{\\rm (A-phase)} =\\frac{1}{\\sqrt{2}}\\left(\\begin{array}{lll} 0 & 0 & 0 \\\\ 0 & 0 & 0 \\\\ 1 & i & 0 \\end{array}\\right), &\\quad & {\\cal C}^{\\rm (CSL)} =\\frac{1}{\\sqrt{3}}\\left(\\begin{array}{lll} 1 & 0 & 0 \\\\ 0 & 1 & 0 \\\\ 0 & 0 & 1 \\end{array}\\right), \\\\ {\\cal C}^{\\rm (polar)} =\\left(\\begin{array}{lll} 0 & 0 & 0 \\\\ 0 & 0 & 0 \\\\ 0 & 0 & 1 \\end{array}\\right),&\\quad & {\\cal C}^{\\rm (planar)} =\\frac{1}{\\sqrt{2}}\\left(\\begin{array}{lll} 1 & 0 & 0 \\\\ 0 & 1 & 0 \\\\ 0 & 0 & 0 \\end{array}\\right). \\end{eqnarray} The corresponding four phases are characterized by different symmetries of their ground state. In particular, the original group SU(3)$_c\\times $SO(3)$_J\\times $U(1)$_{\\rm em}$ of one-flavor quark matter breaks down as follows \\cite{spin-1,spin-1-Meissner,andreas}: \\begin{eqnarray} \\mbox{A-phase:} & & SU(2)_c\\times \\widetilde{SO}(2)_J \\times \\tilde{U}(1)_{\\rm em},\\\\ \\mbox{CSL:} & & \\widetilde{SO}(3)_{J},\\label{csl}\\\\ \\mbox{Polar:} & & SU(2)_c\\times SO(2)_J \\times \\tilde{U}(1)_{\\rm em},\\\\ \\mbox{Planar:} & & \\widetilde{SO}(2)_J \\times \\tilde{U}(1)_{\\rm em}. \\end{eqnarray} It was argued recently that the phase with the lowest free energy is the transverse A-phase \\cite{andreas}.  The value of the spin-1 gap is estimated to be about two or three orders of magnitude smaller than a typical spin-0 gap. At best, it can be about $1\\,\\mbox{MeV}$, but more realistically it is expected to be about $0.1\\,\\mbox{MeV}$, or less. The presence of such a small gap is unlikely to modify considerably any thermodynamic, or even transport properties of dense quark matter. It is fair to mention, however, that the cooling history of stars might be sensitive enough to feel the effect of such small gaps \\cite{B_cool}. It was also suggested that the electromagnetic properties of spin-1 phases can be of phenomenological importance \\cite{spin-1-Meissner}. In contrast to the 2SC and CFL phases, discussed earlier, spin-1 phases reveal the electromagnetic Meissner effect. As one can see from Eq.~(\\ref{csl}), this clearly applies to the CSL phase. It turns out that the electromagnetism is also subject to the Meissner effect when there appears a mixture of any two independent spin-1 condensates made of quarks with different charges (e.g., one is made of up quarks and the other is made of down quarks) \\cite{spin-1-Meissner}. The type-I superconductivity in such a system could have observable effects on the magnetic field relaxation in pulsars.  \\subsection{Summary of the first lecture}  In this lecture, only the main three color superconducting phases of dense quark matter have been discussed. These are the simplest possible phases that can be realized in one-, two-, and three-flavor quark matter. It should be emphasized, however, that the conditions in dense matter were idealized by assuming that the chemical potentials of pairing quarks were equal. In real situation which, for example, is realized in the interior of compact stars, matter happens to be very different from ideal. In particular, the charge neutrality and the $\\beta$ equilibrium are two very important conditions that may modify the quark chemical potentials and, therefore, affect the properties of quark matter. These issues are discussed in the second lecture.  The aim of this lecture was to introduce the general idea of color superconductivity in cold dense baryonic matter. It was argued that such matter is expected to be deconfined and, because of the Cooper instability, it should be color superconducting. To large extent, this phenomenon is the same as low-temperature superconductivity in condensed matter physics \\cite{BCS}. One feature in dense QCD is very remarkable. The attractive interaction between quarks comes from the gauge boson exchange. This is in contrast to the Bardeen-Cooper-Schrieffer theory in solid state physics, where the gauge boson (photon) exchange is repulsive, while an attractive interaction comes from the phonon exchange.  One of the key points of this lecture is the observation that QCD at asymptotic densities is a weakly interacting (although non-perturbative) theory. Moreover, the properties of its color superconducting ground state can be studied analytically from first principles, providing a rare example of an essentially solvable limit in a non-Abelian theory. By itself, this has a fundamental theoretical importance. Also, this result may provide valuable insights in the theory of strong interactions. One of the examples might be the idea of duality between the hadronic and quark description of QCD \\cite{cont}.  Although the use of weakly interacting QCD for the description of dense quark matter is very instructive by itself, it does not seem to be very useful for the quantitative description of the ground state at densities existing inside stars. The reason is that the density of matter in stars is bounded from above by the condition of the hydrostatic equilibrium, $n\\lesssim 10n_0$. At such not so large densities, one cannot treat QCD as a weakly interacting theory. In this situation, the use of various effective models of QCD has proved to be very useful.  \\newpage ",
        "conclusions": ""
    },
    "0410/astro-ph0410100_arXiv.txt": {
        "abstract": "We combine all published NIR  and  optical photometrical observations of NGC 4151 as well as our new unpublished yet data which can be used for determination of time delays between optical and NIR variations. Before we have found change of time delay value for variations in $K$ and $L$ filters for different states of the luminosity. Here we are considering new data for deep minimum after the very high state of the nucleus. We conclude that dust recovering time after  the high state is for any case more then several years. ",
        "introduction": "The review of past results on the optical-NIR time delay  investigations  for this object as well as for other AGNs can be found in Oknyanskij (2002), Oknyankij and Horne (2001). Here we take into account and discuss new result for NGC 4151  Minezaki et al. (2004).  In past  papers  (Oknyanskij, 1993, Oknyanskij et al., 1999) has been found time delay values between optical UBV  and NIR (JHKL)  variations  for two different luminosity cyrcles of NGC~4151 activities (see Fig.1): Case A (1969-1980) and Case B (1990-1998). We have found that time delay  is longer for Case B following to significantly more high level of luminosity in this cyrcle comparative to the past case A.  We  are  going  to  report  about the new time delay  determinations using  optical  UBV   and   NIR data for NGC 4151  in the minimum after high state B (then we call interval 1999-2004 as  Case C).           The past NIR time delays  for NGC~4151  at Cases A and B as well as new ones for Case C (this work and Minezaki at al. 2004) are collected in Table 1.     \\begin{table} \\caption{Collects optical--infrared time delays measured in NGC~4151 }  \\begin{tabular}{lcccccl} \\\\ \\tableline object  & delay         & bands &         reference \\\\      & (d)           &       & \\\\ \\tableline NGC~4151 (Event A) \\\\      & $30-60$       &                      & Penston 1974 \\\\      & $18\\pm6$      & $\\it K(U)$           & Oknyanskij 1994 \\\\      & $26\\pm6$      & $\\it L(U)$           & Oknyasnkij $\\&$ Horne 2000 \\\\ NGC4151 (Event B) \\\\      & $\\sim6$       & $\\it J(UBV)$        & Oknyanskij et al. 1999 \\\\      & $8\\pm4$       & $\\it H(UBV)$        & \\\\      & $35\\pm8$      & $\\it K(UBV)$        & \\\\      & $97\\pm10$     & $\\it L(UBV)$        & \\\\ NGC4151 (Event C) \\\\      & $94\\pm10$       & $\\it H(UBV)$        &  This work \\\\      & $48\\pm2$      & $\\it K(V)$          &  Minezaki at al. 2004 \\\\      & $104\\pm10$      & $\\it K(UBV)$       &  This work \\\\      & $105\\pm10$     & $\\it L(UBV)$        &  \\\\ \\\\ \\tableline \\end{tabular} \\end{table}    \\subsection{Data} Details and needed references about used here optical and NIR data can be found at  Oknyanskij et al.  (1999). Here we are using for analysis published in part UBV and NIR data for Case C till 2004. The details and all references  can be found in our past paper Oknyanskij (1999) and in   Lyuty (2004), which has to be published soon.    \\begin{figure} \\vspace{4cm} \\plotfiddle{oknyanskij_fig1.eps} {5 cm}{0}{70}{70}{-185}{-185} \\caption{Overall $K$ light curve (${F_K}=2.512^{9-K}$) reduced to 12\" aperture in comparison with optical $U$ flux variation (${F_U}=2.512^{11-U}$) through 27\" aperture (solid line). } \\end{figure} ",
        "conclusions": "Time delay between  optical and NIR    variations is a variable value and it is changed with state of the nucleus  activity. We have found that time delay for $K$ variations in the last low state  is about 100 days,  ie  about the same as it  is for $L$ in Cases B and C.   In a case NGC 4151 we have an opportunity to investigate changes of  the  dust hole radius  in different activity  states of  the nucleus and estimate that the time needed for   recovering  or reformation for the dust particles is  longer than several years. Alternative explanations can be: anisotropy of radiation field, shielding  of the central source on a light of sight,  and  also special orientation of the dust region."
    },
    "0410/astro-ph0410270_arXiv.txt": {
        "abstract": "We present results from a coordinated {\\it FUSE}, {\\it HST}/STIS and {\\it Chandra} campaign to study intrinsic UV and X-ray absorption in the outflow of the Seyfert 1 galaxy NGC~7469. Previous non-simultaneous observations of this outflow found two distinct UV absorption components, one of which likely corresponds to the X-ray absorber.  The {\\it FUSE} data reveal that the \\ion{O}{6} absorption in this component has strengthened over time, as the continuum flux decreased. We use  measured \\ion{H}{1}, \\ion{N}{5}, \\ion{C}{4}, and \\ion{O}{6} column densities to model self-consistently the photoionization state of the absorbers. We confirm the physical picture of the outflow in which the low velocity component is a highly ionized, high density absorber located near the broad emission line region, while the high velocity component is of lower density and resides farther from the central engine. ",
        "introduction": "Absorption edges are visible in the X-ray spectra of about one half of all low redshift AGN (Reynolds 1997; George et al.\\ 1998; Crenshaw, Kraemer \\& George 2003), virtually all of which also show high-ionization absorption lines in their UV spectra (Crenshaw et al.\\ 1999), suggesting a connection between the two phenomena. The absorption is thought to be intrinsic to the AGN because the complexes are generally blueshifted with respect to the AGNs, and because many absorbers show variability and/or non-unity covering fractions. (See references in Kriss et al.\\ in this proceeding.)  Most high resolution observations of the intrinsic absorption in AGNs, with the exception of recently published observations of Mrk~279 (Scott et al.\\ 2004), have been performed at different times in the UV and X-rays. The high degree of variability in AGN continua complicates self-consistent photoionization modeling from these non-simultaneous data, preventing firm conclusions about the nature of the relationship between UV and X-ray absorbers.  The intrinsic absorption in NGC~7469 has been studied previously using UV and X-ray observations separated by one year (Blustin et al.\\ 2003; Kriss et al.\\ 2003). These authors found two primary UV components, with outflow velocities of -569 and -1898 km s$^{-1}$, Components 1 and 2, respectively. The {\\it Chandra (CXO)}, {\\it FUSE}, and {\\it HST}/STIS campaign presented here is the first set of simultaneous, high-resolution UV and X-ray spectral observations of NGC~7469. Here, we describe the UV and X-ray observations and the properties of the AGN outflow, and we discuss the geometry of the absorption using the results from photoionization models. ",
        "conclusions": "We summarize our results as follows:  \\begin{enumerate} \\item{Component 1 shows saturated or nearly saturated \\ion{O}{6} absorption. Neutral hydrogen is clearly present, as Ly$\\alpha$ is a strong feature, but Ly$\\beta$ is not detected.  The {\\it FUSE} spectrum near  Ly$\\beta$ is heavily contaminated by molecular hydrogen absorption.  The \\ion{N}{5} and \\ion{C}{4} features associated with this velocity component are weak.}  \\item{Component 2 shows prominent \\ion{O}{6} and Ly$\\alpha$ absorption. As for Component 1, no Ly$\\beta$ is found, but the \\ion{N}{5} and \\ion{C}{4} features are stronger than those of Component 1.}  \\item{We do not detect \\ion{O}{7} or \\ion{O}{8} absorption in the {\\it CXO} spectrum, but we do detect \\ion{O}{8} emission and absorption from H-like and He-like Si, Ne, and Mg at velocities consistent with Component 1. Component 2 has no associated X-ray absorber.}  \\item{Consistent with previous results (Kriss et al.\\ 2003), the covering fraction for Component 1 is $\\sim$0.5, while that of Component 2 is consistent with one.} \\item{We see absorption variability in Component 1 over the three observation epochs (Fig.\\ 1): the \\ion{O}{6} absorption strengthened between the 1999 {\\it FUSE} observations (Kriss et al.\\ 2003) and those presented here, as the continuum flux decreased by a factor of $\\sim$1.5; and the Ly$\\alpha$ absorption increased between 2002 and 2004 as the continuum flux decreased by a factor of $\\sim$2.5.} \\item{The photoionization models combined with the \\ion{H}{1}, \\ion{N}{5}, \\ion{C}{4}, and \\ion{O}{6} column density measurements show that (log(N),U)$\\sim$(20.0,1.0) for Component 1 and (log N,U)$\\sim$(18.5,0.08) for Component 2. The \\ion{C}{4} and  \\ion{N}{5} column densities are particularly useful for breaking the degeneracy in models based on \\ion{O}{6} alone.} \\end{enumerate}  Our conclusions are consistent with the previous physical picture of the outflow (Blustin et al.\\ 2003; Kriss et al.\\ 2003). Component 1 is a highly-ionized, high density absorber that is coincident with or interior to the broad emission line region and related to the X-ray absorption and emission component. Compared with Component 1, Component 2 is a lower density and ionization parameter system located at a larger physical distance from the central engine in the absorbing outflow."
    },
    "0410/astro-ph0410578_arXiv.txt": {
        "abstract": "The electromagnetic radiation of rotating relativistic electron layers is studied numerically using particle-in-cell simulation. This system is general enough to be applied to problems involving the coherent synchrotron radiation (CSR) instability in bunch compressors and possibly in radio pulsars where the CSR instability may be used to explain the high brightness temperature and the observed spectra. The results of the simulation confirm all relevant scaling properties predicted by theoretical models and suggest that their range of validity is bigger than conservative estimates indicate. This is important for radio pulsars where the allowed parameter range is unknown. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410052_arXiv.txt": {
        "abstract": "We utilize new VLA Galactic plane catalogs at 5 and 1.4 GHz covering the first Galactic quadrant ($350^o\\le l\\le 42^o$, $|b|\\le 0.4^o$) in conjunction with the MSX6C Galactic plane catalog to construct a large sample of ultra-compact \\hii regions. A radio catalog of this region was first published by Becker et al. (1994), but we have added new observations and re-reduced the data with significantly improved calibration and mosaicing procedures, resulting in a tripling of the number of 5 GHz sources detected. Comparison of the new 5 GHz catalog and the MSX6C Galactic plane catalog resulted in a sample of 687 matches, out of which we estimate only 15 to be chance coincidences. Most of the matches show red MSX colors and a thermal radio spectrum. The scale height of their Galactic latitude distribution is very small (FWHM of $16'$ or $\\sim 40$ pc). These properties suggest that the sample is dominated by young ultra-compact \\hii regions, most of which are previously uncataloged. ",
        "introduction": "\\label{intro}  Both young and old stars are often embedded in clouds of dust and cold gas, obscuring them at wavelengths shortward of the near-infrared range. This includes star-forming regions, in which new stars and star clusters are born in huge molecular clouds, as well as various classes of evolved stars: e.g., planetary nebulae (PNe) and cool giant stars. Some of these sources are hot enough to ionize the gas in their vicinity, producing radio emission via free-free electron interactions. Cooler molecular gas and dust which still shroud these ionized regions absorb the radiation from the central source and reradiate it in the infrared. As has been shown previously, combining infrared and radio surveys can be a powerful tool in identifying specific classes of objects (i.e., \\hii regions, ultra-compact [UC] \\hii regions, and planetary nebulae) based on the simultaneous presence of the cold molecular and dust phases and the hot ionized gas (White, Becker \\& Helfand 1991; Becker et al. 1994; Szymczak et al. 2002).  Earlier studies have employed different tracers of infrared and radio emission at various levels of sensitivity and resolution according to the targets of interest. Most authors have used photometric measurements from IRAS or ISO to obtain infrared colors (e.g., Becker et al. 1994; Mart\\'{i}n-Hern\\'{a}ndez, van der Hulst, \\& Tielens 2003). Observations from the Midcourse Space Experiment (MSX) satellite mission (Price et al. 2001; Egan, Price \\& Kramer 2003) have rarely been used in combination with radio observations (see, however, work by Cohen \\& Green 2001 and by Kerton 2002). In the radio range, a variety of measures are used depending on the type of sources being investigated and the information needed. Radio emission lines give information on the velocity and distance of the source (Mooney 1992; Kurtz et al. 1994; Hayashi et al. 1994; Bourke, Hyland, \\& Robinson 1995; Walsh et al. 1995; Carpenter, Heyer, \\& Snell 2000; Sevenster 2002; Yang et al. 2002 ; Kerton \\& Brunt 2003 ), while continuum measurements are used to estimate the source ionizing flux (White, Becker \\& Helfand 1991; Becker et al. 1994; Szymczak et al. 2002). A number of extragalactic studies have used the technique of matching infrared and radio catalogs to isolate and study specific populations such as dust enshrouded star-forming galaxies and active galactic nuclei (Condon et al. 2002; Hopkins et al. 2003).  In this work we utilize new VLA Galactic plane catalogs at 5 and 1.4 GHz, the original versions of which were published by Becker et al. (1994) and Zoonematkermani et al. (1990), respectively. We compare the 5 GHz catalog to the most recent version from the MSX. In \\S\\ref{6cm} the new radio catalogs are briefly introduced. In \\S\\ref{iradio} we describe the matching of the 5 GHz radio catalog with the MSX6C catalog, and in \\S\\ref{properties} we study the spatial and color properties of the resulting subset of sources. In \\S\\ref{conc} we summarize our conclusions. ",
        "conclusions": "\\label{conc}  We have presented a new catalog of UC \\hii regions in the first Galactic quadrant. This is the most sensitive, high-resolution catalog of its kind to date, with a completeness for detecting embedded O stars at the $\\ge 75\\%$ level, depending on the optical depth of the nebula to free-free emission. We have investigated the infrared colors, the spatial distributions, the radio spectral slopes, and the morphologies of sources in this catalog. Our main findings are: \\begin{itemize} \\item[1.] The subset of 5 GHz sources with MSX6C colors is tightly confined to the Galactic plane (FWHM of $16'$ or $\\sim 40$ pc). This suggests that the sample is dominated by Population I objects, although our 5 GHz radio catalog and MSX6C infrared data could not distinguish \\hii regions from PNe. \\item[2.] By matching our 5 GHz radio catalog and the MSX6C catalog, most of the blue MSX population is filtered out; the majority of the matching subset belongs to the red MSX population. \\item[3.] By matching our 5 GHz radio catalog and the MSX6C catalog, the non-thermal radio sources are filtered out. The initial bimodal distribution of the radio spectral indices between 1.4 and 5 GHz becomes a single-peaked distribution with $<\\alpha>=0.28\\pm 0.03$. \\item[4.] The MSX6C does not detect all of the Galactic plane thermal sources; relatively weak 5 GHz sources with flat or inverted radio spectra are not detected by MSX, probably because they are below the MSX detection limit. The latitude distribution of these sources suggests a larger extragalactic fraction compared to the subset of MSX matches. We estimate the MSX6C catalog to be $\\ge 80\\%$ complete in detecting thermal sources.  \\end{itemize} The overall picture emerging from these results is that objects of nebular nature, which are detected in the radio, have red infrared colors, flat or inverted radio spectra, and a spatial distribution that is tightly confined to the Galactic plane. These properties suggest that these sources are dominated by young compact \\hii regions, most of them previously unclassified, although PNe and extragalactic sources also make a small contribution to the sample. A study of radio spectral lines could be used to confirm the source classifications and determine distances to produce a 3-dimensional spatial distribution for these objects. Infrared spectroscopy could be used to determine other properties of these sources, such as the shape of the ionizing spectrum, the type of stellar populations, and nebular properties, such as ionization state, gas temperature and density, and metallicity.  The comprehensive compilation from the literature by Paladini et al. (2003) lists only 355 known \\hii regions in our survey area, only 44 of which are smaller than $1'$ in diameter (at 2.7 GHz); thus, our new sample represents a 15-fold increase in the number of compact and UC \\hii regions in this region."
    },
    "0410/astro-ph0410264_arXiv.txt": {
        "abstract": "We report on the analysis of high resolution X-ray spectra of two pre-main-sequence stars: \\object{TWA~5} (observed with {\\em XMM-Newton}) and \\object{PZ~Telescopii} (observed with Chandra/HETGS). \\object{TWA~5} is a classical T~Tauri star in the TW Hydrae association while \\object{PZ~Tel} is a rapidly rotating weak-lined T~Tauri star in the $\\beta$-Pictoris moving group. For both stars we have reconstructed the emission measure distribution and derived the coronal abundances to check for possible patterns of the abundances related to the first ionization potential of the various elements. We have also derived estimates of the plasma density from the analysis of the He-like triplets. We compare the characteristics of our targets with those of other pre-main sequence stars previously analyzed by other authors: \\object{TW~Hya}, \\object{HD~98800} and \\object{HD~283572}. Our findings suggest that X-ray emission from classical T~Tauri and weak-lined T~Tauri stars is produced in all cases by magnetically-heated coronae, except for \\object{TW~Hya} which has unique plasma temperatures and densities. Moreover we derive that \\object{TWA~5} has the same peculiar Ne/Fe abundance ratio as \\object{TW~Hya}. ",
        "introduction": "\\label{intro}  The aim of the present work is to investigate the evolution of stellar X-ray emission during pre-main-sequence (PMS) stages. To address this issue we have analyzed the high resolution X-ray spectra of two PMS stars: the classical T~Tauri star (CTTS) \\object{TWA~5}, and the weak-lined T~Tauri star (WTTS) \\object{PZ~Tel} (\\cite{ArgiroffiDrake2004}). These two stars allow us to probe different evolutionary phases. Mass accretion is still present in \\object{TWA~5} as shown by its asymmetric and broad H$\\alpha$ emission (\\cite{MohantyJayawardhana2003}). \\cite*{JayawardhanaHartmann1999} suggested a possible infrared excess but this finding has been ruled out by subsequent studies (\\cite{MetchevHillenbrand2004,WeinbergeBecklinr2004}), indicating that \\object{TWA~5} no longer harbors large amount of dust. The accretion process has just ended in \\object{PZ~Tel} which has a filled in H$\\alpha$ line (\\cite{SoderblomKing1998}). The analysis of the X-ray emission allows us to infer the characteristics of the emitting plasma including the thermal structure, the electron densities, and the elemental abundances.  \\begin{table}[!b] \\caption{Characteristics of the star sample. Negative values of H$\\alpha$ equivalent width indicate an emission line.} \\label{tab1} \\begin{center} \\leavevmode \\scriptsize \\begin{tabular}{lcccrr} \\hline Name      & Mass         & Spectral & Age      & \\multicolumn{1}{c}{EW(H$\\alpha$)} & \\multicolumn{1}{c}{Distance} \\\\ & ($M_{\\odot}$)& Type     & (Myr)    & \\multicolumn{1}{c}{(\\AA)}         & \\multicolumn{1}{c}{(pc)} \\\\ \\hline \\object{TW~Hya}    & $\\sim0.7$   & K7       & $\\sim10$ & -220.00 \\hspace{0.01\\textwidth} & 56 \\hspace{0.015\\textwidth} \\\\ \\object{TWA~5}     & $\\sim0.5$   & M1.5     & $\\sim10$ &  -13.64 \\hspace{0.01\\textwidth} & 55 \\hspace{0.015\\textwidth} \\\\ \\object{HD~98800}  & $\\sim1$     & K5+K7    & $\\sim10$ &    0.00 \\hspace{0.01\\textwidth} & 47 \\hspace{0.015\\textwidth} \\\\ \\object{PZ~Tel}    & $\\sim1$     & K0       & $\\sim12$ &    0.63 \\hspace{0.01\\textwidth} & 50 \\hspace{0.015\\textwidth} \\\\ \\object{HD~283572} & $\\sim1.6$   & G2       & $\\sim10$ &    1.12 \\hspace{0.01\\textwidth} & 128 \\hspace{0.015\\textwidth} \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  High resolution X-ray spectroscopy studies have been performed for other PMS stars, and in particular for the CTTS \\object{TW~Hya} (\\cite{KastnerHuenemoerder2002,StelzerSchmitt2004}), and for the two WTTSs \\object{HD~98800} (\\cite{KastnerHuenemoerder2004}) and \\object{HD~283572} (\\cite{ScelsiMaggio2004a}). The analysis of \\object{TWA~5} and \\object{PZ~Tel} fills the gap between stars characterized by strong H$\\alpha$ emission lines, like \\object{TW~Hya}, and active stars near the zero age main sequence (ZAMS) with H$\\alpha$ in absorption.  \\begin{figure*}[!t] \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{cargiroffif2.ps}} \\end{center} \\caption{Upper panel: Spectrum of \\object{TWA~5} obtained by adding the RGS1 and RGS2 X-ray spectra previously rebinned by a factor 2. Lower panel: Spectrum of \\object{PZ~Tel} obtained by smoothing the Chandra/MEG first order spectra. \\label{fig2}} \\end{figure*}  \\begin{figure}[!b] \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{cargiroffif1.ps}} \\end{center} \\caption{HR diagram with isochrones (dotted lines) and evolutionary tracks (solid lines) from Siess et al. (2000). \\label{fig1}} \\end{figure} ",
        "conclusions": "In conclusion we have found that Ne/Fe abundance ratio seems to be correlated with evolutionary stages and with spectral type. In particular \\object{TWA~5} has the same Ne/Fe abundance ratio as \\object{TW~Hya}. On the other hand, none of the sample stars shows the same densities and plasma temperatures as \\object{TW~Hya}. We can therefore depict two scenarios: in the first case \\object{TW~Hya} may represent CTTSs with ages $\\sim10$~Myr, but still characterized by high accretion levels, while the other sample stars represent different evolutionary phases. In the second case \\object{TW~Hya}, and/or its environment, may be quite different from the other CTTSs, and therefore it is not straightforward to reconcile its characteristics with those of the other PMS stars.  \\begin{figure}[!t] \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{cargiroffif8.ps}} \\end{center} \\caption{Peak temperature of the $EMD$ for the selected pre-main sequence stars ordered by decreasing H$\\alpha$ emission. \\label{fig8}} \\end{figure}"
    },
    "0410/astro-ph0410608_arXiv.txt": {
        "abstract": "The central stars of planetary nebulae (PNe) are natural targets for \\emph{FUSE} due to their UV brightness. The \\emph{FUSE} spectra of many PNe show absorption features due to circumstellar material in species ranging from H$_2$ and neutrals in the photodissociation region (PDR) to ions resident in the \\ion{H}{2} region. We report results from a program designed to search for nebular components in the H$_2$ Lyman and Werner resonance lines that are responsible for the fluorescent excitation of H$_2$ in strong FUV radiation fields. Our failure to detect H$_2$ in absorption in several PNe with strong near-infrared H$_2$ emission indicates that the molecular material has an asymmetrical or clumpy distribution. We also detect enrichments in the \\emph{s}-process product Ge, find that Fe is not depleted into dust along at least one line of sight through a PN, and show that starlight fluorescence can affect the populations of the excited fine-structure levels of \\ion{O}{1}. ",
        "introduction": "The strong FUV continua of planetary nebula central stars provide excellent backdrops for observing absorption lines produced by circumstellar material as well as foreground interstellar (IS) clouds. The \\emph{FUSE} band contains a uniquely broad suite of spectral features from different phases of material that may be present in the circumstellar environment, providing information on physical conditions along a path that traverses the entire multi-phase nebular envelope.  An important consideration is the need to resolve circumstellar from IS features, for species that may be present in both components. We preferentially selected targets with favorable velocities, such as BD+30$^{\\circ}$3639, in which nebular \\ion{Na}{1} absorption is seen at $v_{helio}$ $\\sim$ --70~km~s$^{-1}$, blueshifted by $\\sim$55~km~s$^{-1}$ with respect to the IS components (Dinerstein, Sneden, \\& Uglum 1995). ",
        "conclusions": ""
    },
    "0410/astro-ph0410322_arXiv.txt": {
        "abstract": "The Andromeda Galaxy (M31) is the closest large galaxy to the Milky Way, thus it is an important laboratory for studying massive dark objects in galactic halos (MACHOs) by gravitational microlensing.  Such studies strongly complement the studies of the Milky Way halo using the the Large and Small Magellanic Clouds.  We consider the possibilities for microlensing surveys of M31 using the next generation of wide field imaging telescopes with fields of view in the square degree range.  We consider proposals for such imagers both on the ground and in space.  For concreteness, we specialize to the SNAP proposal for a space telescope and the LSST proposal for a ground based telescope.  We find that a modest space-based survey of 50 visits of one hour each is considerably better than current ground based surveys covering 5 years.  Crucially, systematic effects can be considerably better controlled with a space telescope because of both the infrared sensitivity and the angular resolution.  To be competitive, 8 meter class wide-field ground based imagers must take exposures of several hundred seconds with several day cadence. ",
        "introduction": "Almost twenty years ago Paczy\\'nski (1986) proposed that dark objects in the Milky Way halo might be found by searching for the rare gravitational microlensing of stars in nearby galaxies: the Large and Small Magellanic Clouds (LMC and SMC) and also M31.  This technique could thus be used to explore the dark matter problem in searching for these objects, dubbed MACHOs.  Both primordial black holes (a candidate for non-baryonic dark matter, known to dominate the clustering matter in the universe) and stellar remnants (a candidate for baryonic dark matter, an unknown component accounting for roughly half of the baryons today) are interesting MACHO candidates.  For a dark halo consisting of solar--mass MACHOs, of order one in one million LMC stars is being lensed (with a magnification of 30\\% or more) at any given time.  These events are transient, developing over typically 100 days.  A monitoring campaign could hope to detect these events, thus shedding light on the nature and composition of the Milky Way halo.  Several groups have performed such campaigns over more than ten years, and have extended the technique to more distant galaxies as well.  The MACHO collaboration (Alcock et al.~2000) monitored the LMC for more than 8 years, concluding from the 5.7 year dataset that there is a population of lenses making a contribution of approximately 20\\% in objects of roughly half-solar mass.  The EROS collaboration (Afonso et al.~2003) has monitored the LMC and SMC over a similar time period, and finds only an upper limit of 25\\% on the microlensing component.  The MACHO collaboration result is intriguing in that while the possibility that the dark matter consists solely of these lensing objects is ruled out, there may be a subdominant non-baryonic dark matter component.  Otherwise, the result speaks to the dark baryon problem.  More than a decade ago, Crotts (1992) pointed out that the Andromeda Galaxy (M31) would be an excellent target for a microlensing survey, in particular because of the asymmetry expected in the rate of halo lens events due to the steep inclination of the M31 disk.  Both the Milky Way and M31 halos could be studied in detail.  Subsequently, this result was discussed in depth (Jetzer 1994; Baltz \\& Silk 2000; Kerins et al.\\ 2001; Baltz, Gyuk \\& Crotts 2003).  Microlensing surveys of M31 require image subtraction, as very few stars can be resolved from the ground.  This is the regime of ``pixel'' lensing (Crotts 1992; Baillon et al.~1993; Tomaney \\& Crotts 1996; Gould 1996).  Several collaborations, including MEGA (de Jong et al.~2004) (preceded by the VATT/Columbia survey (Crotts \\& Tomaney 1996, Uglesich 2001, Uglesich et al.~2004)), AGAPE (Ansari et al.~1999, Auri\\`ere et al.~2001, Calchi Novati et al.~2002, Paulin-Henriksson et al.~2003, Calchi Novati et al.~2003), and WeCAPP (Riffeser et al.~2003), have produced a number of microlensing event candidates involving stars in M31.  The results of the VATT/Columbia survey (Uglesich et al.~2004) indicate the presence of a lensing halo at the 20\\% level, excluding the case of only self lensing at greater than 95\\% confidence.  Gould (1995) noted that the Hubble Space Telescope (HST) could perform a microlensing survey of the Virgo cluster using M87.  Such a survey has been carried out (Baltz et al.\\ 2004), placing only an upper limit that the Virgo halo is less than 60\\% lenses in the solar mass range.  The current observational situation is uncertain.  Ground based programs are moving forward, such as SuperMACHO, MEGA, and AGAPE.  To understand populations of dark objects in galactic halos, wide field imaging telescopes currently being proposed and studied would have large advantages.  In this paper we discuss the potential of wide field imaging for microlensing surveys of M31. For significant progress to be made, ground-based surveys would need to be performed with 8 meter class telescopes, or with night-long integrations on smaller telescopes.  We discuss using the proposed LSST for this purpose.  A space-based wide field survey offers many advantages relative to the ground; we discuss using the proposed SNAP space telescope for microlensing.  Graff \\& Kim (2001) have previously discussed using the SNAP telescope for a microlensing survey of the LMC. ",
        "conclusions": "We have calculated the sensitivity to microlensing of two proposed wide field imaging systems, and found that even with modest programs these would eclipse anything currently available.  More importantly, such wide field multi color surveys would address the systematic difficulties with convincingly identifying microlensing more robustly than has been possible to date.  A space based survey such as SNAP would be especially powerful for this purpose due to the availability of high resolution and infrared imaging.  The chief science goal of the microlensing surveys described in this paper would be to understand the populations of stellar--mass objects (if any) associated with the halos of spiral galaxies.  This speaks to the issue of the dark baryons, still unidentified, and making up roughly half of the baryon budget in the present-day universe.  In addition to dark halo objects, other science products of such surveys would include a census of variable stars (in several filters), and other transient events.  One interesting possibility is the detection of binary lens events with extreme mass ratios, indicative of giant planets.  This has been recently achieved toward the galactic bulge with a solid detection by the OGLE and MOA teams (Bond et al.\\ 2004), followed by another detection by OGLE (Jaroszy\\'nski et al.\\ 2004).  Detecting planets would necessarily involve reducing the survey data in real time, so that microlensing alerts could be issued, and candidates could be followed closely, either from the ground or in space.  We conclude this paper with a simple observation.  The wide field (of order square degree) imaging systems being discussed now have obvious applications for cosmological surveys.  In contrast, there are only a handful of objects on the sky with angular sizes commensurate with these imagers, including the Milky Way bulge and disk, the Magellanic Clouds, and M31.  It seems warranted that at the very least, modest studies of these very few nearby objects should be undertaken with the upcoming wide field technologies."
    },
    "0410/astro-ph0410444_arXiv.txt": {
        "abstract": "The Square Kilometer Array (SKA) will enable studies of star formation in nearby galaxies with a level of detail never before possible outside of the Milky Way.  Because the earliest stages of stellar evolution are often inaccessible at optical and near-infrared wavelengths, high spatial resolution radio observations are necessary to explore extragalactic star formation.  The SKA will have the sensitivity to detect individual ultracompact HII regions out to the distance of nearly 50~Mpc, allowing us to study their spatial distributions, morphologies, and populations statistics in a wide range of environments.  Radio observations of Wolf-Rayet stars outside of the Milky Way will also be possible for the first time, greatly expanding the range of conditions in which their mass loss rates can be determined from free-free emission.  On a vastly larger scale, natal of super star clusters will be accessible to the SKA out to redshifts of nearly $z \\sim 0.1$.  The unprecedented sensitivity of radio observations with the SKA will also place tight constraints on the star formation rates as low as $1 M_\\odot$~yr$^{-1}$ in galaxies out to a redshift of $z \\sim 1$ by directly measuring the {\\it thermal} radio flux density without assumptions about a galaxy's magnetic field strength, cosmic ray production rate, or extinction.  \\vspace{1pc} ",
        "introduction": "\\subsection{The Importance of Studying Extragalactic Star Formation}  Star formation is a critical process -- arguably one of the most fundamental physical processes (next to gravitational collapse) in determining the appearance and properties of the visible universe. The formation of massive stars is particularly important because they have a major effect on the energetics of galaxies: massive stars are responsible for the ionization of the interstellar medium, their stellar winds and supernovae are main sources of mechanical energy, their ultraviolet radiation powers thermal-infrared (IR) luminosities through the heating of dust, they are a main driver of chemical evolution via supernova explosions at the end of their lives, and they may be the sources of gamma ray bursts.  Nevertheless, despite the significant role of massive stars throughout the universe, their birth is not well understood and we are only beginning to piece together a scenario for the youngest stages of massive star evolution.  There has been some progress understanding the early stages of massive star formation in the Milky Way, but the current knowledge about the early stages of massive star evolution in other environments is mediocre at best.  The reasons for this dearth of information about extremely young massive stars in other galaxies are predominantly threefold: (1) the earliest stages of stellar evolution are deeply enshrouded and inaccessible at optical and near-infrared wavelengths; (2) compared to other types of radio sources (e.g. AGN, SNe, SNR), individual stars are relatively faint and require exceptionally sensitive observations; and (3) high spatial resolutions are necessary to disentangle the individual massive stars from their surrounding environment and background contamination.  The study of star formation in environments outside of our own Milky Way is critical for advancing our understanding of stellar and galactic evolution.  The Square Kilometer Array (SKA) will clearly enable investigations of star formation in the Milky Way with an unprecedented level of detail.  However, without similar investigations in other galaxies, it is difficult to place Galactic work in a larger context.  The SKA will revolutionize the field of star formation by enabling detailed studies that have never been possible outside of the Milky Way.  The ability to observe star formation taking place in a range of conditions will enable us to better interpolate and extrapolate details of the star formation process throughout the universe.  For example, we do not know how the properties of star formation depend on various environmental parameters, including metallicity, pressure, turbulence, stellar densities, triggering scenarios (including bars, bubbles, and galactic interactions), or how star formation might differ in nuclear regions or in ``burst'' and quiescent modes. Investigations of all of these issues will be possible with the sensitivity and resolution of the SKA. ",
        "conclusions": ""
    },
    "0410/gr-qc0410129_arXiv.txt": {
        "abstract": "s{In brane-worlds, our universe is assumed to be a submanifold, or brane, embedded in a higher-dimensional bulk spacetime. Focusing on scenarios with a curved five-dimensional bulk spacetime, I discuss their gravitational and cosmological properties. } ",
        "introduction": "In this contribution, I review the basic ingredients of the so-called brane-world models. These  models are based on the  assumption that our universe is a {\\it brane}: a sub-space embedded in a higher dimensional {\\it bulk} spacetime. In contrast with the traditional Kaluza-Klein treatment  of extra dimensions, ordinary matter fields are {\\it confined} on the brane. As a consequence, the size of the extra-dimensions in brane-world models is  not limited by the usual bound, typically $R\\lesssim (1{\\rm TeV})^{-1}$, based on the non-detection of Kaluza-Klein modes in collider experiments up to energies close to $1$ TeV. In brane-worlds, only gravity propagates in the higher-dimensional spacetime and, therefore, the size of the extra-dimensions can be much larger, up to a fraction of a millimeter, corresponding to the current limit of gravity experiments that look for deviations of Newton's law\\cite{Hoyle_etal04}.   Although there were some precursor works on brane-worlds, the huge interest they have encountered in the last few years is due to recent developments in string/M-theory. This has in turn opened up new avenues to tackle  some fundamental problems.  For example, the fact that the four-dimensional Planck mass $M_p$ is in this context only a ``projection'' of the higher-dimensional (fundamental) Planck mass $M_*$, which can thus be lower than $M_p$, offers a new perspective on the hierarchy problem  and suggests the possibility that quantum gravity might be  closer than previously thought. This has been emphasized by Arkani-Hamed, Dimopoulos and Dvali\\cite{add}, who suggested a very simple model based  on a brane embedded in a flat geometry with $4+n$ dimensions, where $n$ dimensions are compactified on a torus of size $R$. Using Gauss' theorem, one can easily show that, on scales $r\\gg R$,  standard 4D gravity is recovered with $M_P^2=M_*^{2+n}R^n$. On scales $r\\ll R$, gravity becomes $(4+n)$-dimensional.  Another important progress was made by Randall and Sundrum\\cite{rs99a,rs99b}, who considered  {\\it curved}, or {\\it warped}, bulk geometries. They showed, in particular, that  {\\it compact} extra-dimensions are not necessary  to obtain a four-dimensional behaviour. The bulk curvature can indeed lead to an {\\it effective compactification}.  In this contribution, I will  focus my attention on brane-worlds characterized by a single extra  dimension   where  the bulk space-time is {\\it curved} instead of flat and where the  self-gravity of the brane is  taken into account. This includes the configurations discussed by Randall and Sundrum.  The outline is the following. In the next section, I will  present in some detail the Randall-Sundrum models and discuss the corresponding effective gravity on the brane. Then, in the subsequent sections, I will turn my attention to cosmological brane models, reviewing a few important topics in brane cosmology. In section 3, the bulk and brane geometries will be presented. Section 4 will be devoted to the so-called dark radiation, or Weyl radiation. Inflation in the brane will be discussed in section 5. Section 6 will discuss some aspects of brane cosmology when  the assumption of isotropy or homogeneity is relaxed. Finally, section 7 will provide some conclusions and perspectives.  Only some aspects of brane-worlds are discussed in this contribution. Much more can be learnt on this subject  from several detailed reviews~\\cite{reviews}. ",
        "conclusions": "In this contribution, I have presented some aspects of brane-world models, covering both the (static) Randall-Sundrum model and its cosmological extensions. Due to lack of time/space, I have not discussed many other  interesting topics in the field. Examples are the brane cosmology of  models involving Gauss-Bonnet corrections; the induced gravity  models, where one includes a 4D Einstein-Hilber action for the brane and which can lead to late-time cosmological effects mimicking dark energy.    There are still many open questions in brane cosmology. Even in the simplest set-up, discussed here, based on a cosmological extension of the Randall-Sundrum model, the evolution of cosmological perturbations has not yet been solved, although some significant progress has been made. The situation is still more complicated  in more sophisticated models, involving a bulk scalar field and/or collision of branes. It must be emphasized that the predictions for the  cosmological perturbations, as observed in the CMB experiments, and their adequation with the present data, is a crucial test for brane-world models for which the early universe is modified. More direct tests of brane-world models involve gravity experiments or  collider experiments. However, if the fundamental Planck mass is too high, such direct experiments cannot see extra-dimensional effects and one must turn to cosmology to try to see indirect signatures from the early universe.  Another direction of  research is to make contact between the brane-worlds, which are still only  phenomenological models, and a fundamental theory like string theory."
    },
    "0410/astro-ph0410672_arXiv.txt": {
        "abstract": "We present \\halpha\\ rotation curves of two low luminosity field galaxies with $r$-band absolute magnitudes $\\Mr=-13.9^{+0.8}_{-0.5}$ and $\\Mr=-14.7^{+0.3}_{-0.2}$ (for $h\\equiv H_0/100\\hubunits =0.7$; the large error bars reflect distance uncertainties). Most previously studied galaxies in this luminosity range are members of groups defined by brighter galaxies, but these two systems, selected from Blanton et al.'s (2004) sample of low luminosity galaxies in the Sloan Digital Sky Survey (SDSS), appear to have no bright companions. The measured rotation speeds at the outer extent of the \\halpha\\ rotation curves, $34.8\\pm 3.8\\kms$ and $30.9\\pm 7.2\\kms$, are larger than the values of $16.0^{+6.1}_{-5.4}\\kms$ and $20.9^{+6.2}_{-5.2}\\kms$ predicted by extrapolating the inverse Tully-Fisher relation of luminous SDSS galaxies to these faint luminosities. However, a previous HI measurement of the first galaxy shows that it has a gas mass similar to its stellar mass, and the total baryonic mass is consistent with that predicted by McGaugh et al.'s (2000) ``baryonic Tully-Fisher relation.'' We find $r$-band dynamical mass-to-light ratios within the radii of the last \\halpha\\ data points (about 1.8 disk scale lengths in each case) of $12.6^{+4.7}_{-4.5}\\, M_\\odot/L_\\odot$ and $4.8_{-2.1}^{+2.5}\\, M_\\odot/L_\\odot$, much higher than the values $\\sim 1\\,M_\\odot/L_\\odot$ expected for the stellar populations. The dynamical properties of these galaxies, including the rotation speeds and evidence for high gas fractions and dark matter domination within the luminous extent of the galaxy, are consistent with those of previously studied faint galaxies in nearby groups. Further studies of the SDSS sample will allow characterization of low luminosity galaxies over the full range of environments in which they reside. ",
        "introduction": "The Sloan Digital Sky Survey \\citep{york00} galaxy redshift survey has an unprecedented combination of large area, depth, and photometric quality, thanks to the combination of a large format camera \\citep{gunn98}, high throughput multi-object spectrographs (A.\\ Uomoto et al., in preparation), careful calibration procedures \\citep{fukugita96,hogg01,smith02}, and an efficient series of data reduction and targeting pipelines \\citep{lupton01,stoughton02,strauss02,blanton03a,pier03}. Recently \\citeauthor{blanton04a} (\\citeyear{blanton04a}, hereafter B04) have searched the SDSS Second Data Release (DR2; \\citealt{abazajian04}) to identify a population of extremely low luminosity field galaxies. With absolute magnitudes $M_r \\sim -13$ to $-16$ (for the Hubble parameter $h\\equiv H_0/100\\hubunits=0.7$ adopted throughout this paper), this population represents a range of luminosities that has previously been accessible to systematic study only in the Local Group and in nearby groups and clusters.\\footnote{The field galaxy survey of \\cite{schombert97}, based on visual inspection of POSS-II plates and HI confirmation from Arecibo, includes some systems in this luminosity range.} We have obtained H$\\alpha$ rotation curves of two of these objects, using the same methods that we are using for a study of the Tully-Fisher (\\citeyear{tul77}, hereafter TF) relation of more luminous SDSS galaxies (J. Pizagno et al., in preparation, hereafter P04). Both objects appear to be ``field'' galaxies rather than satellites of brighter systems. Although the sample is small and the data quality limited, these measurements provide one of the first insights into the dynamical properties of this new population. ",
        "conclusions": "How do our measured rotation speeds compare to expectations based on the TF relation defined by more luminous galaxies?  In P04, we use a sample of 170 galaxies with $-22 < \\Mr < -18$ to measure the forward ($M$ vs. $\\log V$) and inverse ($\\log V$ vs. $M$) TF relations in the SDSS bands. For the $r$-band inverse relation, we fit \\begin{equation} \\eta=a(\\Mr-M_{r,0})+b \\end{equation} with $\\eta \\equiv \\log_{10}\\Vend^i$ and the constant $M_{r,0}=-20.873$ chosen to yield minimal correlation between the errors in $a$ and $b$. We find $a=-0.143 \\pm0.018$ and $b=2.202 \\pm0.004$, yielding predicted values of $\\Vend^i$ for Galaxies 1 and 2 of $V_{\\rm pred}=16.0^{+6.1}_{-5.4} \\kms$ and $V_{\\rm pred}=20.9^{+6.2}_{-5.2} \\kms$, respectively. The uncertainties in $\\Mr$ and in the TF slope both contribute significantly to the errors on $\\Vpred$, and both contributions are non-linear. We have computed upper $1\\sigma$ error bars by separately varying $\\Mr$ by $-1\\sigma$ and $a$ by $+1\\sigma$ and adding the corresponding values of $\\delta \\Vpred$ in quadrature; we follow an analogous procedure for the lower $1\\sigma$ error bar.  The predicted rotation speeds are well below the measured values of $34.8 \\pm 3.8 \\kms$ and $30.9 \\pm 7.2\\kms$, though marginally compatible in the latter case because of the large observational errors. Dotted curves in Figure~2 show rotation curves with the same parameters as the solid curves but scaled by $(V_{\\rm pred}/\\Vend^i)$.  Despite the uncertainties in the measurements, it is clear that the galaxies are rotating substantially faster than predicted by an extrapolation of the inverse TF relation to these faint magnitudes.  To match the observed rotation speeds while maintaining the P04 normalization at $\\Mr \\approx -21$ would require an inverse TF slope of $a \\approx -0.1$, compared to the slope of $-0.143$ measured in the bright galaxy regime.  We can use our measurements to estimate dynamical mass-to-light ratios within the radius of the outermost \\halpha\\ data point in each galaxy. We use the simple mass estimator \\begin{equation} M=\\frac{(\\Vend^i)^2\\Rend}{G}~, \\end{equation} ignoring possible corrections for asymmetric drift or non-circular motions, which would be small compared to our observational error bars. The values of $\\Rend$ are 19 and 11 arc-seconds for Galaxies 1 and 2, respectively, corresponding to 1.7 and 1.8 times the exponential scale lengths measured by GALFIT.  At our estimated distances of $9.9 \\pm 2.9$ Mpc and $20 \\pm 2.2$ Mpc, the physical values of $\\Rend$ are $0.91 \\pm 0.27$ kpc (Galaxy 1) and $1.07 \\pm 0.12$ kpc (Galaxy 2). We measure the fraction of $r$-band light within $\\Rend$ for each galaxy and correct the values of $\\Mr$ in Table~1 accordingly, by 0.24 and 0.14 magnitudes.  We do not apply any internal extinction corrections. Adopting $M_{r,\\odot}=4.67$, we find $M/L=12.6^{+4.7}_{-4.5}\\, M_\\odot/L_\\odot$ for Galaxy 1 and $4.8_{-2.1}^{+2.5}\\, M_\\odot/L_\\odot$ for Galaxy 2. Because $M/L \\propto \\Vend^2/d$ and the fractional uncertainties in these quantities are substantial, we compute the upper $1\\sigma$ error bar by separately varying $d$ by $-1\\sigma$ and $\\Vend$ by $+1\\sigma$ and adding the two changes $\\delta M$ in quadrature, and we follow an analogous procedure for the lower $1\\sigma$ error bar. The $M/L$ value for Galaxy 1 is much higher than the value $\\sim 1.3 M_\\odot/L_\\odot$ expected for a stellar population with the observed $g-r$ color of the galaxy \\citep{bell03}, so this system is strongly dominated by dark matter within $\\Rend$. The case for dark matter domination in Galaxy 2 is less clear because of the large uncertainty in $M/L$, though the central value again corresponds to a large ratio of dark matter to stellar mass within $\\Rend$.  (The \\citeauthor{bell03} [\\citeyear{bell03}] models imply a stellar mass-to-light ratio of $\\sim 0.9 M_\\odot/L_\\odot$ for Galaxy 2's $g-r$ color.) While all disk galaxies become dark matter dominated at sufficiently large radii, these low luminosity systems appear to have high dark matter fractions even within two disk scale lengths.  To put our measurements in the context of previous results, Figure~3 plots inclination corrected circular velocities against $B$-band absolute magnitude for our galaxies (filled circles with error bars) and for galaxies with HI line widths and $M_B \\ga -15$ from \\citeauthor{cot00} (\\citeyear{cot00}; open squares), \\citeauthor{beg04} (\\citeyear{beg04}; open circles), and HI line widths from \\citeauthor{carignan88} (\\citeyear{carignan88}; open triangle).  For our galaxies, we use $\\Vend^i$ as the indicator of circular velocity, and we convert the (AB) $g$-band luminosity to the (Vega) $B$-band luminosity using the galaxy $g-r$ color and $M_B = M_g + 0.365 + 0.46[(g-r) - 0.78]$, obtained using the $K$-correction code of \\cite{blanton03b}. For \\cite{cot00} we use the value of $V_c$ listed in their tables, and for \\cite{beg04} and \\cite{carignan88} we estimate $V_c$ visually from the plotted rotation curves.  We have not attempted to put observational error bars on the literature data points because this would require a detailed assessment of the uncertainties in the distance measurements, which use a somewhat different method in each case. \\cite{cot00} observe galaxies in the Sculptor and Centaurus A groups and assign group distances based on a variety of indicators (see \\citealt{cote97} and references therein). \\cite{beg04} assign both of their galaxies to the NGC 4696 group on the basis of proximity and adopt a distance based on the brightest stars \\citep{huchtmeier00}.  \\cite{carignan88} use a combination of group assignment and brightest stars to infer the distance of DDO 154.  The solid line in Figure~3 shows the $B$-band inverse TF relation extrapolated to the low luminosity regime. We convert the $g$-band relation of P04 to $B$-band using $M_B=M_g+0.19$, appropriate for a galaxy with $g-r=0.4$. Dotted lines show relations with the slope varied by $\\pm 1\\sigma$.  For bright galaxies, the intrinsic scatter about the mean relation is $0.07\\,$dex in $\\log_{10}\\Vend^i$ (P04), smaller than the $0.1\\,$dex observational error bar on Galaxy 1. As noted earlier (on the basis of $r$-band data), Galaxy 1 is rotating substantially faster than the extrapolated TF relation predicts, while the large error bar on $\\Vend$ leaves Galaxy 2 marginally consistent with the TF extrapolation.  The rotation speed of Galaxy 1 appears reasonably in line with that of the other dwarf galaxies, while the rotation speed of Galaxy 2 is noticeably low, perhaps because the optical rotation curve is still rising steadily at our outermost data point.  All of these systems are rotating faster than the TF extrapolation predicts; equivalently, they are fainter than predicted given their rotation speeds.  \\cite{mcgaugh00} show that low luminosity galaxies are fainter than predicted by the (forward) TF relation defined by bright galaxies, consistent with the result shown in Figure~3.  However, faint galaxies are often gas rich, and they show that adding the gas masses inferred from HI observations to the stellar masses leads to a linear relation between $\\log M_{\\rm bar}$ and $\\log V_c$ over four decades in baryonic mass $M_{\\rm bar}$, extending down to a few $\\times 10^7 M_\\odot$.  For our Galaxy 1, \\cite{mat00} report an HI flux integral of 1.24 Jy$\\kms$, which implies an HI mass to optical luminosity ratio of $\\sim 0.95 M_\\odot/L_\\odot$ in $r$-band, or roughly equal mass in neutral gas and stars.  This ratio is below the median of $\\sim 2M_\\odot/L_\\odot$ found (in $V$-band) by \\cite{pildis97}, but within the range spanned by their data.  Simply doubling the luminosity of Galaxy 1 and thus shifting it rightwards by 0.75 mag in Figure~3 would not move it onto the P04 inverse TF relation.  However, \\cite{mcgaugh00} find a steep slope for the (forward) baryonic mass TF relation, and the baryonic mass of $7.5\\times 10^7 M_\\odot$ inferred for Galaxy 1 at a distance of 9.9 Mpc agrees respectably with the value $M_{\\rm bar} = 30.5 (h/0.7)^{-2} (V_c/\\kms)^4 M_\\odot = 5.3 \\times 10^7M_\\odot$ predicted by their fitted relation for $V_c=35\\kms$, given the substantial error bars on both the predicted and measured values.  The luminosity-velocity relations and mass-to-light ratios of low luminosity galaxies offer clues to the role of supernovae and photoionization in regulating star formation, since these feedback effects are generally expected to be stronger in lower mass halos (see, e.g., \\citealt{dekel86,quinn96,thoul96,benson02,somerville02}). The influence of feedback may be different in galaxies that are central objects of their parent dark matter halos and in galaxies that are satellites in more massive halos --- semi-analytic models and hydrodynamic simulations predict systematically different properties for these two populations \\citep{kauffmann93,cole94,somerville99,berlind03,zheng04}. Most of the very low luminosity galaxies studied to date have been discovered or investigated because they are members of groups or clusters defined by brighter galaxies. The two galaxies studied in this paper are not satellites of bright neighbors, but their properties are roughly in line with those of satellite dwarfs. The B04 sample drawn from the SDSS provides the opportunity to study such faint galaxies in detail over the full range of environments in which they appear.  Our present investigation represents a start on this broader program."
    },
    "0410/astro-ph0410391_arXiv.txt": {
        "abstract": "{% We describe the decomposition of QSO absorption line ensembles applying an evolutionary forward modelling technique. The modelling is optimized using an evolution strategy (ES) based on a novel concept of completely derandomized self-adaption. The algorithm is described in detail. Its global optimization performance in decomposing a series of simulated test spectra is compared to that of classical deterministic algorithms. Our comparison demonstrates that the ES is a highly competitive algorithm capable to calculate the optimal decomposition without requiring any particular initialization. ",
        "introduction": "%  The standard astronomical data analysis packages such as the Image Reduction and Analysis Facility (IRAF) and many popular custom-built applications offer only deterministic algorithms for the purpose of parametric model fitting. In practice, however, deterministic algorithms often require considerable operational intervention by the user. In contrast, stochastic strategies such as evolutionary algorithms minimize the operational interaction, a highly appealing feature.  In practice, any parametric model fitting technique reduces to the numerical problem of finding the minimum of an objective function $f:\\mathbb{R}^n\\to\\mathbb{R}$, i.e. constructing a sequence of object parameter vectors \\begin{equation} (\\vec{x}^{(g)})_{g\\in\\mathbb{N}}, \\quad \\vec{x}^{(g)}\\in\\mathbb{R}^n \\label{eq:sequence} \\end{equation} such that $\\lim_{g\\to\\infty}f(\\vec{x}^{(g)})$ is as small as possible. Several classical algorithms are practicable to overcome the problem of minimization. Among the best established strategies are the conjugate gradient, variable metric, and Levenberg-Marquardt algorithms, all gathering information about the local topography of $f$ by calculating its partial derivatives (i.e. the gradient or the Hessian matrix) and thereby ensuring the rapid convergence at a close minimum \\citep[e.g.,][Numerical Recipes]{PressTVF_2002}. The nature of these classical algorithms is deterministic: each member of the sequence $(\\vec{x}^{(g)})_g$ is determined by its predecessor, i.e. the whole sequence is determined by the intial object parameter vector $\\vec{x}^{(0)}$. Exactly this fact turns out to be the cardinal deficiency in case $f$ exhibits many local minima: The success of the algorithms in locating the global minimum of $f$ crucially depends on the adequacy of the initial guess $\\vec{x}^{(0)}$. Hence, the results are biased and several optimization runs are normally necessary. In particular in case of noisy data or strong inter-parameter correlations the problem of finding an adequate initialization $\\vec{x}^{(0)}$ is tedious and often dominates the time needed to complete the optimization. In addition, the classical algorithms are not practicable if the objective function is not continuously differentiable, oscillating, or the calculation of partial derivatives is too expensive or numerically inaccurate.  In contrast, stochastic strategies such as evolutionary algorithms where each member of the sequence $(\\vec{x}^{(g)})_g$ is the result of a random experiment, are not only promising for the purpose of global optimization but also apply to any objective function which is computable or available by experiment. However, the drawback of stochastic optimization strategies is less efficiency: even if an adequate local downhill move exists, a random step will almost always lead astray. Therefore, many more evaluations of the objective function are required before the sequence $(\\vec{x}^{(g)})_g$ converges at a minimum.  Evolutionary algorithms are inspired by the principles of biological evolution. Conceptually, three major subclasses are discerned: evolutionary programming, genetic algorithms, and evolution strategies (ES). While all subclasses apply quite generally, ES are particularly suited for the purpose of continuous parametric  optimization. Largely \\citet[hereafter Paper~A]{HansenN_OstermeierA_2001} and \\citet[hereafter Paper~B]{HansenMK_2003} have established a novel concept of completely derandomized self-adaption in ES which selectively approximates the inverse Hessian matrix of the objective function and thereby considerably improves the efficiency in case of non-separable problems or mis-scaled parameter mappings while even increasing the chance of finding the global optimum in case of strong multimodality.  The interpretation and analysis of QSO absorption lines basically involves the decomposition of line ensembles into individual line profiles. In general, the decomposition of QSO spectra presents an ambigous parametric inverse problem and automatizing the decomposition demands an efficient but at first stable optimization algorithm. In this study we resume the concept of completely derandomized self-adaption introduced in \\citetalias{HansenN_OstermeierA_2001} and test the global optimization performance of the resulting ES when applied to the problem of line profile decomposition in QSO spectra. ",
        "conclusions": "%  All tested classical algorithms require an adequate initial guess to locate the global optimum of the objective function when applied to the problem of spectral decomposition. In contrast, the CMA-ES is demonstrably capable to calculate the optimal decomposition without demanding any particular initialization, and therefore is just asking to be used in automatized spectroscopic analysis software.  The CMA-ES does not guarantee to find the optimal solution: characteristic spectral features are modelled correctly, but features exhibiting just a small attractor volume such as narrow components or tight blends are frequently not distinguished fittingly. Larger populations are expected to improve the chance to hit the global optimum but require a greater number of objective function evaluations. Peak finding algorithms could detect both the proper number (a problem that we have completely left aside) and position of spectral components and provide an improved initialization. But irrespective of whether an automatized analysis software will be realized, the CMA-ES is an elegant and highly competitive general purpose algorithm which is easy to implement. In our opinion, its integration into the standard astronomical data analysis packages will be thoroughly worthwhile and its widespread use will contribute to the further comprehension and improvement of suchlike algorithms."
    },
    "0410/astro-ph0410358_arXiv.txt": {
        "abstract": "The atmospheres of cool, helium-rich white dwarfs constitute an exotic and poorly explored physical regime of stellar atmospheres. Under physical conditions where the temperature varies from $\\rm 1000K$ to $\\rm 10000K$, the density can reach values as large as $\\rm 2 \\ g/cm^{3}$, and the pressure is as high as $\\rm 1 \\ Mbar$, the atmosphere is no longer an ideal gas and must be treated as a dense fluid. Helium atoms become strongly correlated and refraction effects are present. Opacity sources such as $\\rm He^{-}$ free-free absorption must be calculated with a formalism that has never been applied to astrophysical opacities. These effects have been ignored in previous models of cool white dwarf atmospheres. ",
        "introduction": "Very cool white dwarfs (WDs) with $\\it T_{\\rm eff}\\rm<5000K$ are among the oldest objects in our Galaxy. They are of great importance for cosmochronology and for understanding the formation and evolution of the Milky Way. Our main interest is in the atmospheres of these stars as the link between their physical characteristics and the observables. In this contribution we concentrate on the atmospheres rich in helium. Due to the high transparency of helium (compared to hydrogen), our current models predicts for them densities that are typical of liquids, up to $\\rm 2-3 \\ g/cm^{3}$. Under such extreme conditions, the input physics of atmosphere models needs to be reconsidered.  There are several physical effects that distinguish fluid from gaseous atmospheres. They all arise from the strong interaction between particles. Helium atoms become strongly correlated, the refractive index of the fluid departs from unity, and free electrons strongly interact with the medium, which affects the opacity.  Non-ideal effects in the equation of state are dominated by $\\rm He-He$ interactions. For a given density and temperature, they result in increases in the pressure and ionization fraction and a decrease in the adiabatic gradient.  The refractive index inside helium-rich white dwarf atmospheres departs from unity and can be as large as 1.35. We have solved the equation of radiative transfer for refractive stellar atmospheres for the first time (Kowalski \\& Saumon 2004). The impact of the total internal reflection on the radiative equilibrium results in an increase of the temperature in the optically thin atmospheric layers. This decreases the abundance of hydrogen molecules and the $\\rm H_{2}-He$ CIA opacity in atmospheres of mixed composition. Due to strong refraction near the surface, the limb darkening almost disappears.  Fluid atmospheres require a new description of the opacity that is different from that for diluted gases. The calculation of the $\\rm He^{-}$ free-free absorption (or inverse bremsstrahlung) in WD atmospheres must be revised. Preliminary results obtained from quantum molecular dynamics simulations (QMD) suggest that fluid, helium rich, cool WD atmospheres are much more opaque than predicted by current models.  All these effects must be considered to properly model cool WD atmospheres. Realistic models of very cool WDs are of fundamental importance for cosmochronology and to understand the spectra and physical properties of the coolest WDs, especially the recently discovered, peculiar ultracool stars like LHS 3250. ",
        "conclusions": ""
    },
    "0410/astro-ph0410502_arXiv.txt": {
        "abstract": "{ The three basic ingredients - a spinning compact object, an accretion disc and a collimated relativistic jet - make microquasars  a galactic scaled-down version of the  radio-loud AGN.  That explains the large interest attributed  to this new class of objects, which up to now consists of less than 20 members. Microquasars belong to the much larger class of  X-ray binary systems, where there exists a compact object  together with its X-ray emitting accretion disc, but the  relativistic jet is missing. When does an X-ray binary system evolve into a microquasar? Ideal for studying such  kind of a transition is the periodic microquasar LS~I~+61$^{\\circ}$303  formed by a  compact object  accreting  from the equatorial wind of a Be star and with  more than one event of  super-critical accretion and ejection  along the eccentric orbit. For ejections at periastron passage the relativistic electrons suffer  severe  inverse Compton losses by upscattering the UV photons of the Be star  at high energy : At periastron passage Gamma-ray emission has been observed, whereas radio outbursts have never been observed in 20 years of radio flux monitoring. For ejections displaced from periastron passage the losses are less severe and radio outbursts are observed. The radio emission   mapped on scales  from a few AU  to hundreds of AU shows a double-sided relativistic ($\\beta=0.6c$) S-shaped jet, similar to the  well-known precessing jet of \\object{SS~433}. } ",
        "introduction": "Since the beginning of the 1980s radio-galaxies,  quasars, Seyferts, QSO  etc. all are simply classified as  AGNs (``Active Galactic Nuclei'') because the ``energy-engine'' is thought to be the same: A super-massive black hole accreting from its host galaxy. AGNs having  radio-emitting lobes or jets are called radio-loud, the others are called radio-quiet (Ulrich et al. 1997).  In  an X-ray binary system the ``energy-engine'' is  a compact object of a few solar masses accreting from the companion star. Up to now there are almost 250  known X-ray binaries (Liu 2000). Only a small percentuage of them ($<$10\\%) show evidence of a radio-jet and therefore are radio loud applying the same definition as for the AGNs. The radio loud X-ray binaries  subclass (Fig. 1) includes together with the microquasars --objects where high resolution radio interferometric techniques have shown the presence of collimated  jets (Mirabel et al. 1992)-- also   unresolved radio sources with a flat spectrum. This  spectrum can arise from the combination of emission from optically thick and thin regions of an  expanding  continuous jet (Hjellming \\& Johnston 1988;   Fender 2004) as  has been shown by the discovery of a continuous   jet for the flat-spectrum source Cygnus X-1 (Stirling et al. 2001). \\begin{figure*} \\centering \\includegraphics[width=15cm]{massi_fig1.ps} \\caption{ The   basic components of a microquasar - a spinning compact object, an accretion disk and a collimated relativistic jet. The compact object, of a few solar masses,   accretes from a normal star in orbital motion around it.  The mass of the compact object can be  determined by studying  the periodical shift of the optical spectral lines of the normal star  and   establishes, if it is a neutron star or a black hole. The inner part of the disk emits X-rays. The inner radius, three times the Schwarzschild radius, is a few tens of kilometers, the outer radius a factor 10$^3$ larger (the figure is not to scale). Due to magneto-rotational instabilities part of the disk is propelled into a relativistic jet, studied at high resolution with radio interferometric techiques. In some microquasars, like SS433 and LS I+61$^{\\circ}$303, the jet is precessing. If the precession causes the jet to be aligned toward the earth the large variable doppler boosting  mimics the variability of  Blazars and the microquasar in this case is called microblazar. } \\end{figure*}  Besides the first radio loud  X-ray binary system SS433, discovered by chance in 1979 (Margon 1984), the others have been discovered mostly in the last ten years and now  are considered as an ideal, nearby  laboratory for studying the processes of accretion and ejection around black holes.   More generally speaking, the class of the X-ray binaries is formed by  stellar systems of two stars with very different natures:  a normal star (acting as mass donor) and a compact object (the accretor) that can be  either a neutron star or a black hole (White et al 1996). The normal star  orbits around the compact object and therefore the infalling material has some angular momentum ($J$), which prevents  it  from falling directly into the accretor. The stream of matter orbits the compact object with a radius determined  by  $J$  and  the  mass of the compact object ($M_x$). The angular momentum is  redistributed by the viscosity: Some of the material takes angular momentum  and spreads outwards, whereas other material spirals inwards. In this  way a disk is created from the initial ring of matter (Longair 1994; King 1996). Gradually,  the matter drifts inwards until it reaches the last stable orbit, The viscosity has two effects: Beside the transport of the angular-momentum it also acts as a frictional force resulting in the dissipation of heat. The amount of friction depends on how  fast the gas orbits around the compact object. It reaches its maximum at the inner disk, where the matter is heated up to quite high temperatures (tens of millions of  degrees), producing the strong  thermal X-ray radiation giving the  name ``X-ray binaries'' to this class of objects.  In the case of a low vertical magnetic field threading the disk the  plasma pressure dominates the magnetic field pressure and the differentially rotating  disk bends the magnetic field, which is passively wound up (Meier et al. 2001). Due to the compression of the magnetic field lines the magnetic pressure may  become larger than the gas pressure at the surface of the accretion disk, where the density is lower. At this point the gas follows the twisted magnetic field lines, creating  two spinning flows. These extract angular momentum from the surface of the disk (magnetic braking) and enhance the radial accretion. The avalanching material further pulls the deformed magnetic field with it and afterwards magnetic reconnection may happen (Matsumoto et al. 1996). The thickness of the disc is fundamental in this  magneto-rotational process, or better the extent of the poloidal magnetic field frozen into the disc (Meier 2001; Meier et al. 2001; Maccarone 2004). No radio jet is associated at  X-ray binaries in High/soft states, where the X-ray spectrum is dominated by a geometrically thin (optically thick) accretion disc (Shakura \\& Sunyaev 1973). Whereas, numerical results show a jet being launched from an  inner geometrically thick portion of the accretion disc existing (coronal flow/ADAF) when the   X-ray binaries are in their low/hard state (Meyer et al. 2000; Meier 2001).  Among the X-ray binaries an ideal source to study the transition to the microquasar phase is  LS~I~+61$^{\\circ}$303 because it is the only known periodic microquasar.   \\section {The LS~I~+61$^{\\circ}$303 system}  LS~I~+61$^{\\circ}$303 is the only object of this class showing variations in  X-rays (Leahy 2001), optical wavelenghts in both continuum (Maraschi \\& Treves  1981) and line radiation (Zamanov \\& Mart\\'{\\i}  2000; Liu et al. 2000) and at radio wavelenghts with a period equal to the orbital one; the most accurate value of the orbital period is that resulting from radio observations, equal to 26.496 days (Gregory \\& Taylor 1978; Taylor \\& Gregory  1982; Gregory  2002). The fit performed on near infrared data  by Mart\\'{\\i} and Paredes (1995) produced  high  values  for the  eccentricity (e$\\sim$ 0.7-0.8 ) confirmed by optical observations (Casares et al. 2004). The  lower limit to  $i$, the angle formed by the axis of the orbit and the line of sight is  38$^{\\circ}$ (Hutchings \\& Crampton 1981; Massi et al. 2001; Massi 2004). The  phase at the periastron passage is $\\Phi$=0.2 with the phase  referred to the time t$_0$=JD\\,2443366.775, the date of the first radio detection of the system (Gregory \\& Taylor 1978).  Ultraviolet spectroscopy of LS~I~+61$^{\\circ}$303  by Hutchings \\& Crampton (Hutchings \\& Crampton  1981) indicates that the normal star is a main sequence B0-B0.5 star (L$\\sim 10^{38}$ erg sec$^{-1}$, T$_{eff}\\simeq 2.6~10^4$K). The optical spectrum is that of a rapidly spinning Be star. Together with the usual  high velocity  (1000 km s$^{-1}$) low density wind at high latitudes typical for OB stars, Be stars have a dense and slow ($<$100 km~ s$^{-1}$) disk-like wind around the equator (Waters et al. 1988). Equatorial mass loss, due to an interplay of the high rotation and of internal pulsations of the star,  is highly variable and in  some cases  periodical. This is the case for LS~I~+61$^{\\circ}$303, where periodical variations of the mass loss from the Be star have been proved by a  modulation of  H$\\alpha$ emission line with a period of almost 4 years (Zamanov and Mart\\'{\\i} 2000).   The most reliable method to determine the nature of the compact object is the study, as usual  in binary  systems, of the changing radial velocity of the normal companion during its orbit. The   amplitude  ($K_c$) of the  radial velocity variations and the period (P$_{orb}$) of the system define a quantity, called the ``mass function'' (Charles \\& Wagner 1996), $f$, which depends on the inclination $i$  of the orbit, the masses $M_X$ and $M$ of the accretor and its normal companion: $$f ={P_{orb} K_c^3\\over 2\\pi G}={M_X^3 sin^3 i\\over (M_X +M)^2} $$ where  $G$ is the gravitational constant.  Once the inclination, $i$, and the mass of the companion, $M$, are known one can solve for  $M_X$.  Rhoades \\& Ruffini (1974) by taking the most extreme equation of state that produces the maximum critical mass of  a neutron star, established the upper limit of 3.2 $M\\odot$. This absolute maximum mass  provides a decisive method of observationally distinguishing neutron stars from black holes. The problem with LS~I~+61$^{\\circ}$303 is the large range for the mass function allowed by the parameter uncertainties. Optical observations (Casares et~al. 2004) give the mass function $f$ in the range $0.003<f<0.027$. The upper limit for $f$, with i=38 and  $M = 18\\,M_\\odot$ gives  $M_{\\rm X}<3.8\\,M_\\odot$. Therefore  the nature of the accretor is still  an open issue (Massi 2004).   \\section {\\bf Radio emission}  The greatest  peculiarity of LS I +61$^{\\circ}$303 are its periodic radio outbursts  with P=26.496 days (Gregory 2002). In Fig.\\ref{fig:flusso} a typical radio light curve  is shown. The decay of the outburst agrees with that expected for an adiabatically expanding cloud of synchrotron-emitting relativistic electrons (Taylor \\& Gregory 1984). However,  that  model alone fails to fit the peaks at  different frequencies: the flat spectrum during the late portion of the rise in flux density can be reproduced , if together with the adiabatic expansion losses also  a continuos ejection of particles lasting two days is taken into account (Paredes et al 1991).  \\begin{figure}[htb] \\centering \\resizebox{\\hsize}{!}{\\includegraphics[scale=0.15, angle=0]{massi_fig2.ps}} \\caption{ Radio light curve of LS~I~+61$^{\\circ}$303 obtained with the GBI at 8.4~GHz (Strickman et~al. 1998). The shaded area indicates the time interval during the EVN  observation  of Fig. 4 (figure in  Massi et al. 2001).} \\label{fig:flusso} \\end{figure}  The amplitude of each outburst is not randomly varying, but itself periodic with a periodicity of 4.6 years correlated with  the mass loss of the Be star (sect. 2) (Gregory 1999, 2002; Zamanov \\& Mart\\'{\\i}  2000). Also, the orbital phase $\\Phi$ at which the outbursts occur is modulated (Gregory et al. 1999) and varies   in the interval 0.45--0.95 (Paredes et al. 1990) The  $\\Phi$ at periastron passage is   0.2. Therefore one of the fundamental questions concerning the periodic radio outbursts of \\object{LS~I~+61$^{\\circ}$303}  has been: Why are the radio outbursts shifted with respect to the periastron passage?  \\section {\\bf The two peak accretion model}  In order to explain the association between LS~I~+61$^{\\circ}$303 and the gamma-ray source 2CG 135+01/3EG J0241+6103, Bosch-Ramon and Paredes (2004) have proposed a numerical  model  based on inverse Compton scattering. The  relativistic electrons  in the jet are  exposed to  stellar  photons (external Compton) as well as to synchrotron photons (synchrotron self Compton). The model considers  accretion variations along the orbit and predicts a gamma-ray peak at periastron passage where the accretion is higher. EGRET data show indeed a  gamma-ray peak at  $\\Phi$=0.2 (Fig. 3) and catastrophic inverse Compton losses might explain the absence of radio emission at periastron. Therefore, to explain the observed periodic radio outbursts in the phase interval 0.45--0.95 a second accretion/ejection event must occur.  Taylor et al. (1992) and Mart\\'{\\i} \\& Paredes (1995) have shown that for accretion along an eccentric orbit the accretion rate $\\dot{M} \\propto {\\rho_{\\rm wind}\\over v_{\\rm rel}^3}$, (where $\\rho_{\\rm wind}$ is the density of the Be star wind and $v_{\\rm rel}$ is the relative speed between the accretor and wind)  develops  two peaks: the highest peak corresponds to the periastron passage (highest density), while the second peak occurs when the drop in the relative velocity $v_{\\rm rel}$ compensates the decrease in density (because of the inverse cube dependence). Mart\\'{\\i} \\& Paredes (1995) have shown that during both peaks the accretion rate is above the Eddington limit and therefore one expects that matter is ejected twice within the 26.496 days interval. Mart\\'{\\i} \\& Paredes  have found that variations of Be star wind velocity produce a variation in the orbital phase of the second peak. At this second accretion peak the compact object is  far enough away from the Be star, so that the inverse Compton losses are small and electrons can propagate out of the orbital plane. Then an expanding double radio source should be observed, which in fact has been observed by VLBI and MERLIN. \\begin{figure} \\centering \\includegraphics[angle=-90,width=9.5cm]{massi_fig3.ps} \\caption{ EGRET  data (Tavani et al. 1998) folded with the orbital/radio period $P=26.4960$ days. The plot begins at periastron passage $\\Phi\\simeq$0.2 and shows the follow-up of the gamma-ray emission  along one full orbit. At  epoch 2450334JD (i.e. circles in the plot, with empty circles indicating  upper limits) the orbit has been well sampled at all phases: A clear peak is centered at periastron passage 0.2 and 1.2. At a previous epoch (2449045JD; triangles in the plot, with empty triangles indicating  upper limits) the sampling is uncomplete,  still the data show an increase  toward  ($\\Phi \\simeq $0.3) periastron passage and a peak at $\\Phi\\simeq$0.5. The 3 squares refer to a third epoch (2449471JD). The emission is suggested to be   produced via inverse Compton scattering of stellar photons (Taylor et~al. 1992, 1996) and of synchrotron photons (Bosch-Ramon \\& Paredes 2004) by the relativistic electrons of the jet. The ejection is  predicted to occur twice along the orbit, one always at periastron passage and the second at a varing orbital phase (Mart\\'{\\i} and Paredes 1995). Relativistic electrons ejected at $\\Phi$=0.2 suffer severe Compton losses: radio outbursts  never occur at periastron. Radio outbursts occur in the orbital phase interval 0.45-0.95 (Paredes et al. 1990). The gamma-ray peak at $\\Phi\\simeq0.5$  could be associated to a such  second ejection. However, no radio data are available to calculate energy budget/losses of the relativistic electrons (figure in  Massi et al. 2004b). } \\label{FigVibStab} \\end{figure} ",
        "conclusions": "The main points of our review  on the characteristics of the  source LS~I~+61$^{\\circ}$303 in the frame of the two-peak accretion/ejection model are:  \\begin{enumerate} \\item It is still an open issue whether the compact object in this system is a neutron star or a black hole. In fact, taking into account the uncertainty in inclination, mass of the companion and the mass function, the existence of a black hole cannot be ruled out.  \\item The  radio jet  at a scale of hundreds of AU   quite strongly changes its morphology in  short intervals (within 24 hours),  evolving from an initial double-sided jet into  an one-sided jet. This  variation corresponds to a reduction of more than 10$^{\\circ}$ in the angle between the jet and the line of sight. This new alignment severely Doppler de-boosts the counter-jet. Further observational evidence for a precessing jet is recognizable even  at AU scales.  \\item The same population of relativistic electrons emitting radio-synchrotron radiation may  upscatter - by inverse Compton processes - ultraviolet stellar photons  and produce gamma-ray emission. For ejections at the periastron passage gamma-ray flares are expected, but because of severe Compton losses no radio flares, as indeed the data seem to indicate.  \\end{enumerate}  We conclude that  as precession and variable doppler boosting are   the causes of  the  rapid change in the radio-morphology, precession and variable  doppler boosting  are  likely to produce    gamma-ray variations at short time scales. The amplification due to the Doppler factor for Compton scattering of stellar photons by the relativistic electrons of the jet is $\\delta^{3-2\\alpha}$ (where $\\alpha<0$), and therefore  higher than that for synchrotron emission, i.e. $\\delta^{2-\\alpha}$ (Kaufman Bernad\\'o et~al. 2002). LS~I~+61$^{\\circ}$303 becomes therefore the ideal laboratory to test the recently proposed model for microblazars with INTEGRAL and MERLIN observations now and by AGILE and GLAST in the future."
    },
    "0410/astro-ph0410028_arXiv.txt": {
        "abstract": "We investigate how the hierarchical merging of dark matter halos, the radiative cooling of baryons, and the energy feedback from supernovae and active galactic nuclei or quasars combine to govern the amount and the thermal state of the hot plasma pervading groups and clusters of galaxies. We show that, by itself, supernova preheating of the external gas flowing into clusters falls short of explaining the observed X-ray scaling relations of the plasma luminosity $L_X$ or the plasma entropy $K$ versus the X-ray temperature $T$. To account for the scaling laws from rich to poor clusters takes preheating enhanced by the energy input from active galactic nuclei. In groups, on the other hand, the internal impacts of powerful quasars going off in member galaxies can blow some plasma out of the structure. So they depress $L_X$ and raise $K$ to the observed average levels; meanwhile, the sporadic nature of such impulsive events generates the intrinsic component of the wide scatter apparent in the data. The same quasar feedback gives rise in groups to entropy profiles as steep as observed, a feature hard to explain with simple preheating schemes. Finally, we argue a close connection of the $L_X-T$ or the $K-T$ relation with the $M_{\\bullet}-\\sigma$ correlation between the host velocity dispersion and the masses of the black holes, relics of the quasar activity. ",
        "introduction": "% The hot medium pervading many single galaxies and most groups and clusters shines in X rays by thermal bremsstrahlung and line emission, see Sarazin (1988). Simple conditions are found to prevail in rich clusters.  These emit huge powers $L_X\\propto n^2\\,\\sqrt{T}\\, R^3\\sim 10^{44}-10^{45}$ erg s$^{-1}$ in X rays; the temperatures $kT\\approx 5$ keV, measured from the continuum and from high excitation lines, are close to the virial values $kT_v\\approx G\\, M\\, m_p/10\\, R$ in the gravitational wells mainly provided by dark matter (DM) masses $M \\approx 10^{15}\\, M_{\\odot}$ within sizes $R$ of a few Mpcs \\footnote{Throughout the paper $k$ is the Boltzmann constant, $G$ is the gravitational constant, $m_p$ is the proton mass, and $e$ is the electron charge.}. The inferred gas number densities decline outward from central values $n\\approx 10^{-3}$ cm$^{-3}$; so this medium with low $n$ and high $T$ satisfying $kT/ e^{2}\\, n^{1/3}\\sim 10^{12}$ constitutes the best ion-electron plasma in the Universe ever, the intracluster \\emph{plasma} or ICP.  Such a medium is apparently {\\it simple} on the following accounts. Microscopically, it is constituted by pointlike particles in thermal equilibrium. At the macroscopic end, the overall baryonic fractions resulting from the ICP densities and radial distributions inventoried in many clusters (White et al. 1993, Allen \\& Fabian 1994) come to values $m/M \\approx 0.16$; this is close to the cosmic ratio $\\Omega_b/\\Omega_M$ of baryons to DM obtained in the current Concordance Cosmology from the parameters $\\Omega_b\\approx 0.044$ and $\\Omega_M\\approx 0.23$ (see Bennett et al. 2003) \\footnote{This will be used throughout the paper, with its additional parameters $\\Omega_{\\Lambda}\\approx 0.73$ for the dark energy density, and $H_0\\approx 70\\, \\mathrm{h}_{70}$ km s$^{-1}$ Mpc$^{-1}$ for the Hubble constant.}. In addition, the chemical composition is reasonably constant from cluster to cluster, and close to one-third of the solar value (see Matteucci 2003).  However, surprises arise in moving from rich clusters toward poor groups. In fact, similarly simple conditions holding in the intragroup plasma (IGP) would imply that the luminosities retain the gravitational scaling $L_X\\propto T^2$ (Kaiser 1986). This would apply if the IGP passively followed the DM evolution, and retained the key cluster behaviors: $m/M\\approx$ const, i.e., densities $n$ proportional to the DM mass density $\\rho$; and temperatures $T$ close to the virial value $T_v\\propto M/R\\propto M^{2/3}\\, \\rho^{1/3}$.  Instead, the luminosities recently detected or revised (Horner 2001; O' Sullivan, Ponman \\& Collins 2003; Osmond \\& Ponman 2004) are lower by factors of $10^{-1}$ to $10^{-2}$, see Fig.~1. The figure also shows how the emissions from poor groups and large galaxies scatter widely and often downward, a feature of a largely intrinsic nature (Mushotzky 2004).  So in such smaller structures the plasma is surprisingly underluminous and hence \\emph{underdense}. This is an even more surprising result, considering that in the standard hierarchical cosmogony (see Peebles 1993) such earlier condensations ought to be denser, if anything. Moreover, for $kT < 2$ keV the pinch of highly excited metals contributes important line emissions that imply a flatter $L_X\\propto T$, if anything. How the observed steep decrease may come about constitutes a widely debated issue.  Our proposal centers on the energy gained or lost by the baryons through several processes: the gravitational heating driven by the merging events that punctuate the hierarchical growth of DM condensations (``halos'', see Peebles 1993); the radiative cooling of the baryons; and the energy fed back to baryons when they partly condense within galaxies into massive stars then exploding as Type II supernovae (SNe), or accrete onto a central supermassive black hole (BH) energizing an active galactic nucleus (AGN) or a quasar.  The paper is organized as follows. In \\S~2 we use a telling quantity, the plasma entropy, to show that energy feedback from astrophysical sources is needed to explain its high levels in poor clusters and groups. In \\S~3 we show that preheating by SNe alone is not enough. In \\S~4 we consider preheating enhanced by AGNs, compute the resulting X-ray scaling relations, and critically discuss the outcomes of this approach. In \\S~5 we also consider the internal impacts from quasars; we compute their effect on the X-ray scaling relations and on the entropy profiles in groups, and show how these solve the shortcomings of all external preheating scenarios. In \\S~6 we highlight and discuss the main features of our comprehensive approach.  Auxiliary computations and derivations are given in the Appendices. In Appendix A we derive Eq.~(2) of the main text from the hydrostatic equilibrium, and give handy approximations. In Appendix B we reformulate the classic Rankine-Hugoniot jump conditions in a general form that is also valid for accretion shocks/layers, and derive Eqs.~(5) and (6) of the main text. In Appendix C we develop a new family of self-similar hydrodynamic solutions describing the blast waves driven by the internal impacts of quasars, and extensively used in \\S~5. ",
        "conclusions": "% In this paper we have used pilot semianalytic modeling to show that the energy fed back to baryons by AGNs and quasars is essential to fit the recent X-ray data. We find that the AGN external preheating dominates over SN's to yield the scaling laws $L_X\\propto T^3$ and $K\\propto T^{2/3}$ related by Eq.~(1), that constitute fitting trends in clusters if still marginal for groups. But in groups and in large galaxies the quasar impulsive feedback acting from inside takes over, to eject some plasma and further \\emph{depress} $L_X$, while \\emph{enhancing} $K$ and originating \\emph{non-isentropic} profiles.  We stress that the energy added by AGNs plays an \\emph{inverse} role in preheating and in ejection/outflow. This is because in moving from clusters to groups the ratio $\\Delta E/E$ of the added vs. the equilibrium energy is bound to increase. This causes relatively higher external preheating, warmer inflows and \\emph{weaker} accretion shocks; on the other hand, it drives \\emph{stronger} internal blasts causing more mass ejection/outflow. In parallel, the leading shocks of the blasts replace the accretion shocks in the role of increasing the outer entropy production, so originating comparably steep entropy profiles.  Such internal effects, however, are to \\emph{saturate} in galaxies because the values of $\\Delta E/E$ there are limited, lest the impulsive quasar feedback ejects so much of the surrounding baryons as to stop the BH accretion altogether (see Silk \\& Rees 1998); the saturation is what in our calculation yields the lower \\emph{elbows} of the solid and dashed lines presented in Fig.~3. As argued in Cavaliere et al. (2002), the system constituted by the BH and the surrounding baryons self-regulates to the verge of unbinding; the condition $\\Delta E\\approx E$ directly yields $M_{\\bullet}\\approx$ $5\\times 10^8\\, M_{\\odot}\\, (f/5\\times 10^{-2})^{-1}\\, (\\sigma/ 300\\, \\mathrm{km\\, s^{-1}})^5$ in terms of the DM velocity dispersion $\\sigma$. In turn, the latter is found to correlate less than linearly with the velocity dispersion $\\sigma_{\\star}$ of the host galactic bulge (Ferrarese 2002; Baes et al. 2003; A. Pizzella et al. 2005, in preparation); so the BH mass approaches $M_{\\bullet}\\propto \\sigma_{\\star}^4$.  Thus for the \\emph{same} value $f\\approx 5\\times 10^{-2}$ indicated by the average X-ray data for groups we agree (within the observed scatter, see Tremaine et al. 2002) with the observations of the uppermost relic BH masses in the bulges of many local and currently inactive galaxies. Similar BH masses may be also contributed by an initial, supercritical accretion phase (as discussed by King 2003) launching the piston that in turn drives the far-reaching blasts described above. Subsequently, our outgoing self-similar blasts with $\\Delta E(t)/E(<R_s)\\approx$ const stay tuned to the condition $\\Delta E\\approx E$. Specifically, for $\\omega\\rightarrow 2.5$ not only $E(<R_s)\\propto R_s^{5-2\\, \\omega}\\rightarrow$ const holds but also $\\Delta E\\propto t^{2\\, (5-2\\, \\omega)/\\omega}\\rightarrow$ const applies, consistent with fading quasar output; thus at most limited increase of BH mass $M_{\\bullet}$ takes place.  The empirical fact (see Wandel 2002; Vestergaard 2004) that a similar $M_{\\bullet}-\\sigma$ correlation appears to hold also for the currently active and faraway BHs is consistent with our adoption of comparable values for $M_{\\bullet}$ energizing both modes of nuclear activity: the impulsive quasar feedback effective for plasma ejection from groups and galaxies, and the smoother, long-lived AGN outputs sufficient to preheat the gas falling into poor clusters. We are also consistent with the rough equality of the mass densities in BHs derived from the powerful emissions of the quasars in the optical band (see Marconi et al. 2004), and from the weaker and later AGN activity detected mainly in X-rays (Hasinger 2004; Fabian 2004a).  On the other hand, strong and rare (i.e., increasingly sporadic) quasar impacts can also explain the \\emph{scatter} of the X-ray data widening toward smaller systems as poor groups or massive galaxies (Mushotzky 2004). This we trace back to the increasing \\emph{variance} in the occurrence of strong quasar events or even in their coupling level $f$, that concur to dynamically modulate the plasma ejection $\\Delta m/m\\propto f\\, M_{\\bullet}$ and nonlinearly affect $L_X\\propto (1-\\Delta m/m)^2$ when $\\Delta m/m$ approaches $1$. As the hierarchical clustering proceeds toward clusters, instead, the evolution of the quasars cuts down most internal effects; this is because the impulsive contributions to $\\Delta E$ within a structure's dynamical time hardly can keep pace with the increase of the equilibrium energy $E\\propto m\\, T_v$ in such late and massive systems with deep potential wells.  We stress that our upper and lower bounds for $L_X$ illustrated in Fig.~3 by the SN strip and the quasar line comprise nearly \\emph{all} data points, except for a few groups with peculiar features currently under scrutiny (see Mushotzky 2004; Osmond \\& Ponman 2004). So we submit that several pieces of data fit together when considering both the external preheating from AGNs and the internal impact from quasars, with the same average values of $f\\, M_{\\bullet}$. We remark that several authors (e.g., Ruszkowski \\& Begelman 2002; Fabian 2004b; Zanni et al. 2004) have argued the relevance of AGN feedback in explaining the puzzle posed by the ``cool cores'' (Molendi \\& Pizzolato 2001) at the very centers of many clusters. On the other hand, Cavaliere et al. (2002) and Granato et al. (2004) have stressed that powerful quasar impacts in massive spheroids easily quench star formation and produce precociously red giant elliptical galaxies.  To conclude, we stress that energy feedback from AGNs and quasars with an overall coupling around $5\\times 10^{-2}$ to the ambient baryons yields agreement with independent observations in different frequency bands and over different distance scales. Specifically, this paper is focused on the extended X-ray emissions and plasma entropy of poor clusters and groups; but we have also considered at galactic and subgalactic scales the mainly optical correlation of nuclear BH masses vs. host velocity dispersions. At the intermediate scales of early massive galaxies and in the $\\mu$wave/submm band, we have proposed in Lapi et al. (2003) how to catch quasar impacts in the act from resolved Sunyaev-Zel'dovich signals enhanced by overpressure in running blast waves. Such transient events sweeping plasma outward to lower densities (see Figs.~5 and C1) hardly increase the extended X-ray emissions; they instead specifically correlate with pointlike X-rays from a fully active quasar, and/or with strong IR emissions signaling a nascent quasar enshrouded by dust."
    },
    "0410/astro-ph0410734_arXiv.txt": {
        "abstract": "Observations have suggested substantial departures from pressure equilibrium in the interstellar medium (ISM) in the plane of the Galaxy, even on scales under 50~pc.  Nevertheless, multi-phase models of the ISM assume at least locally isobaric gas.  The pressure then determines the density reached by gas cooling to stable thermal equilibrium.  We use numerical models of the magnetized ISM to examine the consequences of supernova driving for interstellar pressures.  In this paper we examine a (200~pc)$^3$ periodic domain threaded by magnetic fields.  Individual parcels of gas at different pressures reach widely varying points on the thermal equilibrium curve: no unique set of phases is found, but rather a dynamically-determined continuum of densities and temperatures.  A substantial fraction of the gas remains entirely out of thermal equilibrium. Our results appear consistent with observations of interstellar pressures.  They also suggest that the high pressures observed in molecular clouds may be due to ram pressures in addition to gravitational forces.  Much of the gas in our model lies far from equipartition between thermal and magnetic pressures, with ratios ranging from 0.1 to $10^4$ and ratios of uniform to fluctuating magnetic field of 0.5--1.  Our models show broad pressure probability distribution functions with log-normal functional forms produced by both shocks and rarefaction waves, rather than power-law distributions produced by isolated supernova remnants. The width of the distribution can be described quantitatively by a formula derived from the work of Padoan, Nordlund, \\& Jones (1997). ",
        "introduction": "Theoretical models of the interstellar medium (ISM) have generally followed Spitzer (1956) in assuming that the interstellar gas is in pressure equilibrium.  Field, Goldsmith \\& Habing (1969; hereafter FGH) demonstrated that the form of the interstellar cooling and heating curves for temperatures below about $10^4$~K allowed a range of pressures in which two isobaric phases could exist in stable thermal equilibrium.  Although the details of this model turned out to be incorrect, due to their assumption of a cosmic ray flux much higher than subsequently observed, Wolfire et al.\\ (1995) demonstrated that photoelectric heating of polycyclic aromatic hydrocarbons or other very small dust grains could recover a two-phase model for the neutral ISM.  McKee \\& Ostriker (1977; hereafter MO77) used a similar framework for the cold ISM, but incorporated the suggestion by Cox \\& Smith (1974) that supernovae (SNe) would create large regions of hot gas with $T\\sim 10^6$~K, to produce a model of a three-phase medium.  They were the first to relax the assumption of global pressure equilibrium, noting that SN remnants (SNRs) would have varying pressures.  They assumed only local pressure equilibrium at the surfaces of clouds in order to determine conditions there.  Measurements of the ISM pressure were performed using {\\em Copernicus} observations of ultraviolet (UV) absorption lines from excited states of C~{\\sc i} by Jenkins \\& Shaya (1979) and Jenkins, Jura \\& Loewenstein (1983).  They found greater than an order of magnitude variation in pressures in the cold gas traced by the C~{\\sc i}, with most of the gas having pressures $P/k < 10^4$~K~cm$^{-3}$, but a small fraction reaching pressures $P/k > 10^5$~K~cm$^{-3}$.  These findings have recently been confirmed and extended using {\\em Space Telescope Imaging Spectrograph} (STIS) observations by Jenkins \\& Tripp (2001).  Bowyer et al.\\ (1995) \\& Bergh\\\"ofer et al.\\ (1998) compared average pressures derived from observations with the {\\em Extreme UV Explorer} of shadows cast by clouds of neutral hydrogen within the Local Bubble with the pressure of the warm cloud surrounding the Sun and found a difference of as much as a factor of 25.  McKee (1996) argued that these estimates were too extreme, but nevertheless concluded that the observations showed a pressure variation of at least a factor of five.  In this series of papers we show that dynamical models of a SN-driven interstellar medium do not produce an isobaric medium controlled by thermal instability, but rather a medium with a broad pressure distribution controlled by the dynamics of turbulence, with rarefaction waves being as important as shocks in setting local pressures.  Although the cooling curve may determine local behavior, the pressure of each parcel is set dynamically by flows primarily driven by distant SNe.  The theory of the probability distribution function (PDF) of compressible turbulence described by Padoan, Nordlund \\& Jones (1997; hereafter PNJ97) seems able to describe the pressure PDFs we find here.  Gazol et al.\\ (2001) suggest that turbulence driven by ionization heating results in nearly half of the gas lying in thermally unstable regions, in agreement with the observations of Heiles (2001).  Earlier dynamical models of the ISM have been reviewed by Mac Low (2000). Models including SN driving were first done in two dimensions by Rosen \\& Bregman (1995), and have since been done by a number of other groups. V\\'azquez-Semadeni, Passot \\& Pouquet (1995) briefly discussed the pressure structure in a two-dimensional hydrodynamical model including only photoionization heating.  In three dimensions, Avillez (2000) used an adaptive mesh refinement code to compute the evolution of a $1\\times 1\\times 20$~kpc vertical section of a galactic disk, while Korpi et al.\\ (1999) used a single-grid magnetohydrodynamical (MHD) code to model an SN-driven galactic dynamo. Slyz et al.\\ (2005) used a single-grid hydrodynamical code to compute supernovae in a periodic box including self-gravity and a star-formation recipe. Both of these papers briefly discuss pressure structure as well. All of these computations show that the interactions between SNRs drive turbulent flows throughout the ISM, and that radiative cooling of compressed regions produces cold, dense clouds with relatively short lifetimes.  We here study in detail the pressure distribution in MHD models of a region in the plane of the SN-driven ISM, while in a subsequent paper we will study the pressure distribution in a stratified ISM.  We use the Riemann framework of Balsara (2000) in single-grid mode on a 200~pc cube with periodic boundary conditions on all sides, including both radiative cooling and heating, as well as isolated SNs, though neither external or self-gravity.  In \\S~\\ref{sec:models} we describe the numerical model.  We then give three different analytic approaches to the question of the pressure distribution in the SN-driven ISM in \\S~\\ref{sec:an-th}, derived from the work of MO77, PNJ97, and Passot \\& V\\'azquez-Semadeni (1998; hereafter PV98).  In \\S~\\ref{sec:numerics} we use our numerical results to demonstrate that the SN-driven ISM is far from isobaric, with order of magnitude pressure variations, but that the distribution of pressures takes on a log-Gaussian form that can be well described. These results are compared to observations in \\S~\\ref{sec:observations}, and the paper is summarized in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have examined the distribution of pressures predicted from three-dimensional simulations of a SN-driven, magnetized ISM, neglecting vertical stratification and self-gravity, but including a distributed heating function.  In all the simulations we have run, we find pressures distributed over rather more than an order of magnitude around the mean.  The pressure PDF can be fit well with a log-normal having width given by the analytic description of PNJ97 based on three-dimensional simulations of isothermal turbulence, and disagreeing with the one-dimensional results of PV98.  The log-normal distribution is predicted by the theory of V\\'azquez-Semadeni (1994) that density fluctuations in a supersonic turbulent flow result from a sequence of random compressions and rarefactions. We find that compressions and rarefactions are indeed symmetrically distributed, lending further support to this idea.  The log normal form differs from the power-law distribution predicted by MO77. The only isobaric regions found in our models are the interiors of young SNRs; these are also the only laminar flow regions in an otherwise turbulent environment.  Most of the mass is in the turbulent gas, however, as is most of the volume, even in models driven with four times the Galactic supernova rate.  This result is consistent with the results of Korpi et al.\\ (1999), Avillez (2000), and Avillez \\& Breitschwerdt (2004).  The last group performed a resolution study to establish that numerical diffusion was not causing hot gas to cool unphysically.  We will study the filling factors of our models in more detail in future work.  We find a broad range of pressures, and a substantial fraction of associated densities, far from the thermal equilibrium values.  This limits the predictive usefulness of phase diagrams based on thermal equilibrium, although thermal equilibrium at the local pressure is still the mildly favored state.  Gas pressures appear to be determined dynamically.  Each individual parcel of gas seeks local thermal equilibrium at the pressure imposed on it by the turbulent flow. As suggested by previous authors, the phase diagram is only locally valid.  Isobaric thermal instability in such an environment will lead to regions of the phase diagram being mildly disfavored, but no more. V\\'azquez-Semadeni et al.\\ (2000) and Gazol et al.\\ (2001) used two-dimensional simulations with more physically-motivated cooling curves at lower temperatures to conclude that the effects of the thermal instability will be barely visible in the overall probability distribution functions (but note that these simulations did not resolve cooling regions well, so the results should be taken with caution).  Our results appear consistent with observations that have repeatedly shown the ISM not to be isobaric, including those by Jenkins \\& Shaya (1979), Jenkins, Jura \\& Loewenstein (1983), Bowyer et al.\\ (1995), Bergh\\\"ofer et al.\\ (1998), and Jenkins \\& Tripp (2001).  Although heating and cooling rates remain important for determining the local density and temperature, they do not produce a global multi-phase medium with well-separated regions of different temperature gas because of the wide range of pressures present and the dynamical processing of the gas.  Rather, a more continuous distribution of temperatures and densities is always present.  The inference that most molecular clouds must be gravitationally bound because of their high observed confinement pressures are called into question by our results.  Regions with densities approaching the overall densities of GMCs, and pressures an order of magnitude above the average interstellar pressure appear in our simulations even in the absence of self-gravity.  This supports suggestions that star-forming molecular clouds may be transient, turbulently-driven, objects (Ballesteros-Paredes et al.\\ 1999ab, Elmegreen 2000, 2002, Hartmann et al.\\ 2001). Objects that do gravitationally collapse from large scales as described, for example by Kim \\& Ostriker (2001), Elmegreen (2002), or Li et al.\\ (2005) will also form molecular gas, but will quickly form starburst knots in a burst of violent, unimpeded star formation.  Magnetic pressures are also broadly distributed, and do not correlate well with thermal pressures.  The ratio of thermal to magnetic pressure $\\beta$ takes on values ranging from 0.1 to as much as 100 in gas with temperatures under $10^4$~K, and much higher values in hot gas. The ratio of fluctuating to mean field in our models $\\langle \\delta B\\rangle / \\langle B \\rangle=0.5$--1, high enough to suggest that field strengths derived from observations of polarized synchrotron emission may be overestimated (Beck 2001)."
    },
    "0410/astro-ph0410214_arXiv.txt": {
        "abstract": "We review our current knowledge of the solar chemical composition as determined from photospheric absorption lines. In particular we describe the recent significant revisions of the solar abundances as a result of the application of a time-dependent, 3D hydrodynamical model of the solar atmosphere instead of 1D hydrostatic models. This has decreased the metal content in the solar convection zone by almost a factor of two compared with the widely used compilation by Anders \\& Grevesse (1989). While resolving a number of long-standings problems, the new 3D-based element abundances also pose serious challenges, most notably for helioseismology. ",
        "introduction": "The chemical composition of the Sun is one of the most important yardsticks in astronomy with implications for basically all fields from planetary science to the high-redshift Universe. Standard practise is to compare the element content in a cosmic object with the corresponding value for the Sun using the normal logarithmic abundance scale $[{\\rm X/H}] \\equiv  {\\log} \\left(N_{\\rm X}/N_{\\rm H}\\right)_* - {\\log} \\left(N_{\\rm X}/N_{\\rm H}\\right)_\\odot \\equiv {\\log} \\epsilon_{\\rm X, *} - {\\log} \\epsilon_{\\rm X, \\odot}$. Abundance ratios such as [C/Fe] are defined correspondingly. Hydrogen is the chosen reference element since it is the most common element in the Universe as a whole and that it normally dominates the continuous opacity in the formation of the stellar spectra. The solar abundances are, however, not only of interest as a reference point in astronomy but are also of profound importance for the understanding of our own solar system and the solar interior.  Most of our knowledge of the overall abundance distribution of the elements in the solar system originates in two quite different sources: the solar photospheric spectrum and pristine meteorites, in particular the so-called C1 chondrites\\footnote{Important and complementary information is provided by studies of for example the solar corona, solar wind, solar energetic particles, solar flares and lunar soil, which however will not be discussed here.} (e.g. Anders \\& Grevesse 1989). Each has their own advantages and disadvantages. The abundances of a very large number of elements can be determined using absorption lines in the solar spectrum. However, since the results are dependent on having realistic models of the solar atmosphere and the line formation process the accuracy can only be expected to be good at the 10\\% (0.04\\,dex) level at best. Isotopic abundance information is only available in very rare instances. Model-independent meteoritic isotopic abundances can be measured to exquisite accuracy in the laboratory but only a very small subsample of meteorites do not show effects of fractionation. In addition, elements like H, He, C, N, O and Ne are all volatile and hence depleted in the meteorites. It is therefore clear that the two methods complement each other well. The present review will describe our knowledge of the solar photospheric chemical composition, in particular in light of recent significant revisions of these abundances following the application of the new generation of three-dimensional (3D) hydrodynamical model atmospheres to the problem. ",
        "conclusions": "Although progress has been made in refining our knowledge of the solar chemical composition, much work still remains. An obvious way forward is to extend the 3D analysis presented here to the remaining elements in the periodic table. Furthermore, the lack of detailed non-LTE line formation calculations is a major cause of concern, not only for the solar abundance determinations but also for much of the stellar work. Equally important, however, is continued investment in improving the necessary atomic and molecular data. Although the impact may not be as profound as the recent surprisingly large revisions of the solar metal content, efforts along these lines would nevertheless be worthwhile."
    },
    "0410/astro-ph0410022_arXiv.txt": {
        "abstract": "Multi-wavelength observations of nearby spiral galaxies have shown that neutral and ionized gas are present up to a few kpc from the disk and that star formation and supernovae probably play an important role in bringing gas into the halo. We have obtained very sensitive \\HI\\ observations of the face-on galaxy \\mbox{NGC 6946} and of the nearly edge-on starburst galaxy \\mbox{NGC 253}. We find high velocity \\HI\\ clouds in \\mbox{NGC 6946} and extra-planar gas with anomalous velocities in \\mbox{NGC 253}. In both galaxies there seems to be a close connection between the star-forming disk and the halo \\HI. In the outer parts of \\mbox{NGC 6946} there is also evidence for recent gas accretion. ",
        "introduction": "\\HI\\ has been detected in the halos of several spiral galaxies. Studies of the edge-on galaxies \\mbox{NGC 891} \\citep{swa97, fra03} and \\mbox{UGC 7321} \\citep{mat03}, and the intermediate inclination galaxy \\mbox{NGC 2403} \\citep{sch00,fra03} have revealed a thick layer of \\HI\\ rotating more slowely than the underlying disk (see Oosterloo et al., this conference, for an overview of these studies). The origin of such a thick layer or halo is unknown. There are various possibilities. One of them is the galactic fountain mechanism according to which active star formation in the disk and supernovae are responsible for blowing the gas into the halo.\\\\ The galaxies mentioned above are seen under suitable orientations for a study of the vertical \\HI\\ distribution and the general kinematics. Vertical motions are instead best observed in face-on galaxies.\\\\ For this purpose we have obtained new, deep \\HI\\ observations with the Westerbork Synthesis Radio Telescope (WSRT) of \\mbox{NGC 6946}, a nearby face-on spiral galaxy showing moderate star formation activity.\\\\ In addition we have chosen \\mbox{NGC 253} as a case of a galaxy with an extreme starburst, in order to trace the effects on the \\HI\\ in such a limiting case. For this galaxy we have used \\HI\\ data from the Australia Telescope Compact Array (ATCA).\\\\ ",
        "conclusions": "\\mbox{NGC 6946} and \\mbox{NGC 253} clearly are two more examples of galaxies with extra-planar \\HI. Yet no consistent picture emerges and a number of processes may be at work. In both cases the data suggest a connection with the star formation in the disk. In \\mbox{NGC 6946} this is more moderate but present over the entire disk. In \\mbox{NGC 253} this is more violent, especially in the inner region. It is not clear why the amount of gas in the halo are so different in the two galaxies. In the outer parts of \\mbox{NGC 6946} we also find evidence for gas accretion.\\\\"
    },
    "0410/astro-ph0410508_arXiv.txt": {
        "abstract": "Phantom dark energy models, with $w < -1$, are characterized by a future singularity and therefore a finite lifetime for the universe.  Because the future singularity is triggered by the onset of dark-energy domination, the universe spends a significant fraction of its total lifetime in a state for which the dark energy and matter densities are roughly comparable. We calculate, as a function of $w$, the fraction of the total lifetime of the universe for which the dark energy and matter densities differ by less than the ratio $r_0$ in either direction.  For $r_0$ = 10, this fraction varies from $1/3$ to $1/8$ as $w$ varies from $-1.5$ to $-1.1$; the fraction is smaller for smaller values of $r_0$.  This result indicates that the coincidence problem is significantly ameliorated in phantom-dominated cosmologies with a future singularity. ",
        "introduction": "The universe appears to consist of approximately 30\\% nonrelativistic matter, including both baryons and dark matter, and 70\\% dark energy (see Ref. \\cite{Sahni} for a recent review, and references therein).  The matter has a density $\\rho_M$ that scales with the scale factor $R$ as $\\rho_M \\propto R^{-3}$.  The evolution of the dark energy density depends on its equation of state, which is usually parametrized in the form \\begin{equation} p_{DE} = w \\rho_{DE}, \\end{equation} where $p_{DE}$ and $\\rho_{DE}$ are the pressure and density of the dark energy. Then the density of the dark energy scales as \\begin{equation} \\label{rhoDE} \\rho_{DE} \\propto R^{-3(1+w)}. \\end{equation} The simplest model for the dark energy is a cosmological constant, for which $w = -1$ and $\\rho_{DE} = constant$. More complex models have been proposed, in which the dark energy arises from a scalar field; these are called quintessence models \\cite{ratra,turner,caldwelletal,liddle,zlatev}.  These models generally give rise to a time-varying $w$ and more complex behavior for $\\rho_{DE}$.  However, all of these models have a similar problem.  Since the matter density decreases more rapidly than the dark energy density, we expect $\\rho_M \\gg \\rho_{DE}$ at times much earlier than the present, while in the far future, $\\rho_{DE} \\gg \\rho_M$.  Thus, we live in a special epoch, when the matter and dark energy densities are roughly comparable.  This has been dubbed the ``coincidence problem.\"  It is possible, of course, that this coincidence is, in fact, just a coincidence, and no deeper explanation is possible. Nonetheless, numerous attempts have been made to provide an explanation for this coincidence. For example, in a certain class of quintessence models with non-standard kinetic terms (``$k$-essence\"), the dark energy density tracks the radiation energy density during the radiation-dominated era, but then evolves toward a constant-density dark energy component during the matter-dominated era \\cite{Arm2,Arm3}, so the coincidence problem is resolved by linking the onset of dark energy domination to the epoch of equal matter and radiation.  (Although see Malquarti et al. \\cite{Malquarti1}, who argued that the basin of attraction for models exhibiting the desired behavior is quite small).  Another possible solution lies with models in which the periods of matter domination and dark energy domination alternate, either because the dark energy arises from a single scalar field with oscillatory behavior \\cite{Dodelson,quintom}, or because the periods of dark energy domination arise from different quintessence fields, with a variety of energy scales \\cite{Griest}. Another possibility is to couple the matter and quintessence fields, so that energy can be transferred between them.  With a suitable coupling, a constant ratio of dark energy density to matter density can be arranged \\cite{Chimento1,Chimento2,Huey}.  Finally, anthropic arguments have been put forward to resolve the coincidence problem (see, e.g., Ref. \\cite{Garriga}), but such arguments are, of course, controversial.  The coincidence problem is significantly ameloriated in phantom dark energy cosmologies, first proposed by Caldwell \\cite{Caldwell}.  Such models, with $w < -1$, can lead to a future singularity, or ``cosmic doomsday\" \\cite{rip,rip2}. (Note that Starobinsky \\cite{Star} previously considered similar models and calculated the time to the future singularity).  The existence of this future singularity is often considered a negative feature of phantom dark energy models, so a great deal of effort has gone into constructing models with $w<-1$ that avoid a future singularity \\cite{Diaz,Sahni2,Nojiri1,Nojiri2,Nojiri3,Woodard,Yu}.  Here we examine a positive feature of phantom dark energy models with a future singularity, previously mentioned in passing by McInnes \\cite{mcinnes2}: since such models produce a finite lifetime for the universe, it is possible to calculate the fraction of time that the universe spends in a ``coincidental state\", with $\\rho_{M} \\approx \\rho_{DE}$.  Because the final singularity is triggered by the onset of dark energy domination, this fraction is generically nonnegligible.  Here we make this argument quantitative.  In the next section, we examine the time evolution of phantom cosmologies, and calculate the fraction of the total lifetime of the universe for which the matter and dark energy densities are roughly similar.  Our results are discussed in Sec. III. ",
        "conclusions": "Our results indicate that a universe with phantom dark energy spends a significant fraction of its total lifetime in a state for which $\\rho_{M}$ and $\\rho_{DE}$ differ by less than an order of magnitude.  Thus, if we simply found ourselves living at a random moment in such a universe, it would not be surprising to find $\\rho_{M}$ and $\\rho_{DE}$ in such a ``coincidental\" state. Clearly, these results do not explain {\\it why} we happen to live in the epoch for which $\\rho_{M} \\approx \\rho_{DE}$; instead, they make this possibility less implausible. These arguments generalize, in an obvious way, to phantom models in which $w$ varies with time.  Note that this is {\\it not} an anthropic argument; it does not rely on the universe being in a state which can support life.  For instance, this argument could also be applied to a model of the universe in which the ratio of the dark energy, matter, and radiation densities was such that the universe passed directly from the radiation-dominated era to the dark-energy dominated era, without an intervening matter-dominated era.  From an anthropic point of view, such a universe is ruled out by the lack of structure formation.  In contrast, the argument presented here would simply indicate that the universe in this case would spend a significant fraction of its lifetime with the radiation and dark energy densities having the same order of magnitude.  Caldwell et al. \\cite{rip} noted that the phantom dark energy cosmology forces us to live in the brief epoch in which bound structures can form, but this again is a different argument from the one we are making here.  These results clearly depend on the fact that a phantom-dominated universe has a finite lifetime.  The only way to calculate an equivalent fraction of time in an infinitely long-lived universe would be if the epochs of dark energy domination and matter domination alternated in a cyclical fashion, as in the models of Refs. \\cite{Dodelson,quintom,Griest}.  In such models one could then determine the probability of living in a ``coincidental epoch\" of the sort in which we presently find ourselves."
    },
    "0410/astro-ph0410352_arXiv.txt": {
        "abstract": "We  have studied the impact of cosmic-ray acceleration in SNR on the spectra of cosmic-ray nuclei in the Galaxy using a series expansion of the propagation equation, which allows us to use analytical solutions for part of the problem and an efficient numerical treatment of the remaining equations and thus accurately describes the cosmic-ray propagation on small scales around their sources in three spatial dimensions and time. We found strong variations of the cosmic-ray nuclei flux by typically 20\\% with occasional spikes of much higher amplitude, but only minor changes in the spectral distribution. The locally measured spectra of primary cosmic rays fit well into the  obtained range of possible spectra. We further showed that the spectra of the secondary element Boron show almost no variations, so that the above  findings also imply significant fluctuations of the Boron-to-Carbon ratio. Therefore the commonly used method of determining CR propagation parameters by fitting secondary-to-primary ratios appears flawed on account of the variations that these ratios would show throughout the Galaxy. ",
        "introduction": "The origin of cosmic-rays (CR) is one of the major problems in modern astrophysics. In that quest two possible strategies may be followed: \\begin{itemize} \\item one may search for high-energy emission from the source regions themselves, for the CR density must be very high therein. Particle acceleration at supernova remnant (SNR) shock waves is regarded as the most probable mechanism for providing Galactic cosmic rays at energies below $10^{15}$ eV on account of the power requirements \\citep{be87}. The recent detections of non-thermal X-ray synchrotron radiation from the four supernova remnants SN1006 \\citep{koyama95}, RX J1713.7-3946 \\citep{koyama97,sla99}, Cas A \\citep{allen97,go01}, and RCW86 \\citep{ba00,bor01,rho02} support the hypothesis that at least CR electrons are accelerated predominantly in SNR. The evidence for CR nucleon acceleration in SNR is at best controversial \\citep{bkv02,reimer02}. \\item one may also construct models for the diffusive propagation of CRs in the Galaxy and compare their predictions with observed CR spectra at earth \\citep[e.g.][]{strong98,maurin01,taillet03}. The relation between primary CRs, which have been accelerated somewhere in the Galaxy, and secondary CRs, which are produced in interactions of CRs with ambient gas, then allows to determine the propagation parameters such as the diffusion coefficient, the size of the diffusion region, and others. Nearly all published studies of that kind assume steady-state conditions. \\end{itemize}  Already in 1969, \\citet{lingenfelter69} found the CR energy density at earth to vary due to nearby SN, given the CR are accelerated at these sites. A few years ago \\citet{pohl98} showed for CR electrons that in case of SN origin the CR electron spectrum would vary in space and time throughout the Galaxy. The purpose of this paper is to present a time-dependent propagation model for CR nucleons. In particular we aim to address the following questions: \\begin{itemize} \\item would a SNR origin of CR nucleons also lead to significant fluctuations of the CR density in the Galaxy, which then would modify secondary-to-primary ratios from their steady-state values? If that was the case, we would have to re-think the way we infer CR propagation parameters from their locally observed spectra. \\item are there any signatures in the CR distribution in the Galaxy, that might permit to infer a SNR origin of CR nucleons on the grounds of locally observed CR spectra and the diffuse Galactic gamma-ray emission? \\end{itemize}  For that purpose we have developed a method to numerically solve the propagation equation based on a series expansion, which allows us to use analytical solutions for part of the problem and an efficient numerical treatment of the remaining equations. As described in detail in section \\ref{sec:computations} this technique allows us to accurately follow the CR propagation on small scales around their sources in three spatial dimensions and time. In section \\ref{sec:signatures1} we discuss the results of these calculations with respect to the CR density distribution and in section \\ref{sec:disskuss} with respect to CR spectra, respectively. ",
        "conclusions": "We  have studied the impact of CR acceleration in SNR on the spectra of CR nuclei in the Galaxy. We found strong evidence that this assumption leads to CR spectra,  which show significant variations in space and time. The behavior of the CR nuclei resembles that of protons, as suggested by first computations \\citep{strong01}, but differs  considerably  from that of CR electrons \\citep{pohl98}.  We have developed a method to numerically solve the propagation equation based on a series expansion, which allows us to use analytical solutions for part of the problem and an efficient distributed computing for the remainder. This method was employed to calculate the CR densities in the Galactic disk with high spatial ($\\Delta z\\,=\\, 20\\,$pc) and temporal ($\\Delta t\\,=\\,1\\,$kyr) resolution for the primary  nuclei $^{16}$O, $^{12}$C and the secondary nucleus $^{11}$B for various distributions of SNR. This method allowed for the first time to calculate the CR densities in the Galactic disk with the high spatial and temporal resolution required to follow the CR density fluctuations caused by an SNR origin. We also studied the impact of a locally enhanced SN rate within Gould's Belt, a nearby star-forming region.  We found strong variations of the CR nuclei flux by typically 20\\% with occasional spikes of much higher amplitude, but only minor changes in the spectral distribution. The locally measured primary CR spectra fit well into the  obtained range of possible spectra. We further showed that the spectra of the secondary element Boron show almost no variations, so that the above  findings also imply significant fluctuations of the Boron-to-Carbon ratio. Therefore the commonly used method of determining CR propagation parameters by fitting secondary-to-primary ratios appears flawed on account of the variations that these ratios would show throughout the Galaxy.  Some indications that the CR flux varies in the Galaxy are given by the observation of the diffuse \\gr emission in our Galaxy. \\citet{digel01} performed \\gr observations of the outer Galaxy. Their analysis of the diffuse \\gr emission from giant molecular clouds in the Monoceros region suggests that the CR flux in the local Galactic arm and the neighboring Perseus arm differ, i.e. the data suggests an enhancement of the CR density in the Perseus arm. These observations would be well explained by the variations in the flux of CR nuclei that results from a SNR origin.  \\citet{hu97} find that the spectrum of the diffuse \\gr emission toward the inner Galaxy can not be explained by the assumption that the locally observed  CR spectrum and electron-to-proton ratio hold throughout the Galaxy. At \\gr energies above some GeV, where the spectrum is presumably dominated by CR-nucleon-induced radiation, they find that the \\gr flux measured exceeds by about 50\\%  that expected if the local CR spectrum and electron-to-proton ratio hold throughout the Galaxy. These findings are { difficult to explain by the calculations presented above as being caused by spectral variations of the CR nuclei flux, for we find these to be rather small. We have, however, neglected the possibility of a dispersion in the source spectral indices, which has been proposed as an explanation of} the GeV excess \\citep[cf.][]{buesching01}."
    },
    "0410/astro-ph0410487_arXiv.txt": {
        "abstract": "The Earth's high energy gamma-ray emission is caused by cosmic ray interactions with the atmosphere. The EGRET detector on-board the CGRO satellite is only the second experiment (after SAS-2) to provide a suitable dataset for the comprehensive study of this emission. Approximately 60\\% of the EGRET dataset consist of gamma photons from the Earth. This conference contribution presents the first results from the first analysis project to tackle this large dataset. Ultimate purpose is to develop an analytical model of the Earth's emission for use in the GLAST project. The results obtained so far confirm the earlier results from SAS-2 and extend them in terms of statistical precision and angular resolution. ",
        "introduction": "The Earth's high-energy gamma-ray emission (the ``gamma-ray albedo'') has been studied very little from space. The most precise measurements so far were obtained in 1972/73 with the SAS-2 satellite \\cite{thompson} but have very limited photon statistics. Other space measurements include \\cite{kraushaar,guryan,galper}.  The gamma-ray albedo is caused by the interaction of cosmic rays and solar wind with the Earth's atmosphere. The study of the albedo is interesting mainly for two reasons:  \\begin{enumerate} \\item From a geophysical point of view, it can serve to improve our understanding of the interaction of cosmic rays and solar wind with the atmosphere and the Earth's magnetic field.  \\item For gamma-ray-astronomical observations from satellites, the albedo constitutes an intense background. It is important to know the properties of this background in detail in order to be able to (a) avoid it as much as possible by specific pointing strategies and data filtering (either online or offline),  and (b) possibly use it as a  calibration source. \\end{enumerate}   \\noindent The EGRET high-energy gamma-ray telescope \\cite{thompson2} was operational as part of the Compton Gamma-Ray Observatory  satellite (CGRO) for 9.1 years from 1991 until 2000. CGRO had a circular orbit at approx. 450 km altitude (the altitude changed over time) with an inclination of 28.5$^\\circ$. From this perspective, the Earth's horizon appears at a zenith angle of $\\approx$111$^\\circ$. The instrument was almost always in pointed observation mode which meant that the Earth occulted the object of interest for on average nearly 50\\%  of the time, or in other words, the Earth was in the 80$^\\circ$ diameter field of view (FOV) for more than half the time.  In order to avoid the bright gamma-ray albedo of the Earth, the spark-chamber trigger of EGRET was automatically modified as the Earths horizon moved in and out of the FOV. Still, the number of albedo photons registered by EGRET is large. The total cleaned photon dataset comprises 5.2 million events. Roughly 60\\% of these were observed from zenith angles $> 105^\\circ$. More than half of the archival EGRET data are albedo photons!  In view of the upcoming GLAST mission \\cite{michelson}, the interest in the Earth's albedo increased, and it was decided to use the rich EGRET dataset to obtain the most detailed picture of the Earth's albedo available so far. The application of this knowledge to GLAST will be especially direct because GLAST is foreseen to fly an orbit very similar to that of CGRO.  The results presented here are from an ongoing analysis project and make use of only about 10\\% of the available data. In this first step, the analysis method was established and results of previous observations confirmed. The project will continue to study the complete data set, and look in addition at longitudinal and latitudinal variations of the measured parameters. ",
        "conclusions": ""
    },
    "0410/astro-ph0410164_arXiv.txt": {
        "abstract": "We examine a sample of low redshift ($10<d<150$ $h^{-1}$ Mpc) field galaxies including galaxies with luminosities as low as $M_r - 5 \\log_{10} h \\sim -12.5$, selected from the Sloan Digital Sky Survey Data Release 2 (SDSS). The sample is unique in containing galaxies of extremely low luminosities in a wide range of environments, selected with uniform and well-understood criteria.  We present the luminosity function as well as the broad-band properties of low luminosity galaxies in this sample. A Schechter function is an insufficient parameterization of the $r$-band luminosity function; there is an upturn in the slope for $M_r - 5 \\log_{10} h > -18$. The resulting slope at low luminosities in this sample is $\\alpha_2 \\sim -1.3$. However, we almost certainly miss a large number of galaxies at very low luminosities due to low surface brightness selection effects, and we estimate that the true low luminosity slope may be as steep as or steeper than $\\alpha_2 \\sim -1.5$. The results here are consistent with previous SDSS results and, in the $g$-band, roughly consistent with the results of the Two degree Field Galaxy Redshift Survey. Extremely low luminosity galaxies are predominantly low surface brightness, exponential disks, the majority of which are red. ",
        "introduction": "\\label{motivation}  Extremely low luminosity galaxies in the field are not a well-studied population, because they are difficult to find except in deep, wide-field surveys. On the other hand, one of the major problems facing models of galaxy formation is why cold dark matter models predict larger numbers of low mass galaxies than observed, either in the field or in the halos of larger galaxies. Thus, understanding the properties of the lowest luminosity galaxies (which are presumably also very low mass) might shed light on some of the mysteries surrounding the production of galaxies in low mass halos.  Most of the observational work regarding low luminosity galaxies has concentrated on galaxies in clusters. The Virgo cluster photometric sample of low luminosity galaxies (\\citealt{binggeli85b}), for which the background contamination is low, indicates that the low luminosity galaxies in clusters tend to be red and exponential in profile shape (e.g. \\citealt{barazza03a}).  The low luminosity slope of the luminosity function in the Virgo cluster appears to be $\\alpha \\sim -1.3$. A number of other purely photometric studies have found similarly steep slopes, usually as steep as or steeper than $\\alpha \\sim -1.4$ (\\citealt{depropris95a,trentham98a,secker97a}), and possibly a variation of the low luminosity slope depending on the cluster (\\citealt{driver98a}). However, it is worth noting that deriving luminosity functions from photometry alone requires background subtraction, an uncertain procedure which may bias one's estimates towards steep slopes (\\citealt{valotto01a}).  Recently, \\citet{deady02a} have studied the properties of spectroscopically confirmed dwarf galaxies in the Fornax cluster. They report that most of the low luminosity galaxies in that cluster are dE galaxies that lie on the optical color-magnitude and magnitude-surface brightness relationships for dwarf galaxies. Notably, they find a population of ultra compact dwarfs (UCDs) which tend to be very red and (as their name suggests) small even for their luminosity.  The other location in which astronomers have studied very low luminosity galaxies is of course in the Local Group (\\citealt{mateo98a}). It is in this environment that we have the best constraints on the luminosity function at very low luminosities (\\citealt{vandenbergh92a, trentham05a}). In the Local Group, ``late type'' dwarfs (spirals or irregulars) tend to be blue, though a large fraction ($\\sim 50\\%$) of the known dwarfs ($M_V > -14$) are ``early type'' (dwarf spheroidals and ellipticals) and have red colors typical of high luminosity galaxies.  In the past few years, there have been several new detections of low luminosity objects in the Local Group (\\citealt{ibata95a, armandroff98a, armandroff99a, whiting99a, zucker04a, willman04a}).  There are a number of efforts to collect data on low luminosity galaxies in other areas. \\citet{impey96a} have searched for low luminosity galaxies nearby by selecting low surface brightness galaxies.  More recently, \\citet{karachentsev04a} released a catalog of nearby galaxies selected from the literature, from eyeball inspection of photographic plates, from HI surveys, and from infrared surveys. Again, many of the galaxies were found by selecting low surface brightness objects which were followed up in the optical or radio. The resulting catalog contains galaxies as low luminosity as $M_B \\sim -8$. \\citet{gildepaz03a} have defined a set of ``blue compact dwarf'' (BCD) galaxies, having culled these from a number of different catalogs with varying selection criteria. A new very low luminosity galaxy ($M_I \\sim -10$) has been recently discovered in the field (\\citealt{pasquali04a}). There are also viable candidates (as yet unconfirmed) in small groups (\\citealt{flint01a}). The HI Parkes All-Sky Survey (HIPASS; \\citealt{meyer04a}) have a neutral-gas selected sample which includes many objects with low optical luminosities and surface brightnesses (West et al. in preparation).  Previous large field galaxy redshift surveys have provided samples of low luminosity galaxies selected without regard to environment: for example, the Center for Astrophysics redshift survey (\\citealt{huchra83a}), the Las Campanas Redshift Survey (\\citealt{shectman96a}), the IRAS Point Source Catalog redshift survey (\\citealt{saunders00a}), and the Two degree Field Galaxy Redshift Survey (2dFGRS; \\citealt{colless01a}). The Sloan Digital Sky Survey (SDSS; \\citealt{abazajian04a}) represents an improvement on these and other previous surveys. It has deep surface brightness limits in its photometry, covers a wide area of sky, and provides highly complete follow-up spectroscopy. On the other hand, its photometric and spectroscopic pipelines are not optimized for finding such low redshift objects, resulting in problems in creating a reliable low redshift catalog. In \\citet{blanton04a} we created such a catalog using extra processing and careful inspection of galaxies in the SDSS catalog. The result is not perfect; in particular, there is clearly incompleteness at the luminous, nearby end and at low surface brightness. However, the selection process, which is homogeneous, fairly well-defined, and unbiased with respect to environment, yields a catalog whose properties we can compare to theoretical predictions.  Our focus here is not on studying a necessarily {\\it complete} census of nearby galaxies --- impossible in any case given the current sky coverage of the SDSS --- but instead on describing the properties of a well-defined sample. For that reason, we have not included every object with a known (low) redshift in the catalog, but only those whose selection criteria we understand --- that is, those in the SDSS spectroscopic survey.  Our sample is publicly distributed as part of the NYU Value-Added Galaxy Catalog (\\citealt{blanton04a}). As described below, that catalog includes all of the parameters used to make the figures in this paper, including atlas images.  In this paper we present the catalog and a basic description of their properties.  Section \\ref{data} describes the SDSS, how we have selected our galaxies, and what properties we have measured. Section \\ref{surface} investigates the surface brightness completeness of our sample. Section \\ref{properties} describes qualitatively the distribution in space of these galaxies, the distribution of their basic photometric properties relative to those of more luminous galaxies, and their luminosity function.  Section \\ref{summary} summarizes our results.  Where necessary, we have assumed cosmological parameters $\\Omega_0 = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 100$ $h$ km s$^{-1}$ Mpc$^{-1}$ (with $h=1$). All magnitudes in this paper are $K$-corrected to rest-frame bandpasses using the method of \\citet{blanton03b}, unless otherwise specified.  Because of the small range of look-back times in our sample (a maximum of around 700 Myr), we do not evolution-correct any of our magnitudes. ",
        "conclusions": "\\label{summary}  We have presented a sample and basic properties of a set of low redshift ($z <0.05$) galaxies in the SDSS. Most interestingly, this sample is the only sample of galaxies extending down to $M_r \\sim -12.5$ which is unbiased with respect to environment.  From our measurements of the properties of these galaxies and of their luminosity function, we conclude: \\begin{enumerate} \\item The slope of the luminosity function at low luminosities ($M_r - 5\\log_{10} h > -17$), here denoted $\\alpha_2$, is at least as steep as $-1.3$, though it is almost certainly steeper, and perhaps as steep as $-1.5$. \\item Galaxies at these low luminosities are low in surface brightness, close to exponential, and predominantly blue. \\item Our estimate of the luminosity function is strongly affected by surface brightness selection effects at low luminosities. We have taken some care in estimating the completeness in this regime. \\item Incompleteness at low luminosities hardly affects our estimated total luminosity density because even with our completeness corrections the low lumosity galaxies do not contribute significantly. \\item There is not a large discrepancy between the SDSS and 2dFGRS determinations of the luminosity function. \\item Comparing void, field, and cluster luminosity functions in the SDSS yields the conclusions that the exponential cut-off at high luminosities is a strong function of environment and that $\\alpha_2$ is not a function of environment at low density, but may become significantly steeper at high density. \\end{enumerate}  There are a number of important improvements we can make to the measurements described here: \\begin{enumerate} \\item With later data releases of the SDSS, we can explore more of the local volume, of which we do not yet have a fair sample. In particular, we will be observing regions surrounding and including the Virgo cluster. This increased volume will allow a better measurement of the environmental dependence of the properties of low luminosity galaxies. \\item Deeper SDSS spectroscopy to $m_r < 19.5$, targeted at galaxies thought to be at $z<0.1$ according to photometric redshifts, may result in improved statistics on the abundance of low luminosity galaxies (Lin et al, in preparation). \\item Improved analysis of the SDSS imaging will allow us to detect low surface brightness galaxies. First, in the Southern Equatorial stripe of the SDSS, we have many epochs of imaging, which will allow us to detect galaxies by their diffuse light. Second, \\citet{willman01a} and \\citet{zucker04a} continue to search for very nearby galaxies detected in resolved stars. These samples will allow us to evaluate the effects of surface brightness completeness on the low luminosity slope empirically. \\item In order to improve the sample at the luminous end, we will supplement the SDSS redshifts with known large, bright galaxies from the astronomical literature, in a statistically meaningful way. \\end{enumerate} These last two items will allow us to probe the regime below 10 $h^{-1}$ Mpc, which we have excluded here because of the current incompleteness of the catalog in that regime. In addition to examining even lower luminosity galaxies, we will be able to perform a more direct comparison with the work of \\citet{karachentsev04a} at these distances.  With the data in hand, we will be studying further the environmental dependence of galaxy properties, analyzing the images to better understand the morphology of galaxies over a large range of luminosities, and seeking to perform follow-up observations to better understand the mass-to-light ratios, star-formation histories, and dust content of extremely low luminosity galaxies in the field."
    },
    "0410/astro-ph0410602_arXiv.txt": {
        "abstract": "We report on the abundance of interstellar neutral oxygen (OI) for 26 sightlines, using data from the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}), the {\\it International Spectroscopic Explorer} ({\\it IUE}), and the {\\it Hubble Space Telescope} ({\\it HST}).  OI column densities are derived by measuring the equivalent widths of several ultraviolet absorption lines, and subsequently fitting those to a curve of growth.  We consider both our general sample of 26 sightlines and a more restrictive sample of 10 sightlines that utilize {\\it HST} data for a measurement of the weak 1355 \\AA{} line of oxygen, and are thus better constrained due to our sampling of all three sections of the curve of growth.  The column densities of our {\\it HST} sample show ratios of O/H that agree with the current best solar value if dust is considered, with the possible exception of one sightline (HD 37903).  We note some very limited evidence in the {\\it HST} sample for trends of increasing depletion with respect to $R_V$ and $\\fHmol$, but the trends are not conclusive.  Unlike a recent result from \\citet{Cartledge2}, we do not see evidence for increasing depletion with respect to $\\nH$, but our {\\it HST} sample contains only two points more dense than the critical density determined in that paper.  The column densities of our more general sample show some scatter in O/H, but most agree with the solar value to within errors.  We discuss these results in the context of establishing the best method for determining interstellar abundances, the unresolved question of the best value for O/H in the interstellar medium (ISM), the O/H ratios observed in Galactic stars, and the depletion of gas-phase oxygen onto dust grains. ",
        "introduction": "\\label{s:intro} As the third most abundant element in the Galaxy after hydrogen and helium, oxygen is an important component in interstellar chemistry and dust grain models.  The exact abundance of gas-phase oxygen relative to hydrogen has been in question for quite some time.  The question of whether oxygen accumulates onto dust grains to varying degrees in different environments has also been investigated.  While recent research seems to be moving towards a picture of constant oxygen depletion from the gas, questions remain, especially with regard to the possible dependence of oxygen depletion on cloud density.  There is much work to be done in examining sightlines with more extreme physical properties---sightlines that have been previously ignored due to observational considerations.  Interstellar absorption lines in the ultraviolet were first detected with sounding rockets, e.g. \\citet{MortonSpitzer}.  Detailed studies of neutral oxygen in the interstellar medium became possible with the introduction of ultraviolet satellites such as {\\it Copernicus} and {\\it IUE} in the 1970s.  Studies such as those by \\citet{Morton74}, \\citet{Morton75}, \\citet{Zeippen}, \\citet{deBoer79}, and \\citet{deBoer81} each focused on the study of an individual sightline, finding that gas-phase oxygen was depleted relative to the solar value.  \\citet{MortonEtAl} also studied the column densities of several different atomic species, including neutral oxygen, for five sightlines.  The first extensive survey to determine the general interstellar value of O/H was reported by \\citet{York}.  Their study included 53 sightlines.  Of these, oxygen column densities were determined in 14 cases, while column density upper limits were determined in 16 cases and column density lower limits were determined in 20 cases (they did not report on OI for the final three sightlines).  Their analysis used the weak 1355 \\AA{} line of OI, relying on the assumption that NI and OI have a similar $b$-value.  They determined column densities by using this $b$-value to apply an appropriate saturation correction to the 1355 \\AA{} line when necessary.  In a few cases, the 988 \\AA{} triplet of OI was used, but further details were not given.  The conclusion of this study was that interstellar O/H is between 40\\% and 70\\% of the assumed solar value, from \\citet{Withbroe}, of approximately 690 oxygen atoms per million hydrogen atoms (later studies have assumed different solar values of the solar O/H ratio, and the correct value is still very much in dispute; see \\S\\ref{ss:dust} for a brief discussion).  Additionally, \\citet{York} concluded that there is no evidence of systematically enhanced depletion as reddening (i.e. total hydrogen column density) increases.  Rather, they found some limited evidence that sightlines with $\\logHtot<20.5$ are more depleted (with the mean O/H ratio approximately 40\\% solar) than sightlines with $\\logHtot>20.5$ (with the mean O/H ratio approximately 70\\% solar).  \\citet{Keenan} explored 26 sightlines---including one sightline from \\citet{York}---and found an average O/H ratio that is 50\\% of the solar value by \\citet{Withbroe}.  In contrast to the results of \\citet{York}, \\citet{Keenan} did not find a significant difference between sightlines with $\\logHtot<20.5$ and those with $\\logHtot>20.5$.  An important recent survey of interstellar oxygen was performed by \\citet{Meyer}.  Their study covered 13 sightlines and involved detailed analysis of the 1355 \\AA{} line of OI using the Goddard High Resolution Spectrograph (GHRS) on the {\\it Hubble Space Telescope}.  These 13 sightlines included four for which the column density had been previously determined, and three for which a saturation correction (estimated from MgII and NI) was applied.  The study found the mean O/H to be $319\\pm14 \\ppm$, with the largest deviation from the mean less than 18\\%.  Assuming as many as 180 oxygen atoms in the dust phase per million hydrogen atoms, a value consistent with many theoretical models (see \\S\\ref{ss:dust}), this corresponds to a total O/H ratio approximately two-thirds of the \\citet{Grevesse} solar value, $\\OHsolar=741\\pm130 \\ppm$.  Additionally, \\citet{Meyer} did not find any systematic, statistically significant variation of O/H with respect to the total column density of hydrogen ($\\Htot$), the average volume density of hydrogen along the sightline ($\\nH$), or the molecular fraction of hydrogen ($\\fHmol=2\\NHmol/[2\\NHmol+\\NHI]$).  We note that the \\citet{Meyer} study was confined to stars with $A_V \\lesssim 1.0$ mag.  Studies since \\citet{Meyer} may exhibit systematic differences from previous studies, in both total oxygen column density and O/H ratio relative to solar, for the following reasons: \\begin{itemize} \\item Revision of the $f$-value for the weak 1355 \\AA{} transistion (see \\S\\ref{ss:1355line}) from $1.25\\times10^{-6}$ to $1.16\\times10^{-6}$, as suggested by \\citet{Welty} \\item Revision of the solar O/H ratio from approximately $700 \\ppm$ or more \\citep{Withbroe, Grevesse} to approximately $550\\ppm$ or less \\citep{Allende, Holweger, Asplund} \\end{itemize} In comparing results with \\citet{Meyer}, more recent authors \\citep{Cartledge, Andre, Cartledge2} have made a linear correction to the reported O/H ratios of \\citet{Meyer} to reflect the change in $f$-value.  With the change, the mean \\citet{Meyer} value becomes $\\OH=343\\pm15\\ppm$.  Additionally, assuming $\\OHdust$ is as large as $180\\ppm$, the total O/H ratio can be reconciled with the revised solar value.  \\citet{Cartledge} examined 11 sightlines using the Space Telescope Imaging Spectrograph (STIS) onboard {\\it HST} to observe the 1355 \\AA{} line.  This study used two methods: profile fitting, where velocity information from other atomic species is incorporated into the fit of the oxygen line; and the apparent optical depth (AOD) method, where it is assumed that the line does not contain unresolved saturated structure and the measurement represents a lower limit on column density if the assumption fails (profile fitting can suffer from the same problem if unresolved, saturated structure is present and none of the velocity information from other species reveals such structure).  \\citet{Cartledge} analyzed a subset of seven sightlines, and compared them with the 13 sightlines measured with GHRS from \\citet{Meyer}.  They found that five of the seven sightlines measured with STIS lie below the O/H mean (adjusted for the new $f$-value of the 1355 \\AA{} line) determined by \\citet{Meyer}.  These five sightlines have five of the six highest hydrogen column densities in the combined sample of 20 sightlines.  This result is similar to O/Kr ratios reported in the same work, where some sightlines with higher total oxygen content exhibit oxygen depletion relative to krypton (a noble gas that is likely to be found almost exclusively in the gas phase).  Despite this, the oxygen column densities for only two sightlines have 1-$\\sigma$ uncertainties that exclude the possibility of consistency with the \\citet{Meyer} mean for O/H.  The evidence that there is oxygen depletion with respect to krypton, however, is somewhat more convincing.  \\citet{Moos} were among the first to use {\\it FUSE} to measure the equivalent widths of several different transistions for neutral oxygen, and to subsequently measure oxygen column densities by using curves of growth.  They did so for seven sightlines less than $180 {\\rm \\: pc}$ in length.  Despite large uncertainties and potential systematic errors from inhomogeneous determinations of atomic hydrogen, the seven sightlines were found to have O/H ratios from 234 to $479\\ppm$---consistent within a factor of 2.1.  Excluding the sightline BD+28$^\\circ$2411, the range is 363 to $479\\ppm$---consistent within a factor of 1.3.  Within the uncertainties, these results are consistent with \\citet{Meyer}.  The next major survey of interstellar oxygen was by \\citet{Andre}.  This survey measured the oxygen content of 19 sightlines using STIS.  Similar to \\citet{Cartledge}, \\citet{Andre} used both profile fitting and the AOD method to determine column densities.  Their analysis produced a mean $\\OH=408\\pm13\\ppm$.  The standard deviation in their sample was $59\\ppm$.  \\citet{Andre} included the previous results of \\citet{Cartledge} and \\citet{Meyer} to find a mean $\\OH=362\\ppm$, with a standard deviation of approximately 20\\%.  \\citet{Andre} took the average of values by \\citet{Allende} and \\citet{Holweger}, $517\\pm58\\ppm$, as the solar value of O/H.  Both the O/H mean of the 19 sightlines of the \\citet{Andre} sample and the O/H mean of the combined sample can be reconciled with this solar value if dust is considered.  The most recent interstellar oxygen survey was by \\citet{Cartledge2}, probing 36 sightlines using the same methods as \\citet{Cartledge}, and adding previous results for 20 sightlines, discussed results for a cumulative sample of 56 sightlines.  The results of \\citet{Cartledge2} reinforce the conclusions of \\citet{Cartledge} with many more data points.  Specifically, a 4-parameter Boltzmann function is fit to the ensemble of O/H ratios when plotted against $\\nH$, with an apparent transition centered at $\\nH\\approx1.5\\percc$.  This phenomenon is somewhat surprising in that it is not clear why $\\nH$ should accurately trace the environment of a sightline; rather it is more commonly expected that reddening parameters, e.g. $A_V$ or $E_{B-V}$, should be a better trace of the individual clouds that dominate a given sightline.  However, the trend seen by \\citet{Cartledge2} is not new; results by \\citet{SnowRach}, \\citet{Jenkinsetal}, \\citet{Spitzer}, and many other references cited therein have shown that the sightline characteristic with the most consistent correlation for enhanced depletion of many different elements is $\\nH$.  \\citet{Spitzer} discussed the possibility that $\\nH$ traces the superposition of three phases of the ISM: a warm, low-density phase, a cold, moderate-density phase, and a cold, high-density phase.  Regardless of the theoretical difficulties of determining a comprehensive picture for why $\\nH$ traces the interstellar environment so well, this most recent result of \\citet{Cartledge2} is the first time that this has been detected for oxygen.  The fit found by \\citet{Cartledge2} suggests that $\\OHdust=192\\pm51 \\ppm$ in the higher density sightlines, compared to $\\OHdust=86\\pm51 \\ppm$ in lower density sightlines.  As discussed in that paper, the increase of over $100\\ppm$ is difficult to account for with most dust models.  Of note, \\citet{Cartledge2} adopt a slightly different value for the solar oxygen abundance than \\citet{Andre}.  \\citet{Cartledge2} adopt $\\OHsolar=476\\pm50 \\ppm$, the weighted average of \\citet{Holweger} and \\citet{Asplund}, as opposed to the straight average of \\citet{Allende} and \\citet{Holweger} adopted by \\citet{Andre}.  Taking all these results into account, the emerging picture is that interstellar O/H is relatively constant over a wide range of environments.  If ${\\rm O_{dust}/H}\\lesssim180\\ppm$, then the interstellar O/H ratio can be reconciled with the solar value.  There are two exceptions to the observed constancy.  The first exception is the \\citet{Cartledge} study that hints at the evidence of, but does not conclusively prove, enhanced depletions in sightlines with higher total oxygen and hydrogen column densities.  The second exception is the \\citet{Cartledge2} study that shows a small but clear depletion effect as $\\nH$ becomes larger than $1.5\\percc$.  This is in strong contrast to the results for many other atomic species such as iron, silicon, and aluminum, species that show highly enhanced depletion in sightlines with higher reddening.  One assumption concerning these studies that should be noted is that the ionization energies of OI and HI are very similar, at $13.618 {\\rm \\; eV}$ and $13.598 {\\rm \\; eV}$, respectively.  This means that the fractional contribution to the total oxygen abundance from ionized oxygen (OII, OIII, etc.) will be mirrored by the fractional contribution to the total hydrogen abundance from HII.  A charge exchange reaction between oxygen and hydrogen also constrains this relationship \\citep{Field}.  Since the dissociation energy of molecular hydrogen is only 4.476 eV, the assumption that O/H is independent of ionization fraction holds across all molecular fractions as well.  In particular, the fact that a cloud has a high molecular fraction would tend to reveal that the ionization fraction is very small, since the majority of molecular hydrogen would be destroyed before significant ionization of atomic hydrogen would occur.  The goal of this study is to add to the current body of knowledge concerning interstellar oxygen by determining the O/H ratio for several sightlines observed by {\\it FUSE}.  We supplement the {\\it FUSE} data with {\\it HST} and {\\it IUE} data where possible.  Here we report on our measurements of the equivalent widths of several neutral oxygen absorption lines for each sightline, and our fits of these to curves of growth, from which we determined the $b$-values and column densities of oxygen.  The curve of growth method was used in the hope of avoiding assumptions about $b$-values and subsequent saturation corrections.  We then looked for systematic differences in O/H with respect to total hydrogen and various reddening parameters.  As a side goal, we aimed to determine the utility of {\\it FUSE} for observing oxygen absorption lines shortward of $\\sim1100$ \\AA{}, a wavelength region that is not covered by {\\it HST} or {\\it IUE} and has been observed only to a limited extent with {\\it Copernicus}, but where a number of weak-to-medium strength OI lines are found.  With some exceptions, e.g. \\citet{Zeippen}, studies of interstellar oxygen have been limited to analysis of the 1355 \\AA{} line of oxygen which, though weaker than the {\\it FUSE} lines, is sometimes strong enough to be saturated.  While the more recent of these studies have taken into account the profiles of other species to formulate a comprehensive picture of the cloud structure along a given sightline, there remains the possibility of saturation effects for individual velocity components.  It is not yet known whether studies such as this paper and \\citet{Moos} can equal or perhaps improve upon the results of previous studies by taking into account all the information that is now available to us through {\\it FUSE} in the form of many far-UV OI lines. ",
        "conclusions": "\\label{s:results} The value we adopt to summarize our O/H results depends on the sample we include.  If we take the results for the 10 sightlines for which {\\it HST} data were available, we find a mean O/H of $421^{+47}_{-33}\\ppm$, with a standard deviation of 29\\%.  As discussed further in \\S\\ref{ss:methods}, it is fair to consider these sightlines separately since the {\\it HST} data provide the best-constrained equivalent widths for the 1302 and 1355 \\AA{} lines, and the presence of these two lines gives the widest range in $\\log{f\\lambda}$, constraining all three parts of the curve of growth.  This mean value for the O/H ratio agrees very closely with \\citet{Andre} and the low-density sightline mean of \\citet{Cartledge2}.  If we include all 26 sightlines for which we determined O/H (including our adopted solutions for the 10 two-point curves of growth), our mean O/H is $468^{+103}_{-46}\\ppm$, with a standard deviation of 98\\%.  Given the conservative error bars, this value is still consistent with a solar abundance and also relatively consistent with, though larger than, the values measured by \\citet{Andre} and \\citet{Cartledge2}.  We plot the logarithmic ratio of gas-phase oxygen to hydrogen against total hydrogen column density and distance in Figures \\ref{fig:Htot} and \\ref{fig:distance}, plotting both the general sample and the {\\it HST} sample separately to show the difference in the samples.  The plots show the scatter in the O/H ratio for our less well-constrained lines of sight and no trend is apparent, except for perhaps with respect to distance to the illuminating star, i.e., the length of the sightline.  The most significant scatter in our O/H determinations occurs almost exclusively for sightlines greater than 1100 pc in length.  If, contrary to the reservation expressed in \\S\\ref{ss:outliers}, we accept our column densities for these sightlines, the reason for these deviations may be that over such distances we are sampling local inhomogeneities.  However, this would be in contrast to the results of \\citet{Andre}, who find consistent O/H ratios for many sightlines greater than 1000 pc in length.  Additionally, if we are observing local inhomogeneities, it is difficult to know where they occur with respect to distance since our measurements of equivalent width are integrated over the entire sightline.  However, with the previous discussion of \\S\\ref{ss:outliers} in mind, noting the large qualitative uncertainties and potential alternate solutions for many of these sightlines, the fact that the greater scatter in our results occurs for sightlines greater than 1000 pc in length is likely coincidental.  When we consider the sample of 10 {\\it HST} sightlines separately, any possible trend disappears.  The {\\it HST} sightlines are at most 1100 pc in length and do not show any increased scatter with respect to distance.  The {\\it HST} sample is plotted against $A_V$, $E_{B-V}$, $\\nH$, $R_V$, and $\\fHmol$ in Figures \\ref{fig:Htot}-\\ref{fig:molecularfraction}.  In the {\\it HST} sample we find the O/H ratios for two sightlines (HD 24534 and HD 37903) do not agree within the errors with the mean value found by \\citet{Andre} or the mean value found by \\citet{Cartledge2} for sightlines with a lower density than their derived critical density they of $\\nH=1.5\\percc$.  As discussed in \\S\\ref{ss:compare}, there is some discrepancy between our result and that of \\citet{Cartledge} for HD 37903.  However, whether our O/H ratio or that of \\citet{Cartledge} is adopted, it is the smallest O/H ratio of all 10 sightlines in the {\\it HST} sample.  Our derived O/H ratio for HD 154368 is similar to that of HD 24534, but with much larger error bars that include the \\citet{Andre} and low-density \\citet{Cartledge2} means.  There is the possible hint of two trends in the {\\it HST} sample.  HD 24534, HD37903, and HD 154368, with the lowest O/H ratios in the sample, also have the three largest values of $R_V$, the ratio of total to selective extinction, in the {\\it HST} sample.  These three sightlines, along with HD 27778, are the only four sightlines in the {\\it HST} sample with molecular hydrogen fractions greater than 0.5.  However, both of these trends are far from definite and we do not want to overstate the case based on only these three points in a sample of just 10 points.  Some implications of the possibility of a trend with $R_V$ are discussed in \\S\\ref{ss:dust}.  In addition to the lack of a trend with $A_V$ and $E_{B-V}$, we do not see a trend with respect to $\\nH$, the trend that has been seen by \\citet{Cartledge2}.  However, comparing our sample size of 10 sightlines to the 56 sightlines of \\citet{Cartledge2} indicates the absence of evidence is not necessarily the evidence of absence.  In our {\\it HST} sample, only HD 27778 and HD 210839 are more dense than the critical density determined by \\citet{Cartledge2}.  Both of these sightlines have O/H ratios larger than the \\citet{Andre} and low-density \\citet{Cartledge2} means, though consistent within the errors.  HD 24534, HD 37903, and HD 154368 all have $\\nH$ between 1.1 and $1.4\\percc$.  \\subsection{Evaluation of Methods} \\label{ss:methods} The difference between our {\\it HST} and general samples sheds some light on the methods that we should use.  The {\\it HST} sample uses data on the 1355 \\AA{} line.  Even if the line is somewhat saturated, it is weak enough that it produces a moderately strong constraint on the column density.  When combined with results from other lines, such as the 1039 \\AA{} line, we can determine a $b$-value rather than merely assuming one.  Hindering this, however, is the fact that the velocity structure may be very complex, requiring multi-component curves of growth for the most accurate picture (see \\S\\ref{ss:bvalues}).  Our conclusion regarding the utility of the {\\it FUSE} data for determining neutral oxygen column densities is mixed.  We note the case of HD 24534 discussed in \\S\\ref{ss:compare} where measurements of absorption lines in the {\\it FUSE} wavelength region led to a significant revision in $b$-value and resulting column density compared to a previous study \\citep{SnowXPer}, when used in conjunction with {\\it HST} measurements of the 1355 \\AA{} line, a result confirmed by another study \\citep{Knauth}.  This result indicates the possible importance of using {\\it FUSE} data in oxygen column density determinations.  However, {\\it FUSE} lines without support from {\\it HST} data can often lead to poorly constrained column densities.  Even though {\\it IUE} can measure the two lines at 1302 \\AA{} and 1355 \\AA{}, uncertainties in the {\\it IUE} spectra complicate matters, and the 1355 \\AA{} line is typically not observed.  Our evaluation of the {\\it FUSE} data in this context is that the far-UV lines should be viewed as a very useful tool in conjunction with good data on the 1355 \\AA{} line, but the far-UV data alone have some limitations if such additional data are not available.  Compared to recent studies that use profile fitting, the curve of growth method is inferior when using the lines found in {\\it FUSE} without {\\it HST} measurements of the 1302 \\AA{} and 1355 \\AA{} lines.  However, when {\\it HST} measurements of these lines are available, a curve of growth with the addition of the medium-strength {\\it FUSE} lines seems to provide results similar to profile fitting in most cases.  In \\S\\ref{ss:compare} we discussed the fact that there are two cases of discrepancy where the evidence seems to favor previous profile fitting results (HD 37903 and HD 207198) based on consistency with the line profiles.  In both of these cases the reason for the discrepancy seems to arise from the fact that the sightline has a complex velocity structure that cannot be easily approximated by a single cloud.  However, in one case where a single cloud is a good approximation (HD 27778), the curve of growth indicates that the profile fitting methods were unable to detect some saturation, a suggestion that is reinforced by the line profiles of the stronger lines.  Our biggest points of concern with respect to systematic errors are the inhomogeneous determinations of hydrogen (particularly for sightlines where only the reddening relationship was used to determine total hydrogen) and the fact that both sightlines with an {\\it IUE} measurement of the 1355 \\AA{} line (HD 42087 and HD 152236) had anomalously high O/H ratios.  Future work will undoubtedly allow us to resolve the first issue and refine our results.  With respect to the second point, as discussed in \\S\\ref{sss:HD42087} and \\S\\ref{sss:HD152236}, there is a great deal of uncertainty with both of the 1355 \\AA{} line measurements for these sightlines, including possible contamination from stellar lines.  It would be presumptuous to assume a systematic error in the {\\it IUE} data based on only two data points that are admittedly not well-determined.  Furthermore, if our derived column densities are accurate, this would likely be the result of a selection effect:  the {\\it IUE} data has the poorest S/N and resolution of all the data used, and we would only be likely to conclusively observe the 1355 \\AA{} line in a sightline with a very large oxygen column density.  \\subsection{$b$-values} \\label{ss:bvalues} The mean $b$-value of the {\\it HST} sample is $7.7^{+0.4}_{-0.3}\\kmpers$ with a standard deviation of 49\\%.  If we include all 26 sightlines, we find a mean $b$-value of $9.4^{+0.5}_{-0.3}\\kmpers$ with a standard deviation of 47\\%.  These $b$-values are somewhat larger than the typical $b$-values ($\\sim {\\rm few}\\kmpers$) assumed for oxygen in many previous studies, e.g. \\citet{York}.  It may also be an indication that the $b$-value for neutral oxygen is inconsistent with the $b$-values of other atomic species, which may provide insight into variability of composition in clouds along a given line of sight.  For example, these average $b$-values are on the upper end of the NI $b$-values found by \\citet{York}.  Since NI has an ionization potential that is slightly smaller than oxygen, it is possible that cloud components along the line of sight subject to ionizing influences (e.g., shocks and radiation) contain a smaller than average ratio of neutral nitrogen to neutral oxygen, while at the same time being the cloud components with the largest velocity dispersions due to the energy imparted by these ionizing influences.  This would allow for the possibility that the $b$-value of oxygen is larger than the $b$-value of nitrogen.  If this scenario of differing $b$-values for different atomic species is found to be common, it may provide some insight into the dynamics and composition of various cloud components along a given sightline.  For the sightlines that have data from {\\it HST}, we can attempt to directly measure the ``effective'' $b$-value of the 1302 \\AA{} and 1355 \\AA{} lines, albeit through different methods.  The $b$-value of the 1355 \\AA{} line can be inferred through the Gaussian width of the fit.  The $b$-value of the 1302 \\AA{} line is derived through the fit of a Voigt profile.  In the context of this paper, a Voigt profile is any saturated profile, one that may include damping wings.  In all cases of the 1302 \\AA{} line with HST data, the damping wings were apparent.  When strong damping wings are present, the $b$-value and column can be determined simultaneously with reasonable accuracy (in our measurements, the 1302 \\AA{} line was never strong enough for the profile to be completely independent of $b$-value).  The 1039 \\AA{} line, which was usually well-fit by a Gaussian but occasionally required the fitting of a Voigt profile, can also be fit with the same Voigt profile program in all cases to determine $b$-value and column density (in cases of minimal saturation, the profile reduces to a Gaussian).  However, because the 1039 \\AA{} line is usually on the flat portion of the curve of growth, there can be a fairly wide range of solutions for the $b$-value and column density that produce reasonable fits.  Yet if one parameter ($b$-value or column density) is fixed (for consistency with the other profiles), the other parameter can be determined.  Fitting Voigt profiles to lines below 1000 \\AA{} (discussed in \\S\\ref{ss:1000lines}), where the spectrum had many other interfering lines and a much more uncertain continuum, was used in a few cases to determine equivalent widths, but is less useful for direct determinations of column density and $b$-value.  Through an iterative process of fitting one of the three major lines (1039 \\AA{}, 1302 \\AA{}, and 1355 \\AA{}) for both column density and $b$-value, and then generating simulated absorption profiles for the other lines for consistency, we examined the nature of the nine sightlines where we had high-resolution {\\it STIS} data.  In this process, it was initially assumed that the absorption lines were well-represented by a single-component cloud along the line of sight with a specific $b$-value.  Using a combination of results from this process and results from the literature where applicable, we also engaged in multiple-component fitting and profile generation.  Using this method, the sightlines towards HD 206267 and HD 207538 are the simplest cases to check for consistency.  The column density and $b$-value derived through the curve of growth method generated reasonable fits within the errors to the observed profiles of the three major lines.  The evidence that these sightlines are not well-represented by a single-cloud approximation is very limited from the oxygen lines.  These two sightlines are the only sightlines in our {\\it HST} sample that have not previously been included in oxygen studies (see \\S\\ref{ss:compare}), but optical spectra of these sightlines have been published in \\citet{Pan}.  While the optical spectra of elements such as potassium do reveal multiple cloud components within $\\sim10\\kmpers$ of each other for these sightlines, this is consistent with the $b$-values of $6.3\\kmpers$ (HD 206267) and $5.4\\kmpers$ (HD 207538) that were determined with the curve of growth method and verified through profile generation.  Another relatively simple case is the sightline towards HD 27778.  Using profile fitting methods, \\citet{Cartledge} find two major cloud components in this sightline.  However, the weak 1355 \\AA{} line does not show independent evidence of this, and can be fit well by a single Gaussian.  A consistent picture of column density and $b$-value can be achieved with the single-component approximation.  The sightline towards HD 210839 was reported on by \\citet{Andre}.  Observations of the 1355 \\AA{} line make it apparent that there is a complex velocity structure along this line of sight.  We used the velocity measurements of \\citet{Andre} to generate profiles for the three major lines that fit the data well.  The column density used to generate these profiles was the same as the column density independently measured by \\citet{Andre} and through our single-component curve of growth analysis.  The curve of growth analysis measured a $b$-value of approximately $9\\kmpers$, which also generates relatively accurate profiles in the single-component approximation.  This is anecdotal evidence, along with our earlier analysis of HD 206267 and HD 207538, that sometimes even complex velocity structure can be well-approximated by a single $b$-value.  Other sightlines showed an interesting property: consistent single-component profiles could be generated for the various absorption lines with a consistent column density, but with the $b$-value increasing with the strength of the line.  This is physically possible if the lines are not best represented by a single-component cloud, but rather there are one or more small, undetected components hidden in the noise near the weak 1355 \\AA{} line.  These small components would not significantly contribute to the overall column density, but would contribute to the profile of the stronger lines.  While this may be physically reasonable, simple attempts to model this phenomenon through profile generation were unsuccessful.  For the sightline towards HD 24534, appropriate profiles for all three major lines can be generated with a consistent column density of $\\logOI=17.83$, but with increasing $b$-values under the single-component assumption: $2.0\\kmpers$ for the 1355 \\AA{} line, $6.5\\kmpers$ for the 1039 \\AA{} line, and $8.2\\kmpers$ for the 1302 \\AA{} line.  This suggests that that the line of sight is not best represented by a single cloud component, but that our column density is accurate.  The curve of growth method suggests this same column density, and a $b$-value closest to the value measured for the 1039 \\AA{} line.  The curve of growth analysis caused us to reject an old literature value for the oxygen column density towards HD 24534 (see \\S\\ref{ss:compare}), and the profile fitting of the stronger lines confirms this.  The sightlines towards HD 185418, HD 192639, and HD 207198 present similar problems.  In all three cases, multiple components are obvious in the weak line 1355 \\AA{} or have been assumed elsewhere in the literature.  Yet generating profiles for the 1039 \\AA{} and 1302 \\AA{} lines in a given sightline using a column density consistent with the weak line and the measured or assumed multiple components of the weak line does not produce good results, in that the generated profiles of the stronger lines are either weaker or narrower than observed in the data.  Relatively consistent results can be obtained through the single-component approximation for a consistent column density but increasing $b$-values with the strength of the absorption line.  Once again, this suggests that there are one or more hidden weak cloud components.  As discussed in \\S\\ref{ss:compare}, the sightline towards HD 37903 is complicated because of absorption lines from vibrationally excited $\\Hmol$ in the spectrum.  However, the strongest $\\Hmol$ lines that might influence the measurement of the oxygen lines have rotational excitation levels of $J=7$ or greater, and it is unlikely that the potential contamination from $\\Hmol$ is solely responsible for the observed increase of $b$-value in the stronger lines with the single-component approximation.  In keeping with \\citet{Cartledge}, we measure the the apparent $b$-value of the 1355 \\AA{} line to be $1.6\\kmpers$; however, the $b$-value of the 1302 \\AA{} line is greater than $9\\kmpers$ when approximated as a single component.  We may once again infer there are hidden weak components near the narrow 1355 \\AA{} line.  As discussed in \\S\\ref{ss:compare}, a moderate range of column densities, when paired with the correct $b$-value, can produce a reasonable approximation to the 1302 \\AA{} line.  If we assume that this model of very weak undetected components is correct, determining the best $b$-value and column density for this type of distribution along the line of sight would best be served by the use of a multi-component curve of growth.  While the curves of growth we have used contain two separate damping wing regions to accommodate the differing damping constants of the 1039 \\AA{} and 1302 \\AA{} lines, a multi-component curve of growth in this context refers to a curve of growth with the assumption of multiple cloud components, possibly each with its own $b$-value.  This type of curve of growth possesses a ``multi-tiered'' flat region.  For lack of constraining data points, we have not used multi-component curves of growth, though some of our sightlines are suggestive of the need for them.  If indeed many or most sightlines are better represented by a multi-component curve of growth, it partially mitigates the discrepancy between our study and previous studies discussed at the beginning of this section, as previous studies have measured smaller $b$-values for the weak 1355 \\AA{} line than our measurements through the curve of growth method.  Multi-component curves of growth might be a good compromise between the two methods.  Whereas previous studies such as \\citet{York} applied a saturation correction based on $b$-value and single-component curves of growth might have a higher $b$-value (corresponding to an approximate ``effective'' $b$-value for multiple components) and not apply the same correction, multi-component curves of growth would most effectively consider the constraints from all the observed absorption lines.  The difficulty would be in obtaining a quality constraint on each line.  Future work with higher S/N spectra and the well-constrained detection of lines below 1000 \\AA{} will hopefully allow us to explore this important issue in greater depth.  \\subsection{Dust Grains} \\label{ss:dust} We have quoted a theoretical upper limit from \\citet{Cardelli} on the amount of oxygen that is tied up in interstellar dust grains, ${\\rm O_{dust}/H}\\lesssim180\\ppm$.  This limit comes from an assumption on the composition of grains, and the limited amount of other elements (Fe, Si, Mg, etc.) that are available to combine with oxygen in the form of silicates in these grains.  Using the mean of our {\\it HST} sample, we find ${\\rm O_{dust}/H}\\approx79\\ppm$, with the assumption that the total (gas + dust) oxygen content of the ISM is equal to the solar value and with the upper limit determined by the lower gas-phase O/H ratio of our best determined sample.  We do not estimate errors on this value since it is highly dependent on the solar value assumed.  Given the conservative errors on each sightline and the cumulative range in the solar O/H ratio spanned by the results of \\citet{Allende}, \\citet{Holweger}, and \\citet{Asplund}, this upper limit on dust-phase oxygen abundance may be uncertain by a factor of two or more.  \\citet{Andre} use similar reasoning (but a slightly lower ISM O/H ratio and a slightly higher solar O/H ratio) to determine ${\\rm O_{dust}/H}\\approx109\\ppm$.  If it is first assumed that maximum amount of oxygen available for dust (${\\rm O_{dust}/H}\\lesssim180\\ppm$) is in fact in dust and the current measurement of dust is the percentage resilient to shocks, then the \\citet{Andre} result implies that 60\\% of dust grains containing silicates is resilient to shocks.  The results of our {\\it HST} sample are in approximate agreement with this (we find 44\\% resiliency).  An alternate interpretation to this presumed resiliency is the possibility that there are more oxides than silicates in grains, allowing for the enhanced depletion of elements such as iron, magnesium, and silicon, without requiring as much depletion of oxygen.  We noted that our results hint at a trend of increasing depletion with increasing $R_V$.  Papers by \\citet{CCM1, CCM2} have shown that most features of the UV extinction curve can be determined to good approximation if the single parameter of $R_V$ is known.  Other studies such as \\citet{MathisWallen} and \\citet{MathisWhiffen} have shown that larger values of $R_V$ require grain size distributions weighted towards larger particles.  Thus, a trend of enhanced depletion with $R_V$ would imply that the depletion of oxygen is tied to grain size.  In the coldest, most dense environments large grains can accrete materials directly from the gas phase to mantles surrounding their cores, and the previously mentioned theoretical limits on oxygen in dust (based on the assumption that all oxygen in dust is found in silicates and oxides and limited by abundances of other elements) do not necessarily apply.  This would be a means for the enhanced depletion of oxygen.  In fact, if these skewed size distributions are the result of accretion, abundant elements such as carbon, nitrogen, or oxygen would be the required source of accretion material as refractory elements such as iron and calcium can only account for a limited increase in total grain volume \\citep{CCM2}.  However, grain coagulation from collisions can also account for skewed size distributions and in some cases is a better fit to the total extinction per hydrogen nucleus \\citep{CCM2}.  The problem with any interpretation of our result for the amount of oxygen in dust is our underlying assumption that the solar abundance of oxygen is the correct value with which to compare the total oxygen abundance in the ISM, and therefore the value from which we derive the amount of oxygen in dust.  As discussed in \\S\\ref{s:intro}, the accepted solar ratio of O/H has undergone significant revision in recent years, raising questions of how well this value is known.  To this point, the most recent values \\citep{Allende, Holweger, Asplund} together suggest that, within the 1-$\\sigma$ errors, the true solar O/H ratio is only known to a precision of $\\sim30\\%$ the average value adopted in this paper.  Published values of solar O/H before \\citet{Withbroe} show an even larger spread \\citep{Sofia}.  \\subsection{O/H of Galactic Stars} \\label{ss:stellarOH} Even if the solar ratio of O/H was well known, would it be appropriate to use it as the standard value for the total ISM oxygen abundance?  This question has been raised concerning krypton (which should not be depleted onto grains) by \\citet{CardelliMeyer}, who find the average Kr/H ratio in the ISM to be $\\sim60\\%$ of the solar ratio.  Subsequent studies \\citep{Meyer, Andre} have reviewed potential reasons for a deficit in the local ISM of elements such as krypton and oxygen:  enrichment of the presolar nebula, migration of the Sun from a smaller Galacto-centric distance, and infall of metal-poor gas onto the ISM.  Another possibility is metallicity enhancement resulting from the late infall of exoplanets or other metal-rich material onto the parent star (in this case the Sun).  If convection has begun in the parent star, the metals would remain in the photosphere and the observed metallicity would be enhanced relative to the actual metallicity.  Our study does not place strict enough constraints to comment on the relative probability of these scenarios.  However, the fact remains that the solar value of O/H is still in some question, and even if this were not the case, our estimation of oxygen in dust is very uncertain given that we cannot assume the solar value of O/H is equivalent to the average value in the ISM.  \\citet{SnowWitt} and \\citet{SofiaMeyer} both performed surveys of B, F, and G stars, finding that B stars have O/H ratios of $\\sim350\\ppm$, while F and G stars have O/H ratios of $\\sim450\\ppm$.  Both of these values are smaller than the current best values for the solar O/H ratio.  Additionally the O/H ratios for B stars, expected to best represent the ISM due to their recent formation, tend to be significantly lower than the average measured gas-phase O/H ratio of the ISM.  Similarly, the O/H ratios for F and G stars are not much larger than the average measured gas-phase O/H ratio of the ISM.  Both samples show spreads that are larger than the uncertainties, a result that is in contrast with studies such as \\citet{Meyer} and \\citet{Andre}, which show a very consistent O/H ratio across many sightlines.  The question must be asked as to how there can be such a large variation in the O/H ratios in stars that formed from a well-mixed ISM, where O/H is constant.  The answer likely lies in aspects of stellar formation that are not yet fully understood.  For example, \\citet{SnowJGR} suggests that ambipolar diffusion in star-forming clouds might be able to carry away some ions and thus deplete metals in the material that stars form from.  Furthermore, photoablation of the disks of the most massive stars may prevent the type of late metallicity ``pollution'' discussed above, thus explaining why B stars have significantly lower O/H ratios than F and G stars.  A better understanding of these issues will give further insight into whether or not the elemental abundances of the Sun or any subset of observed stars can be used as a good proxy for the abundances of the ISM.  \\subsection{Future Work} \\label{ss:Future} There is much work to be done concerning oxygen in the interstellar medium.  While this study is consistent with the picture that there is a constant depletion (relative to the solar value) of gas-phase oxygen abundance relative to hydrogen and that enhanced depletions in dense regions have not been conclusively observed, we have not increased the overall constraints already established by studies such as \\citet{Meyer}, \\citet{Cartledge, Cartledge2}, and \\citet{Andre}.  Perhaps with continued investigation into sightlines of very high reddening, we will conclusively observe our suggested trend with $R_V$.  However, this will be very difficult to do with current instrumentation, as the highly-reddened sightlines desired produce spectra with the poorest S/N.  Further work with {\\it FUSE} and the possibility of a second {\\it FUSE} satellite, with wavelength coverage to observe the important 1039 \\AA{} line and other lines blueward of 1000 \\AA{}, would help to constrain the curve of growth in many cases.  Finally, future work by infrared space telescopes SIRTF and the {\\it James Webb Space Telescope} ({\\it JWST}) should help to constrain the composition and amount of interstellar dust through solid-state IR features, perhaps allowing an appropriate standard of the total (gas + dust) O/H ratio to be determined from a direct measurement of dust-phase oxygen, rather than inferring dust-phase oxygen from an assumed total O/H ratio.  Unfortunately, the futures of UV instruments that are or were to be installed on {\\it HST}, specifically STIS and the Cosmic Origins Spectrograph (COS), are currently unknown.  There is also work to be done concerning other atomic species.  If oxygen does not become significantly depleted onto grains in dense environments, what elements do make up the grains?  Additionally, is the total amount of material in dust grains consistent with current models?  Is there conclusive evidence of a local elemental deficit?  Why are the krypton and sulfur abundances in the local ISM significantly subsolar?  Studies such as those by \\citet{Cartledge, Cartledge2} and \\citet{Andre} have reconciled the ISM value of O/H with the solar value if dust is included, but the opposite question remains:  is there enough oxygen to form the large number of grains that are required by the depletion of other elements?   In summary, we measure equivalent widths for neutral oxygen absorption lines along 30 sightlines, and determine column densities for 26 of these sightlines.  Our best-constrained column densities agree very closely with the interstellar O/H ratios derived by \\citet{Andre} and \\citet{Cartledge2}.  Our less well-constrained column densities show some scatter in O/H, but many agree (to within errors) with the solar value we adopt of O/H=$500\\pm75\\ppm$.  We find an O/H mean of $421^{+47}_{-33}\\ppm$, with a standard deviation of 29\\%, for our 10 best-determined points, and an O/H mean of $468^{+103}_{-46}\\ppm$, with a standard deviation of 98\\%, for all 26 sightlines included in our final analysis.  The two cases of an apparent supersolar abundance have questions, as discussed in \\S\\ref{ss:outliers}.  There are two potential cases of significantly enhanced depletion in our sample, the sightlines toward HD 37903 and HD 199579.  Our results for neither sightline are conclusive, as the O/H ratio of HD 37903 may still be consistent with the more moderately enhanced depletion found by \\citet{Cartledge2}, and profile fitting of the 1039 \\AA{} line in the HD 199579 spectrum suggests that a smaller $b$-value and larger column density is a better solution.  We argue for a revision of the neutral oxygen column density in the sightline towards HD 24534, studied previously by \\citet{Snow154368}.  Our revision is consistent with a recent result by \\citet{Knauth}.  We also suggest that the literature value for the neutral oxygen column density in the sightline towards HD 27778 \\citep{Cartledge} may need revision, a suggestion supported by profile fitting of the stronger oxygen lines.  A curve of growth analysis will be accomplished best in the near future by the simultaneous use of {\\it FUSE} and {\\it HST} data, as {\\it FUSE} data alone cannot consistently provide a strong constraint on the column density of oxygen.  When {\\it HST} data are available, we argue that our curve of growth method produces results similar to profile fitting methods in most cases and can be considered an appropriate parallel technique for determining elemental column densities.  Our general sample does not show convincing evidence of enhanced depletion of oxygen that is systematic with respect to any reddening parameters along the line of sight (e.g., $A_V$, $n_H$, etc.).  However, the enhanced depletion effect with respect to $\\nH$ found by \\citet{Cartledge2} is small and would not be seen in our general sample.  Our better-constrained {\\it HST} sample does not show the effect found by \\citet{Cartledge2}, though this is based on a minimal number of points.  In the {\\it HST} sample, we find some very limited evidence that depletion increases with respect to $R_V$ and $\\fHmol$, but the results are far from conclusive.  Further confirmation of these possible trends and the trend with respect to $\\nH$ will be important as we move closer to a definitive answer for the overall abundance of interstellar oxygen."
    },
    "0410/astro-ph0410434_arXiv.txt": {
        "abstract": "{ We present {\\it XMM-Newton} data centered on one of the MUNICS Near Infrared Cluster Survey fields (S2F1) and we discuss the X-ray properties of the 6 X-ray emitting EROs found. For one of them we have already obtained the redshift using near--infrared spectroscopic data, while for the remaining 5 EROs the analysis is based on photometric redshifts. We find evidence for the presence of an X--ray obscured QSO in at least 5 out of the 6 X-ray emitting EROs. For these 5 objects we derive intrinsic (2--10 keV) luminosities in excess of 10$^{44}$ erg s$^{-1}$ and intrinsic column densities higher than 10$^{22}$ cm$^{-2}$. These values have been obtained through a basic X-ray spectral analysis for the three brightest sources and through the analysis of the hardness ratios  for the remaining two. All of these 5 X--ray emitting EROs appear extended in the optical/near--infrared bands indicating that the host galaxy emission dominates at these wavelengths. This suggests that the hosted AGNs are likely to be absorbed also in the optical/near--infrared bands: i.e. they are likely X--ray obscured possible type~2 QSOs. For the remaining ERO the presence of an AGN is suggested both by its high 0.5-2 keV luminosity (L$_{0.5-2\\rm keV}\\sim10^{43}$ erg s$^{-1}$) and by its X-ray-to-optical flux ratio. In this case the quality of the present data prevents us from placing firm constraints on the AGN type hosted. Finally, the near-IR spectrum obtained for one of the 6 EROs classifies the host galaxy as an elliptical at $z\\simeq1.7$ with a stellar mass well in excess of 10$^{11}$ M$_\\odot$. This result corroborates the possible link between the QSO activity and the formation of massive spheroids. ",
        "introduction": "As shown by {\\it Chandra} and {\\it XMM-Newton} observations, a sizable fraction of the 2--10 keV selected sources are associated with Extremely Red  Objects (EROs, R-K$>$5). In particular, about 15\\% of the 2--10 keV sources down to a flux limit  of 10$^{-15}$ erg cm$^{-2}$ s$^{-1}$ are EROs (e.g. Rosati et al. \\cite{Rosati}; Mainieri et al. \\cite{Mainieri}). This already significant fraction should be considered as a lower limit since several X--ray sources are not yet covered by deep K--band observations. Unfortunately, good--quality optical and/or X--ray spectra are available only for a few X--ray emitting EROs,  preventing us from firmly establishing their physical nature. The X--ray and optical/near--IR photometric properties, the only information available for the large majority of the X--ray emitting EROs studied so far, strongly suggest that the bulk of this population is composed of obscured AGNs at least at the brightest (F$_X>$10$^{-15}$ erg cm$^{-2}$ s$^{-1}$) X--ray fluxes (e.g Alexander et al. \\cite{Alexander02}, \\cite{Alexander03}; Barger et al. \\cite{Barger02}, \\cite{Barger03}; Stevens et al. \\cite{Stevens}; Szokoly et al. \\cite{Szokoly}, Brusa et al. \\cite{Brusa04}). Although dusty starbursts are also observed as EROs (e.g. Brusa et al. \\cite{Brusa02}), their fraction is not high enough to dominate the population of X--ray emitting EROs at these X--ray fluxes. Indeed, the large majority of X--ray emitting EROs have measured X--ray--to--optical flux ratios  typical of AGN (i.e. F$_X$/F$_{opt}>$0.1) and evidence of obscuration comes from both their hard X--ray colors and their extremely  red optical--to--near--infrared colors. In particular,  hard X--ray colors are generally associated with a large amount of circumnuclear gas (N$_H$\\footnote{As usually done in the recent literature, we use here N$_H$=10$^{22}$ cm$^{-2}$ as the dividing N$_H$ value between X--ray unabsorbed and absorbed sources.}$>$10$^{22}$ cm$^{-2}$) which blocks the softer X--ray emission, while red optical colors indicate that large amounts of dust are absorbing the typical intrinsic blue continuum of the AGN (Elvis et al. \\cite{Elvis}, Risaliti \\& Elvis \\cite{Risaliti}). The few cases for which  either X--ray and/or optical spectra are available seem to confirm the presence of X--ray and/or optical obscuration (e.g. Lehmann et al. \\cite{Lehmann}; Vignali et al. \\cite{Vignali}; Mainieri et al. \\cite{Mainieri}; Willott et al. \\cite{Willott}; Roche et al. \\cite{Roche}; Stevens et al. \\cite{Stevens}; Mignoli et al. \\cite{Mignoli}; Szokoly et al. \\cite{Szokoly}; Brusa et al. \\cite{Brusa}, \\cite{Brusa04}) and, more important, indicate that most of the X--ray emitting EROs host high--luminosity AGNs (L$_X>$10$^{44}$ erg s$^{-1}$, i.e. QSOs). These findings open the possibility that a significant fraction of X--ray emitting EROs could host X--ray obscured and optically absorbed (i.e type~2) QSO (for details on this class of sources see e.g. Norman et al. \\cite{Norman}; Stern et al. \\cite{Stern}; Della Ceca et al. \\cite{Della Ceca03} and references therein). These sources are predicted in large numbers by the synthesis model of the Cosmic X--ray Background (38\\% and 16\\% of Cosmic X--ray Background should be due to QSO2 following the Gilli et al. \\cite{Gilli} and Ueda et al. \\cite{Ueda} models, respectively).  To study the nature of X--ray emitting EROs, we have obtained  two {\\it XMM-Newton} observations respectively on the  S2F1 (AO--2 period) and on the  S2F5 (AO--3 period) fields of the {\\it MUnich Near-IR Cluster Survey} (MUNICS, Drory et al. \\cite{Drory}, Feulner et al. \\cite{Feulner}). These fields ($\\sim160$ arcmin$^2$ each) are covered by photometric observations in the B, V,  R, I, J and K' bands (Drory et al. \\cite{Drory}) down to limiting magnitudes of R$\\simeq$24 mag and K'$\\simeq$19.3 mag (50\\% completeness limits for point sources). Thirty-six out of the 170 EROs present in these two fields have a K' magnitude brighter than 18.5 and are already under investigation by an ongoing spectroscopic survey performed by our group (see  Saracco et al. \\cite{Saracco1},  \\cite{Saracco2}).  Here we discuss the X-ray properties of 6 X-ray emitting EROs found in the S2F1 field. The {\\it XMM-Newton} data centered on the S2F5 field have not yet been made available by the {\\it European Space Agency}. A comprehensive analysis of the whole sample of near-IR selected EROs including the stacked analysis and the statistical X-ray properties performed on both fields will be presented in a forthcoming paper.  Throughout this paper we assume H$_0$=70 km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_M$=0.3, $\\Omega_{\\Lambda}$=0.7. All the magnitudes are in the Vega system. ",
        "conclusions": "In this paper, new {\\it XMM-Newton} data have been used to study and to discuss the nature of 6 X--ray emitting EROs found in the MUNICS S2F1 field. For these 6 sources optical and near--infrared photometry as well as photometric redshifts are available from the MUNICS catalog. For one of the 6 objects, ERO S2F1\\_443, we have presented and analyzed the near--infrared low--resolution spectrum, deriving  the spectroscopic redshift and unveiling the nature of the host galaxy.  Five (71, 443, 507, 551, 714) out of the 6 EROs discussed here have X--ray properties matching those expected for X--ray obscured QSO. For 3 of them (71, 507, 551) these properties have been derived by performing a basic X--ray spectral analysis, while for the other 2 (443 and 714) we used their X--ray broad--band photometry (i.e. X--ray emission, hardness ratio analysis, radio properties). Due to the extended appearance of these 5 sources in the optical and near--infrared images, i.e. they are galaxy dominated at these wavelengths, it is very likely that the QSOs are also strongly absorbed in the optical/near--infrared bands, i.e. they are also type~2 QSO. This is consistent with the fact that their (2--10 keV)--to--optical flux ratios vs. their 2--10 keV luminosities are in agreement with the relation described by Fiore et al. (\\cite{Fiore03}) for type~2 AGNs. In these 5 X--ray emitting EROs we find evidence for the presence of X--ray obscured type~2 QSO candidates.  For the remaining ERO, S2F1\\_493, the presence of an AGN is strongly  suggested by the high 0.5--2 keV luminosity and by the value of its  (0.5--2 keV)--to--optical flux ratio ($\\sim$1). In this case the hypothesis of very high intrinsic obscuration ($>$10$^{24}$ cm$^{-2}$) is not discarded. On the other hand, the quality of the data prevents us from placing reliable constraints on the AGN type hosted by this ERO. Due to its optical X--ray broad--band photometry, if an obscured AGN is present, it would have an X--ray luminosity in the QSO regime.  The only QSO host galaxy for which we have collected spectroscopic observations so far is an elliptical galaxy at z$\\simeq$1.7 with a stellar mass well above 10$^{11} M_\\odot$. This result suggests possible link between the  QSO activity and the formation of the bulk of the stars in massive spheroids."
    },
    "0410/astro-ph0410744_arXiv.txt": {
        "abstract": "Using the code {\\sc autostructure}, extensive calculations of inner-shell atomic data have been made for the chemical elements He, C, N, O, Ne, Na, Mg, Al, Si, S, Ar, Ca, Cr, Mn, Fe and Ni. The results are used to obtain up-dated opacities from the Opacity Project, OP. A number of other improvements on earlier work have also been included.  Rosseland-mean opacities from OP are compared with those from OPAL. Differences of 5 to 10\\% occur. OP gives the `$Z$-bump', at $\\log(T)\\simeq 5.2$, to be shifted to slightly higher temperatures. The opacities from OP, as functions of temperature and density, are smoother than those from OPAL.  The accuracy of the integrations used to obtain mean opacities can depend on the frequency-mesh used. Tests involving variation of the numbers of frequency points show that for typical chemical mixtures the OP integrations are numerically correct to within 0.1\\%.  The accuracy of the interpolations used to obtain mean opacities for any required values of temperature and density depend on the temperature--density meshes used. Extensive tests show that, for all cases of practical interest, the OP interpolations give results correct to better than 1\\%.  Prior to a number of recent investigations which have indicated a need for downward revisions in the solar abundances of oxygen and other elements, there was good agreement between properties of the sun deduced from helioseismology and from stellar evolution models calculated using OPAL opacities. The revisions destroy that agreement. In a recent paper Bahcall \\etal argue that the agreement would be restored if opacities for the regions of the sun with \\mbox{$2\\times 10^6 \\lesssim T \\lesssim 5\\times 10^6$ K} ($0.7$ to $0.4 R_{\\sun}$) were larger than those given by OPAL by about 10\\%. In the region concerned, the present results from OP do not differ from those of OPAL by more than 2.5\\%. ",
        "introduction": "First results for opacities from the Opacity Project (OP) were published by Seaton \\etal (1994, to be referred to as Paper I) and were in good general agreement with those from the OPAL project (Rogers and Iglesias, 1992) except for regions of high densities and temperatures. In a more recent paper, Badnell and Seaton (2003, Paper II) confirmed a suggestion made by Iglesias and Rogers (1995) that the discrepancies in those regions were due to the omission of important inner-shell processes in the OP work. Further comparisons between OPAL results and those from OP with inclusion of all important inner-shell processes are given by Seaton and Badnell (2004, Paper III) for the case of the 6-element mix (H, He, C, O, S and Fe) of Iglesias and Rogers (1995). In Papers II and III, the required inner-shell atomic data were computed using the code {\\sc autostructure} (Badnell, 1986, 1997). Similar computations of the required inner-shell atomic data for all cosmically-abundant elements % (He, C, N, O, Ne, Na, Mg, Al, Si, S, Ar, Ca, Cr, Mn, Fe and Ni) have now been made. Some further improvements and up-dates of the OP work are also considered in the present paper. ",
        "conclusions": "OP opacities have been up-dated by inclusion of all contributing inner-shell processes and some further refinements have been made.  Comparisons with OPAL show agreement generally to within 5 to 10\\%.  Using earlier estimates of metal abundances, good agreement was obtained between computed solar models and results obtained from helioseismology for $\\rcz$ and other data. With recent downward revisions in metal abundances that agreement is destroyed. It is shown by Bahcall \\etal (2004b) that agreement would be restored if opacities in the region just below the base of the convection zone, \\mbox{$2\\times 10^6 \\lesssim T \\lesssim 5\\times 10^6$ K} ($0.7$ to $0.4 R_{\\sun}$), were 10\\% larger than those from OPAL . In that region we find $\\kr($OP$)$ never to be larger than $\\kr($OPAL$)$ by more than 2.5\\%. Results for opacities in the region near the base of the convection zone have been discussed in some detail.  Studies have been made of the accuracies of the OP integrations over frequency, used to obtain Rosseland means, and of the OP procedures used for interpolations in temperature and density. It is concluded that the OP Rosseland-means should be numerically correct to within 1\\% or better."
    },
    "0410/astro-ph0410058_arXiv.txt": {
        "abstract": "We are developing a Compton telescope based on high resolution Si and CdTe imaging devices in order to obtain a high sensitivity astrophysical observation in sub-MeV gamma-ray region. In this paper, recent results from the prototype Si/CdTe semiconductor Compton telescope are reported. The Compton telescope consists of a double-sided Si strip detector (DSSD) and CdTe pixel detectors, combined with low noise analog LSI, VA32TA. With this detector, we obtained Compton reconstructed images and spectra from line gamma-rays ranging from 81~keV up to 356~keV. The energy resolution is 3.8~keV and 7.9~keV at 122~keV and 356~keV, respectively, and the angular resolution is $9.9^{\\circ}$ and $5.7^{\\circ}$ at 122~keV and 356~keV, respectively. ",
        "introduction": "Gamma-ray universe in the energy band of 100~keV to 10 MeV is a good observational window to study nucleosynthesis and particle acceleration in high energy astrophysical objects. However, in this energy band, observation sensitivity have been still limited, due to high background, difficulty in detecting photons, and incapability of imaging. A Compton telescope tracking the Compton scattering process of gamma-rays by position and energy sensitive detectors is difficult but a promising approach to bring a breakthrough in this energy band.  In the MeV region, COMPTEL\\cite{Schoenfelder} on-board {\\it CGRO} ({\\it Compton Gamma-Ray Observatory}) was the first successful Compton telescope and brought many pioneering results\\cite{Schoenfelder2}. In a Compton telescope, deposited energy and position of the gamma-ray interactions with the detector are acquired. When a gamma-ray photon is scattered in one detector and absorbed in another detector, the incident energy of the gamma-ray and the scattering angle can be determined as, \\begin{eqnarray} E_{\\rm in} &=& E_1 + E_2\\label{eq:comp_ene}\\\\ \\cos\\theta &=& 1 - \\frac{m_e c^2}{E_2} + \\frac{m_e c^2}{E_1 + E_2}\\label{eq:comp}, \\end{eqnarray} where $E_1$ is the energy of the recoil electron, $E_2$ is the energy of the scattered photon and $\\theta$ is the scattering angle. The latter equation presents the incident direction as a cone in the sky.  As shown in equation (\\ref{eq:comp}), incident direction of a gamma-ray is calculated from the position and the energy information of interactions. Therefore, employing semiconductor imaging spectrometers as components of Compton telescopes, with their good energy and position resolution, will drastically improve the angular resolution and hence the sensitivity. In this point of view, advanced Compton telescopes based on Si, Ge, CZT and CdTe are proposed and development is on-going in many groups\\cite{Takahashi04NAR,Takahashi03,MEGA,Kurfess,TIGRE,CZT}. A Compton telescope also enables us to measure polarization of gamma-rays, by obtaining the azimuthal distribution of Compton scattering\\cite{Kamae,polarimetry}.  Based on our CdTe and Si detector technologies\\cite{Takahashi02,Tanaka,Nakazawa,Fukazawa,Tajima03}, we are working on a new Si/CdTe semiconductor Compton telescope to explore the universe in the energy band from 100~keV to a few MeV. In this paper, we describe  experimental results obtained with a prototype Compton telescope which consists of one layer of a Double-sided Si strip detector (DSSD) and one layer of CdTe pixel detectors.  Detailed performance of the DSSD and the vision of next-generation gamma-ray missions using this technique are reported elsewhere\\cite{Fukazawa,Takahashi04}. Demonstration of polarization detection are reported in Mitani {\\em et al.}\\cite{Mitani04} and Tajima {\\em et al.}\\cite{Tajima04} ",
        "conclusions": "Si/CdTe semiconductor Compton telescope is attractive to explore the universe in sub-MeV gamma-ray region. We constructed a prototype Compton telescope which consists of one layer of DSSD and one layer of CdTe pixel detectors. With this prototype detector, we succeeded in Compton reconstruction of images and spectra in the energy band from 80~keV to 400~keV. The obtained angular resolution is $5.7^{\\circ}$ and $9.9^{\\circ}$ at 122~keV and 356~keV, respectively, and the energy resolution is 3.6~keV (FWHM) at 122~keV and 7.6~keV (FWHM) at 356~keV. The lower energy coverage down to 80~keV owes to good energy resolution and lower energy threshold of our detectors. We also succeeded in reconstruction of three-hit events by Compton scattering and K-escape, which becomes significant under some configuration of Compton telescopes. With these results, we have experimentally demonstrated the high capabilities of a Si/CdTe Compton telescope."
    },
    "0410/astro-ph0410091_arXiv.txt": {
        "abstract": "{ Observations of up to ten carbon monoxide (CO and isotopomers) transitions are presented to study the interstellar C$^{18}$O/C$^{17}$O ratio towards 21 positions in the nearby ($d$$\\sim$140\\,pc) low-mass star forming cloud $\\rho$~Oph. A map of the C$^{18}$O $J$=1--0 distribution of parts of the cloud is also shown. An average $^{12}$C$^{18}$O/$^{12}$C$^{17}$O isotopomeric ratio of 4.11 $\\pm$ 0.14, reflecting the $^{18}$O/$^{17}$O isotope ratio, is derived from Large Velocity Gradient (LVG) calculations. From LTE column densities we derive a ratio of 4.17$\\pm$0.26. These calculations also show that the kinetic temperature decreases from about 30~K in the cloud envelope to about 10~K in the cloud cores. This decrease is accompanied by an increase of the average molecular hydrogen density from 10$^4$~cm$^{-3}$ to $\\ga$10$^5$~cm$^{-3}$. Towards some lines of sight C$^{18}$O optical depths reach values of order unity. ",
        "introduction": "Abundance ratios of interstellar isotopomers are a powerful tool to study the chemical evolution of the Galaxy. One such ratio is that of the rare species of oxygen, $^{18}$O and $^{17}$O, as measured from the isotopomers of CO. For the galactic disk and -center region, Penzias (\\cite{penzias}) reported average $^{18}$O/$^{17}$O ratios of 3.65$\\pm$0.15 and 3.5$\\pm$0.2, respectively. He also found that the $^{18}$O/$^{17}$O ratio, determined from the integrated line intensity ratios $\\int$T[$^{12}$C$^{18}$O(1$-$0)]d$v$/$\\int$T[$^{12}$C$^{17}$O(1$-$0)]d$v$, shows no significant gradient with galactocentric distance $R$ out to 10~kpc (solar circle: $R_0$=8.5\\,kpc): the $^{18}$O/$^{17}$O ratios of the galactic disk and -center are, within the limits of observational accuracy, identical. Models of the chemical evolution of the Galaxy by Prantzos et al. (\\cite{prantzos}) suggest that after a few Gyr the ratios in the Galaxy should be independent of galactocentric radius. There is, however, a discrepancy between the interstellar medium (ISM) values and the much higher (5.5; Anders \\& Grevesse \\cite{anders}) solar system one. Heikkil\\\"a et al. (\\cite{heikkila}) obtain a low value of 1.6$\\pm$0.3 in the LMC, while in the nuclear starbursts NGC\\,253 and NGC\\,494 5 $^{18}$O/$^{17}$O $\\sim$ 6.5 (Harrison et al. \\cite{harrison}; Wang et al. \\cite{wang}). These results suggest that the $^{18}$O/$^{17}$O ratio depends on metallicity.  The sources observed by Penzias are located in a limited range of galactocentric radius (and therefore metal abundance), and we (Wouterloot et al., in preparation) have reobserved these sources with higher angular resolution and have extended our study to sources out to $R$=16\\,kpc. While Penzias (\\cite{penzias}) only observed the $J$=1--0 transition, our observations also include the $J$=2--1 rotational lines. The goal of the present paper is to study in detail excitation and opacity effects that could affect the measured $^{18}$O/$^{17}$O ratios and radial gradients on small scales. These effects are usually either ignored or physical parameters are derived by assuming a fixed $^{18}$O/$^{17}$O ratio, so that a more careful study is desirable.  We have chosen the \\object{$\\rho$~Oph cloud} ($d$$\\sim$140\\,pc) because of the large range in column densities found therein, and because it is in the solar neighbourhood so that a high linear resolution can be attained towards this object. Twenty one positions were selected for observations from a C$^{18}$O(1--0) map of Wilking \\& Lada (\\cite{wilking}) to have a range of C$^{18}$O(1--0) intensities. Towards the positions with the strongest lines we not only observed the $J$=1--0 and 2--1 lines of four isotopomers ($^{12}$C$^{16}$O, hereafter $^{12}$CO; $^{13}$C$^{16}$O, hereafter $^{13}$CO; $^{12}$C$^{18}$O, hereafter C$^{18}$O; $^{12}$C$^{17}$O, hereafter C$^{17}$O), but we also measured the $J$=3--2 lines of C$^{18}$O and C$^{17}$O. ",
        "conclusions": "In calculating the isotope ratios we did not distinguish between the different velocity components in the cloud (with typical velocity differences of about 2~kms$^{-1}$). The reason is the non-Gaussian line shape of many components which are confused by the hyperfine structure of C$^{17}$O (3 components for $J$=1--0 and 9 components for $J$=2--1).  Assuming a depth along the line of sight of 0.4 pc (10\\arcmin ), the crossing time for a velocity of 2~kms$^{-1}$ would be 2.0 10$^5$ yr, which is probably much less than the lifetime of the cloud. This is at least 5.5 10$^6$ yr, the age of the Upper Scorpius subgroup of the Scorpius-Centaurus OB association (de Geus et al. \\cite{geus}). Wilking et al. (\\cite{wilkinglada}) obtained an upper limit for the ages of the T Tauri stars in $\\rho$ Oph of 3.0 10$^6$ yr. This would provide enough time for sufficient mixing of the isotopic constituents of the gas. In addition the types of stars which produce $^{17}$O and $^{18}$O are not (yet) present in the $\\rho$ Oph cloud (they do exist in the nearby Upper Scorpius subgroup of the older Scorpius-Centaurus association), so both isotopes are not locally produced and it is unlikely that they are enhanced in some parts of the cloud by this mechanism (see e.g. Henkel \\& Mauerberger \\cite{henkelmau}). It also seems unlikely that stellar winds from these associations can alter the composition of the cloud significantly.  The results for the C$^{18}$O/C$^{17}$O ratios are summarized in Table~\\ref{meanratios}, with the transition and the method used.  The first three entries in Table~\\ref{meanratios} indicate that the observed ratio is very dependent on the transition used. We note that Penzias (\\cite{penzias}) corrected his ratios only for the difference in frequency and not for optical depth effects. This suggests that also the $J$=1--0 ratios used in Penzias' (\\cite{penzias}) galactic study need some more analysis. This will be discussed in more detail together with new measurements in a forthcoming paper. Within the uncertainties the average $J$=1--0 value is equal to the number obtained by Penzias (\\cite{penzias}) for the galactic plane, 3.65$\\pm$0.15. The weighted average ratio of our LTE ratios is 4.17$\\pm$0.26, whereas the LVG calculations resulted in a ratio of 4.11$\\pm$0.14. Because this method combines all observational data we consider this the best result. Bensch et al. (\\cite{bensch}) detected the almost certainly optically thin transitions $^{13}$C$^{18}$O and $^{13}$C$^{17}$O(1--0) towards pos. 1. Their observed $^{13}$C$^{18}$O/$^{13}$C$^{17}$O ratio, corrected for the frequency difference is 4.23$\\pm$0.53. Using escape probability models they derived a ratio $^{18}$O/$^{17}$O of 4.15$\\pm$0.52, which is identical to (but with a relatively large uncertainty because only one position was observed) the presently derived average value. Note that the position used by Bensch et al. (\\cite{bensch}) had to be omitted from our calculations because of an undetermined optical depth.  Recently Ladd (\\cite{ladd}) observed C$^{18}$O and C$^{17}$O(1--0) towards some 600 positions in the Taurus clouds and derives a ratio of 4.0$\\pm$0.5, in agreement with the present result. However, the C$^{18}$O/C$^{17}$O ratio appears to decrease with increasing integrated C$^{17}$O(1$-$0) intensity. Ladd then concludes that the ratio in the inner parts is 2.8$\\pm$0.4, due to larger self shielding of C$^{18}$O in the outer parts. The range of integrated C$^{17}$O(1--0) intensity in Taurus (0.2 to 0.6 Kkms$^{-1}$) is much smaller than in $\\rho$~Oph because of the smaller line widths. We think that this ratio of 2.8 is not real: our LVG models (using a smaller velocity gradient of 2.0 km\\,s$^{-1}$\\,pc$^{-1}$ and an intrinsic C$^{18}$O/C$^{17}$O ratio of 4.0) can reproduce the decrease in ratio with $I$(C$^{17}$O) if positions with lower $I$(C$^{17}$O) have a lower density and maybe kinetic temperature than the points with higher $I$(C$^{17}$O), which is quite possible. In addition, it appears that when self-shielding is significant for C$^{18}$O, at $N$(H$_2$) $>$ 1.22 10$^{21}$ (Frerking et al. \\cite{frerking}), or $A_{\\rm v}$ $>$ 1.3 mag (Bohlin et al. \\cite{bohlin}), this corresponds to a far-UV extinction by dust of $>$13 mag (Aannestad \\& Purcell \\cite{aannestad}). This implies that beyond this extinction (the (C$^{17}$O) self-shielding H$_2$ column density is even higher) there is too little UV radiation left to affect the C$^{18}$O/C$^{17}$O ratio, in contradiction with the suggestion by Ladd (\\cite{ladd}). In $\\rho$~Oph a systematic decrease of C$^{18}$O/C$^{17}$O with $I$(C$^{17}$O) is not seen (see Fig. \\ref{coratios}), except towards pos. 1.  \\begin{table} \\caption[]{Derived C$^{18}$O/C$^{17}$O ratios compared with some previous results. \\label{meanratios}} \\begin{flushleft} \\begin{tabular}{lrll} \\hline\\noalign{\\smallskip} Ratio & trans. & method & omit pos.\\\\ C$^{18}$O/C$^{17}$O & & \\\\ \\hline\\noalign{\\smallskip} 3.53$\\pm$0.11 & $J$=1$-$0 & freq. corr. & 1 \\\\ 3.06$\\pm$0.11 & $J$=2$-$1 & freq. corr. & 1 \\\\ 2.78$\\pm$0.18 & $J$=3$-$2 & freq. corr. & 1 \\\\ 4.07$\\pm$0.32 & $J$=1$-$0 & from $N_{\\rm lte}$ & 1,3,4,6\\\\ 4.35$\\pm$0.44 & $J$=2$-$1 & from $N_{\\rm lte}$ & 1,2,3,4,6,\\\\ &           &        & 7,9,11 \\\\ 4.17$\\pm$0.26 & \\multicolumn{3}{l}{average of $J$=1--0 and 2--1} \\\\ &           &from $N_{\\rm lte}$ & \\\\ 4.11$\\pm$0.14 & \\multicolumn{3}{l}{$J$=1--0,2--1,3--2; from LVG}  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 3.65$\\pm$0.15 & $J$=1$-$0 & Penzias (\\cite{penzias}) gal. disk & \\\\ &           & freq.corr.& \\\\ 4.15$\\pm$0.52 & $J$=1$-$0 & Bensch et al. (\\cite{bensch}) \\\\ &           & from $^{13}$C$^{18}$O/$^{13}$C$^{17}$O \\\\ &           & from exc. model\\\\ 4.23$\\pm$0.53 & $J$=1$-$0 & Bensch et al. (\\cite{bensch}) \\\\ &           & from $^{13}$C$^{18}$O/$^{13}$C$^{17}$O \\\\ &           & freq. corr.\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table}"
    },
    "0410/astro-ph0410572_arXiv.txt": {
        "abstract": "We observed Hubble V, the brightest HII region complex in the dwarf irregular galaxy NGC~6822, at near-infrared (1.8--2.4 $\\micron$) wavelengths using the Cooled Grating Spectrometer 4 (CGS4) at the United Kingdom Infra-Red Telescope (UKIRT). The line emission maps of Hubble V show the typical structure of a photo-dissociation region (PDR) where an ionized core, traced by compact He I emission (2.0587 $\\micron$) and Br$\\gamma$ emission (2.1661 $\\micron$), is surrounded by an outer layer traced by molecular hydrogen (H$_2$) emission. The measured line ratios of $\\hto$ (2.2477 $\\micron$) / $\\hozs$ (2.1218 $\\micron$) from 0.2 to 0.6 and the un-shifted and un-resolved line profiles suggest that the H$_2$ emission originates purely from a photo-dissociation region (PDR). By comparing the H$_2$ results with a PDR model, we conclude that Hubble V includes dense ($10^{4.5} \\cmv$) and warm PDRs. In this environment, most of the H$_2$ molecules are excited by far-UV photons (with a field strength of $10^{2-4}$ times that of the average interstellar field), although collisional processes de-excite H$_2$ and contribute significantly to the excitation of the first vibrational level. We expect that Hubble V is in the early stage of molecular cloud dissolution. ",
        "introduction": "The dwarf irregular galaxy NGC~6822 is a member of the Local Group. Because of its proximity ($d = 500$ kpc; \\citealt{mca83}), we can resolve its molecular clouds and star forming regions on parsec scales (1 arcsec $\\simeq$ 2.4 pc). The Large Magellanic Cloud (LMC) and the Small Magellanic Cloud (SMC) are much closer to us than NGC~6822. However, they reside so close to the Galaxy that the star forming process may be influenced by the tidal force associated with the Galactic gravitational field.  On the contrary, NGC~6822 is far more isolated and star formation processes are dictated only by the local conditions in NGC~6822 itself.  The optical view of NGC~6822 is dominated by a large, well defined, central bar and many bright HII regions and OB associations.  The brightest and largest HII region complexes, Hubble I, III, V, and X \\citep{hub25}, are located at the northern end of the bar. The oxygen abundance measured in the HII regions (12 + log(O/H) = 8.23; \\citealt*{leq79,pag80,ski89}) is 2 times smaller than the Galactic value (12 + log(O/H) = 8.52 in Orion; \\citealt{pei77}) and between those in the LMC (12 + log(O/H) = 8.43) and the SMC (12 + log(O/H) = 8.02; see \\citealt{duf84} and the references therein).  With a total visible extent of 50 pc (about 20 arcsec on the sky), Hubble V is the brightest HII region complex in NGC~6822.  Visual images show the structure of the bright core and the large, diffuse halo that surrounds the core and extends towards the northwest (\\citealt*{ode99,isr03}; see Fig.~\\ref{fig_imap}).  The core contains a compact cluster of bright blue stars.  The halo is overlaid on to the eastern part of the OB association, Hodge OB 8 \\citep{ode99}. The age of the Hubble V complex is about 4 Myr and there is no evidence for multiple star formation events in the past \\citep{ode99,bia01}.  \\citet{wil94} observed Hubble V in CO emission at high spatial resolution (6.2 x 11.1 arcsec) and reported finding a molecular cloud complex in the Hubble V region (Fig.~\\ref{fig_imap}(a)). According to the visual image (Fig.~\\ref{fig_imap}(b)) based on a multi-color composite from the Hubble Space Telescope (HST), it is likely that dark clouds surround all parts of the core cluster except in the west. \\citet{isr03} found a compact source of $K$-band emission to the south of the visual core. This $K$-band peak is not seen in the visual images so they argued that it could be another compact star cluster that is highly obscured by the dark cloud.  In this paper, we present the results of our near-infrared (near-IR) observations of NGC~6822 Hubble V at high spatial resolutions. The observations and data reduction are described in Section~\\ref{observation}. We present the morphology of the ionized region and photo-dissociation region (PDR) and the physical conditions in the PDR in Section~\\ref{results}. In Section~\\ref{discussion} we derive some physical parameters by comparing our observations with predictions of a PDR model \\citep{ste89}. We also compare with previous CO observations to illustrate the structure of the molecular clouds and discuss the evolution of the Hubble V star forming region.    \\begin{figure*} \\special{psfile=\"fig1.eps\" angle=0 vscale=78 hscale=78 voffset=-545 hoffset=5} \\vspace{193mm} \\caption{(a) Blue optical image and contours of CO $J = 1 - 0$ integrated intensity (from \\citealt{wil94}) over an area covered by our slit scanning observations with the 40 line mm$^{-1}$ (hereafter l/mm) grating. The contour intervals are (-4, -3, -2, 2, 3, 4, 5) $\\times$ Jy beam$^{-1}$ (1 $\\sigma$). (b) The visual image of NGC~6822 Hubble V observed by \\citet{ode99} and \\citet{bia01} using the NASA/ESA/STScI Hubble Space Telescope (HST). The four other figures of contours are reconstructed from the scanning of 12 slit positions. The horizontal axis is along the slit and each row corresponds to each slit. These contour maps show views of NGC~6822 Hubble V in (c) near-infrared continuum and integrated line emission of (d) atomic hydrogen, (e) helium, and (f) molecular hydrogen. Contour levels are linear and increase from 1-$\\sigma$ RMS noises with the same intervals; (c) $9 \\rm \\times 10^{-18}\\;W\\;m^{-2}\\;\\micron^{-1}\\;arcsec^{-2}$, (d--f) $1 \\rm \\times 10^{-19}\\;W\\;m^{-2}\\;arcsec^{-2}$. The open circle at (x, y) = (38.5, 6.0) pixel marks the position $\\rm \\alpha = 19^h 44^m 52\\fs85, ~ \\delta = -14\\degr43\\arcmin12\\farcs8$ (J2000) on the sky and the pixel scale is 1.22 arcsec $\\times$ 1.2 arcsec. The dotted boxes indicate the areas over which the spectra in Fig.~\\ref{fig_40spectrum} are averaged. The solid arrow in (f) shows the direction of the slit used for the echelle observations. The physical conditions at the five positions along the slit (marked by the ticks; hereafter positions A, B, C, D, and E from the north to the south) are discussed in the text. \\label{fig_imap}} \\end{figure*} ",
        "conclusions": "We performed near-IR spectroscopic observations of the Hubble V complex, the brightest HII region complex in the dwarf irregular galaxy NGC~6822.  From low spectral resolution ($\\lambda / \\Delta \\lambda \\simeq 800$) $K$-band (1.8 - 2.4 $\\micron$) slit scanning, we obtained high spatial resolution ($\\sim$ 1 arcsec) maps of the region in 2.2 $\\micron$ continuum, 2.0587 $\\micron$ He I, 2.1661 $\\micron$ Br$\\gamma$, and 2.1218 $\\micron$ $\\hoz$ emission. The morphology of Hubble V in the near-IR is typical of HII regions/PDRs; the compact He I and Br$\\gamma$ emitting region is surrounded by the H$_2$ emitting region. The detailed distribution of the H$_2$ emission in Hubble V is observed for the first time.  Our high spectral resolution observations ($\\lambda / \\Delta \\lambda \\simeq 15000$) of the $\\hoz$ and $\\htos$ lines, when combined with our excitation analysis, indicate that the region around the core of Hubble V is a dense PDR and suggest that there is no significant shock activity. The moderate $\\htos$ / $\\hozs$ ratios (0.2 -- 0.6) are explained by high densities ($n_{H_2} \\ga 5 \\times 10^4 \\; \\cmv$) and the possibility of shocks are excluded by the gas kinematics. The H$_2$ lines have the same systematic velocity as the cold molecular gas (traced by CO) and the line profiles are spectrally unresolved. This conclusion implies that there is no detectable YSO activity around the core of Hubble V.  By comparing the observed results with a PDR model, we estimate a gas density of $n_T \\, \\simeq \\, 10^{4.5} \\, \\cmv$ and incident UV field strength of $\\chi = 10^{2-4}$. This means that the environment around the core of Hubble V is dense and warm enough so that most of the H$_2$ molecules excited by far-UV photons are collisionally de-excited and the excitation to the $\\upsilon = 1$ level is dominated by collisions. The physical parameters estimated in this work are consistent with those independently derived by \\citet{isr03}; the strength of the UV field estimated from radio continuum flux densities and the gas density estimated from CO line ratios assuming a PDR environment.  The distribution of the near-IR continuum and the line emission from the ionized regions in the southern part of Hubble V confirm the existence of a hidden star cluster inside a dark molecular cloud. Hubble V seems to be in the early stage of molecular cloud dissolution after having finished its star formation activity $\\sim 4$ Myr ago.  But the progress of evolution seems slower than in the Galaxy, due to the intrinsically larger sizes and masses of the molecular clouds."
    },
    "0410/astro-ph0410566_arXiv.txt": {
        "abstract": "We present full Stokes observations from SPINOR in the \\ion{Ca}{2} infrared triplet and the \\ion{He}{1} multiplet at 1083~nm from which some properties of spicules have been derived. There are important advantages in multi-line observations, particularly from different elements. We find that the orientation of the plane of polarization is very different for the Ca and He lines, which provides the first direct model-independent evidence for magnetic fields in spicules. Our data shows that the Ca and He lines have almost identical widths. Since the Ca atom is 10 times heavier than He, we are able to conclude that most of the broadening is non-thermal ($\\simeq$16~km~s$^{-1}$) and to set an upper limit of 13~kK to the spicular temperatures. The bisectors of the lines span a velocity range of over 15~km~s$^{-1}$ for the He line and 30~km~s$^{-1}$ for the Ca ones. The vertical gradient of line-of-sight velocities is also very different for both elements. We obtain 2.8~km~s$^{-1}$~Mm$^{-1}$ from He versus 6.4~km~s$^{-1}$~Mm$^{-1}$ from Ca. These properties, and others from similar observations, should be taken into account in future physical models of spicules. ",
        "introduction": "\\label{sec:intro}  Spicules are elongated, nearly-vertical structures visible in emission above the solar limb in chromospheric lines. First observed over 100 years ago (\\citeNP{S1877}), spicules have thus far eluded our efforts to understand their physical nature and origin. Once viewed as a very attractive problem for solar physicists, interest seemed to fade away somewhat since the early 1970s. Most of the knowledge gathered until then, mainly from spectroscopic observations of H$\\alpha$, He D$_3$ or Ca~H and~K, is compiled in a comprehensive review by \\citeN{B68}. Since then, very few works have been published with new observational properties of spicules and some authors have concentrated instead in novel modeling efforts (the most recent by \\citeNP{DPEJ04}).  However, this scenario seems to be about to change drastically. Modern spectro-polarimeters capable of observing chromospheric lines off of the solar limb are providing new observational constrains that show considerable promise to expand our knowledge on spicules (and other chromospheric structures). In addition to the present work, there are two other papers in press with new spectro-polarimetric observations of \\ion{He}{1} lines, namely the 1083~nm multiplet (\\citeNP{TBMC+04}) and D$_3$ (\\citeNP{LAC04}). Hopefully, this surge is only the beginning of a renewed interest in spicule observations. New data is needed both to inspire and to constrain future theoretical models.  We present here for the first time observations of spicules in the \\ion{Ca}{2} infrared triplet (in particular the 849.8 and 854.2~nm lines), alongside the \\ion{He}{1} multiplet at 1083~nm. The dataset analyzed here was recorded using the new SPINOR (Spectro-Polarimeter for INfrared and Optical Regions) instrument, at the Dunn Solar Telescope (DST) in the Sacramento Peak Observatory. We believe that multi-line observations from different elements (ideally with very different atomic weights) have important advantages to determine some magnetic, thermal and dynamic parameters of spicules. For example, we present evidence of the existence of a magnetic field in spicules without having to resort on complex Hanle effect modeling. In this sense, our work can be considered the first model-independent verification of the magnetic nature of spicules (although such fields have been obtained from Hanle effect measurements by \\citeNP{TBMC+04}; \\citeNP{LAC04}). Since the thermal broadening depends strongly on the atomic weight, multi-line observations are also helpful to set boundaries to this type of broadening. ",
        "conclusions": "The observations that we present in this work provide some new insights into the physical origin of spicules. By observing lines of both Ca and He we are able to set constrains on the non-thermal broadening of the lines and the temperature of the emitting plasma. The change in the orientation of the polarization plane in the Ca and He lines is a convincing model-independent evidence for magnetic fields in prominences. A quantitative determination of the field strength and orientation would require detailed modeling of the line formation, which is beyond the scope of the present work. The different behavior of the emission wings in He and Ca (broader and narrower than a Gaussian, respectively) is certainly intriguing. It seems appealing to consider this as a clue for the nature of the non-thermal broadening mechanism.  In spite of recent efforts, spicules still remain poorly understood and we feel that more data are urgently needed. We plan to expand the results presented here in a forthcoming paper with new observations and more detailed analyses. In particular we intend to use an infrared camera and the new SPINOR modulator to obtain better signal-to-noise ratios in the infrared (which, as stated above is one of our main limitations here). Detailed Non-LTE modeling would be desirable and it seems to be well understood now, at least for the Ca lines both in the Zeeman (\\citeNP{SNTBRC00a}) and Hanle (\\citeNP{MSTB01}) regimes."
    },
    "0410/astro-ph0410085_arXiv.txt": {
        "abstract": "The pursuit of the core collapse supernova explosion mechanism continues. While such efforts have been undertaken over the last four decades, it is only in the last decade that multidimensional models have been developed, and only in the last few years that significant components of supernova models, such as the neutrino transport, have been modeled with sufficient realism. We can now identify what are arguably the important components that must be included in any realistic three-dimensional supernova model. The challenge we now face is to include all of these in the models. ",
        "introduction": "Four hundred years after the observation of SN1604 by Galileo Galilei, the mechanism for core collapse supernova explosions remains unknown. Nonetheless, one-, two-, and three-dimensional simulations performed thus far have shown that (1) neutrino transport, (2) fluid instabilities, (3) rotation, and (4) magnetic fields, together with proper treatments of (5) the sub- and super- nuclear density stellar core equation of state, (6) the neutrino interactions, and (7) gravity will be important. This applies to the explosion mechanism  and to phenomena associated with supernova explosions, such as neutron star kicks, supernova neutrino and gravitational wave emission, and the polarization of supernova light. Not surprisingly, current two- and three-dimensional models have yet to include (1)--(4) with sufficient realism. One-dimensional spherically symmetric models have achieved a significant level of sophistication but, by definition, cannot incorporate (2)--(4), except phenomenologically. Fully general relativistic spherically symmetric simulations with Boltzmann neutrino transport do not yield explosions \\citep{lmtmhb01}, demonstrating that some combination of (2), (3), and (4) is required to achieve this. (While the neutrino interactions and equation of state are important, we do not expect changes in them to lead to {\\it qualitative} changes in the models.) It is important to note that layering the dimensionality and the physics will be needed to achieve a complete understanding of the supernova mechanism and phenomenology \\citep{mlmhtb01}. The above mentioned one-dimensional simulations, for example, provided the first proof that (2), (3), and (4), in some combination, are {\\it necessary} ingredients in the recipe for explosion.  The past modeling efforts alluded to above have illuminated that core collapse supernovae may be neutrino driven, MHD driven, or both \\citep{s84,wm93,hbhfc94,bhf95,jm96,mw99,fw04,janka04}, but uncertainties in the current models prevent us from being able to single out any one of these possibilities. And it may be that more than one possibility is realized in Nature. Nonetheless, were a supernova neutrino driven, magnetic fields will likely have an impact on the dynamics of the explosion. Similarly, were a supernova MHD driven, the neutrino transport will dictate the dynamics of stellar core collapse, bounce, and the postbounce evolution, which in turn will create the environment in which the MHD-driven explosion would occur. Thus, while reduction will allow us to sort out the roles of each of the major physics components listed above, we will not obtain a quantitative, and perhaps even qualitative, understanding of core collapse supernovae until all components and their coupling are included in the models. ",
        "conclusions": "Much work remains to determine the core collapse supernova explosion mechanism and to understand all of the phenomena associated with these explosions. Despite significant progress and the accumulation of a wealth of understanding over the last four decades, as well as a number of specific clues, we still cannot answer some of the most fundamental questions. And, certainly, multidimensional supernova modeling is only in its infancy. However, guiding observations are becoming increasingly sophisticated and supercomputers are becoming increasingly powerful. Consequently, we are now presented with a unique opportunity to respond to this challenge."
    },
    "0410/astro-ph0410750_arXiv.txt": {
        "abstract": "We present an analysis of stellar distributions for the young stellar clusters GGD~12-15, IRAS~20050+2720, and NGC~7129, which range in far-IR luminosity from 227 to 5.68~$\\times 10^{3}~L_{\\odot}$ and are all still associated with their natal molecular clouds.  The data used for this analysis includes near-IR data obtained with FLAMINGOS on the MMT Telescope and newly obtained wide-field 850~$\\mu$m emission maps from SCUBA on the JCMT. Cluster size and azimuthal asymmetry are measured via azimuthal and radial averaging methods respectively.  To quantify the deviation of the distribution of stars from circular symmetry, we define an azimuthal asymmetry parameter and we investigate the statistical properties of this parameter through Monte Carlo simulations.  The distribution of young stars is compared to the morphology of the molecular gas using stellar surface density maps and the 850~$\\mu$m maps. We find that two of the clusters are not azimuthally symmetric and show a high degree of structure.  The GGD~12-15 cluster is elongated, and is aligned with newly detected filamentary structure at 850~$\\mu$m.  IRAS~20050+2720 is composed of a chain of three subclusters, in agreement with \\citet {chen97}, although our results show that two of the subclusters appear to overlap. Significant 850~$\\mu$m emission is detected toward two of the subclusters, but is not detected toward the central subcluster, suggesting that the dense gas may already be cleared there.  In contrast to these two highly embedded subclusters, we find an anti-correlation of the stars and dust in NGC~7129, indicating that much of the parental gas and dust has been dispersed.  The NGC~7129 cluster exhibits a higher degree of azimuthal symmetry, a lower stellar surface density, and a larger size than the other two clusters, suggesting that the cluster may be dynamically expanding following the recent dispersal of natal molecular gas. These analyses are further evidence that embedded, forming clusters are often not spherically symmetric structures, but can be elongated and clumpy, and that these morphologies may reflect the initial structure of the dense molecular gas.  Furthermore, this work suggests that gas expulsion by stellar feedback results in significant dynamical evolution within the first 3~Myr of cluster evolution. We estimate peak stellar volume densities and discuss the impact of these densities on the evolution of circumstellar disks and protostellar envelopes. ",
        "introduction": "It is now commonly accepted that most stars form in clusters of 100 or more stars \\citep{ll03,porr03,carp00}.  In addition, young OB stars nearly always appear in groups and clusters of lower mass stars \\citep{tpn99,mege02}. Clearly, understanding the process of forming stars in clusters is of considerable importance.  Although we have an increasingly detailed picture of relatively isolated star formation in nearby dark clouds such as those in the Taurus complex, our understanding of formation in clusters has been hindered by greater observational and theoretical challenges.  Observationally, the challenges of studying embedded clusters are their distance, their spatial density, and their association with high column density molecular clouds.  While the Taurus complex provides a number of examples of isolated star formation at 160~pc, most clusters are at distances of 300~pc or greater \\citep{ll03,porr03}, and thus require greater angular resolution and sensitivity to study. The stars in these regions may be closely spaced, again requiring high angular resolution to resolve the individual stars.  Finally, many clusters are deeply embedded, and are thus obscured to observation at visible wavelengths.  For this reason, the importance of young stellar clusters was not realized until infrared cameras with detector arrays became available \\citep{lada92}.   With the latest generation of wide field cameras on 6-10 meter class telescopes, near-IR observations can now probe a large sample of young clusters with the sensitivity to detect objects well below the hydrogen burning limit, the angular resolution to resolve high density groupings of stars, and the field of view necessary to observe the distribution of stars over multi-parsec distances.  Furthermore, the {\\it Spitzer} Space Telescope is now providing detailed images of young clusters in the mid-IR, allowing us for the first time to identify young stars with disks and infalling envelopes efficiently in clusters out to 1~kpc and beyond \\citep[e.g.,][]{mege04,whit04}.  Theoretically, the problem of star formation in clusters is perhaps an even greater challenge, as a realistic analysis must include a network of non-linear processes, such as turbulent motions and shocks in the natal molecular cloud, potential collisions and multi-body interactions between young stars and protostars, and the destructive influence of radiation and winds from the newly formed stars.  Numerical models of turbulent fragmentation have been successful at simulating the formation of clusters (or at least clusters of protostellar gravitationally bound cores) and these simulations are beginning to yield a wealth of theoretical predictions which can now be investigated observationally \\citep[e.g.,][]{kb00,khm00,bbb03,bbv03,gamm03,li04}. For example, numerical simulations by \\citet{bbb03} and \\citet{bbv03} of turbulent, initially spherically symmetric, Jeans unstable molecular clouds produce collapsing, non-uniform filamentary density structures which are formed through the dissipation of kinetic energy by hydrodynamic shocks, allowing gravity to overcome turbulent support and force collapse.  In these simulations, the gas eventually fragments into small ($<200$~AU), self-gravitating, sub-virial clumps, which are interpreted as protostars.  The models produce multiple groups of protostars which eventually merge to form large clusters. However, these simulations have obvious limitations: there are numerical problems associated with following the collapse of gas from parsec to 100~AU size scales, the models assume idealized initial conditions, and the models also do not take into account the feedback from the newly formed stars.  In particular, the later stages of the evolution in these models may have questionable relevance to real clusters given that most young stellar clusters do not evolve into bound clusters, in part due to the dispersal of most of the binding mass before it forms stars \\citep{ll03}.  With increasingly sophisticated numerical models beginning to predict the spatial distribution and densities of young stars, it is an appropriate time to explore the configuration of young stellar clusters through wide-field infrared and submillimeter imaging, both to test current models and to guide the development of future models.  Through the analysis of the stellar surface density in young ($<3$~Myr) stellar clusters still associated with their natal molecular gas, and through a comparison of the corresponding dust and gas density structure, we can study the initial configuration of young clusters and their subsequent dynamical evolution.  We are now engaged in a multi-wavelength study of the structure of young stellar clusters, including ground-based near-IR data, {\\it Spitzer} mid-IR imaging, and ground-based submillimeter and millimeter-wave imaging \\citep{gute04,mege04}. While testing cluster formation models is a long-term goal of the project, this paper is primarily focused on developing quantitative techniques for analyzing the spatial structure of young stellar clusters and applying these techniques to new observations of three clusters.   Subsequent papers will apply these techniques to the full sample of clusters observed in our survey. It is our hope that this work will encourage modelers to provide more accessible observable quantities from their simulations.  In this paper, we present an initial analysis of near-IR and submillimeter imaging of three clusters using near-IR images from FLAMINGOS on the 6.5~meter MMT telescope and 850~$\\mu$m continuum maps from SCUBA on the JCMT.  We restrict our near-IR analysis to an $H-K$ color-based extinction analysis and the distribution of sources detected in the $K$--band.  The $K$--band is the most sensitive to embedded young stars \\citep{mege94}, and thus gives the most unbiased view of the distribution of point sources.   Although the membership of individual sources cannot be established by the $K$--band data alone, statistical methods can be used to infer the overall distribution of young sources.  This approach is particularly appropriate for studying the densest parts of young clusters, where the density of young stars dominates over the density of field stars in the line of sight. In a future paper, we will use data from {\\it Spitzer}'s Infrared Array Camera to find the distribution of sources with infrared excess emission in all three of the regions presented; this approach is particularly effective in probing regions with a lower density of young stars than the background stellar density.  The three cluster regions chosen for this analysis exhibit signs of ongoing star--formation such as active molecular outflows, H$_{2}$O masers, and associated IRAS point sources with significant far-IR luminosities ranging from 227~to~$5.68 \\times 10^{3}~L_{\\odot}$ (see Table~\\ref{littable}).  A brief summary of the current knowledge of these three regions follows.  \\subsection{GGD~12-15\\label{ggd1215}} GGD~12-15 is a young cluster region associated with the Monoceros molecular cloud, located at a distance of 830~pc \\citep{hr76}.  It was originally identified from several patches of optical reflection nebulosity with associated Herbig-Haro (HH) objects \\citep{ggd78}.  There is a wide-spread (1.13~pc) stellar cluster detectable at near-IR wavelengths with an estimated 134~objects \\citep{hoda94,carp00} and a moderate velocity bipolar molecular outflow oriented northwest-southeast centered on an H$_{2}$O maser and secondary peak in 800~$\\mu$m emission \\citep{lhd90}. A compact H~II region, with flux consistent with excitation by a B0.5 zero--age main sequence star, was detected in radio continuum emission \\citep{rodr80}.  The primary peaks of 800~$\\mu$m dust emission \\citep{lhd90}, $^{13}$CO emission, and C$^{18}$O emission \\citep{ridg03} are also approximately centered on the H~II region.  Furthermore, a cluster of point-like radio continuum sources has been detected in the vicinity of the H~II region \\citep{grg00,grg02}.  Some of these objects have near-IR and mid-IR spectral energy distributions consistent with embedded protostars \\citep{pt03}.  \\subsection{IRAS~20050+2720} IRAS~20050+2720 is a relatively small but very active site of embedded star formation originally identified as an IRAS point source.  It is associated with the Cygnus Rift, located at a distance of 700~pc \\citep{wilk89}.  Multipolar outflows have been found centered on the main IRAS source position, one of particularly high velocity (+/-~70~km/s) \\citep{bft95}.  \\citet{chen97} performed a near-IR investigation of this object (completeness limit of $K =$~16.2), and reported an estimated 100 cluster objects arranged in three tightly packed subclusters which they labeled A, B, and C.  The most reddened of these is subcluster~A, located at the IRAS point source position, as well as the peak of $^{13}$CO and C$^{18}$O emission \\citep{ridg03}.  The area around this subcluster has significant near-IR nebulosity, and two embedded sources were detected in $L$--band (3.5~$\\mu$m) and narrow-band $M$ (4.8~$\\mu$m) imaging that were not detected at shorter near-IR wavelengths \\citep{chen97}.  \\citet{chin01} performed a submillimeter and millimeter study and found a strong emission source at the IRAS point source position, with a ridge extending to the south. They also detected a diffuse ridge to the northwest at 1.3 mm and three point--like millimeter sources outside the 2.3$^{\\prime}$ field of view of their 450~$\\mu$m and 850~$\\mu$m maps.  These are likely protostars or pre-stellar cores associated with the cluster.  \\subsection{NGC~7129}  NGC~7129 is a large region of reflection nebulosity that is located 1~kpc from the Sun \\citep{raci68}.  It is illuminated by two Be stars and an associated cluster of lower mass stars \\citep{hoda94}.  This cluster is located within a cavity west of a kidney-shaped molecular cloud which is sharply defined in both CO \\citep{ridg03} and submillimeter emission \\citep{fms01}. A nebulous filament which traces the cavity wall is the dominant feature of the region at both optical and near-IR wavelengths.  The activity of forming young stellar objects has been detected in several places within the molecular cloud. Within the cavity wall is a third massive star, LkH$\\alpha$~234, and its embedded companion \\citep{wein96,cabr97}.  A jet discovered at visible wavelengths \\citep{ray90} with near-IR and mid-IR counterparts \\citep{cabr97} is centered on these two objects.  The jet is blue-shifted and points southwest into the cavity, suggesting that it may be the counterjet to the red-shifted molecular outflow to the northeast.  Another deeply embedded intermediate mass protostellar object, FIRS2, is located three arcminutes to the south of the main cluster \\citep{epc98}.  FIRS2 is located at the region's primary peak in $^{13}$CO emission \\citep{bech78} and is associated with a multipolar molecular outflow, suggesting that it may have one or more young companions \\citep{fuen01,misk01}.  Many other Herbig-Haro objects have been detected around the region which are associated with sites of outflow activity \\citep{es83,hl85}.  Several large-scale outflows have been observed in molecular hydrogen emission \\citep{eisl00}.  {\\it Spitzer} 3~to~8~$\\mu$m and 24~$\\mu$m imaging has revealed a significant extended cluster of embedded Class~I~and~II young stellar objects to the northeast of LkH$\\alpha$~234 \\citep{muze04,gute04}. ",
        "conclusions": "\\subsection{The Initial Configuration and Evolution of Clusters}  As we have shown, characterizing young stellar clusters with azimuthal averaging methods obscures a key component in their analysis.  Those clusters presented that are still highly embedded, GGD~12-15 and IRAS~20050+2720, have a high degree of asymmetry in their stellar distributions compared to NGC~7129. Given the similarity of the asymmetric stellar distributions to the filamentary structures detected in 850~$\\mu$m emission, we argue that the stellar structure observed is a direct result of asymmetry in the distribution of the natal gas. In contrast, the NGC~7129 cluster is approximately circularly symmetric, and occupies a cavity in its natal cloud.  Low average and peak stellar densities within the cluster, overall lack of significant stellar density structure, and a significantly larger cluster core radius are circumstantial evidence that recent expulsion of the bulk of the gas mass from the cavity has allowed the cluster to dynamically expand to its current state.  In the two embedded clusters there is also evidence that the star--forming gas and dust are being disrupted.  While stellar density morphology seems related to dust emission morphology in the embedded clusters, peak stellar densities are often anti--correlated to dust emission maxima, suggesting that either local extinction variation severely affects our sensitivity to stars or stellar feedback mechanisms are already dispersing natal material locally near high stellar densities.  For example, in GGD~12-15, the peak of stellar density is located at a local minimum in 850~$\\mu$m emission, suggesting that the central stars may be opening a cavity in the dense central cloud core.  Similarly, subcluster~B of IRAS~20050+2720 is just outside the region of peak 850~$\\mu$m emission, possibly indicating that it too has dispersed much of its gas, as suggested by \\citet{chen97}.  Subcluster~B may be entering a phase of dynamical expansion similar to what we suggest to explain the current state of the NGC~7129 cluster, while the remaining filament is still actively forming stars in subclusters~A~and~C and the small millimeter emission cores detected nearby \\citep{chin01}.  Analytical investigations and numerical simulations of the dynamics of young clusters following rapid expulsion of their natal gas have shown that dynamical expansion and dissolution can occur on short time scales in young clusters \\citep{am01,kpm99}.  Indeed, over the course of 1~Myr an unbound star moving at 1~km/s can travel 1~pc, suggesting that the high stellar densities observed in young clusters can only last while the system is gravitationally bound by the mass of the natal molecular cloud. The observations presented in this work further these arguments.  Understanding the impact of forming stars on their parent cloud is crucial to understanding the complex dynamics of embedded cluster evolution.  By analyzing stellar density distribution morphology in relation to molecular cloud structure, observational analyses can more adequately address the clear link between star formation, gas expulsion, and the dynamics of the clusters and how these processes guide the evolution of young clusters.  \\subsection{The Impact of Clustering on Disk and Envelope Evolution}  Simulations of massive protostellar disk interactions suggest that truncation and fragmentation of the disks are likely for close approaches with other stars on the order of the disk radius \\citep[e.g.,][]{boff98}.  A recent investigation into the effects of close approaches on classical T~Tauri disks as a function of a variety of parameters suggests tidal truncation and mass loss are possible, as well as enhanced accretion \\citep{pfal05}. Furthermore, the gravitational instability planet formation scenario may require a close approach with another star to perturb the protoplanetary disk to begin the process \\citep{boss02}. If close approaches are likely for a significant fraction of pre-main sequence stars in an embedded cluster, this would have a profound effect on disk evolution and perhaps planet formation in a clustered star--forming environment.  Observations other than those presented here provide further clues that envelope and disk interactions may happen frequently and indeed may be significant in the evolution of disks and the formation of planets.  Several disks around pre-main sequence stars in the Orion Nebula Cluster (ONC) have been directly observed with {\\it HST} \\citep{mo96}; these disks have radii ranging from as small as 50~AU to as large as 1000~AU, and their central stars have derived ages of 0.8-3.0~Myr.  Given the high stellar densities measured in the ONC, as well as the wide range of disk radii and sharp disk edges observed, tidal truncations from close encounters are a likely explanation \\citep{mo96}.  A 10~AU gap in the disk of $\\sim$1~Myr old T~Tauri star CoKu Tau/4 has been discovered in observations from the {\\it Spitzer} Infrared Spectrograph, suggesting the presence of a planet \\citep{forr04,dale04}.  This suggests that giant planet formation may occur very early in a protoplanetary disk's lifetime, an argument in favor of planet formation via gravitational instability. Current orbital models for the recently discovered distant minor planet Sedna \\citep{btr04} argue for a close interaction from a passing star, suggesting that our solar system may have formed in a clustered environment.  One of the strengths of the stellar surface density mapping method adopted for this work is its ability to probe stellar densities over a wide dynamic range with adequate spatial resolution.  While stellar collisions are very unlikely at the stellar densities reported in this work \\citep[cf.][and note that lower average stellar masses further increase the timescales in Figure~1 of that study]{bb02}, close approaches on the order of protostellar envelope or large T~Tauri disk radii \\citep[$R = 10^3$~AU,][]{ma01,mo96} at these densities may be much more common.  The time--scale for a given star to experience a close encounter when among a cluster of stars of stellar density $n$ and velocity dispersion $v_{disp}$ can be roughly approximated via a simple analysis outlined in \\citet{bt87}.  They express this time--scale, $t_{coll}$, as: \\begin{displaymath} \\frac{1}{t_{coll}} = 16\\sqrt{\\pi}nv_{disp}R_*^2(1+\\frac{GM_*}{2v_{disp}^2R_*}) \\end{displaymath} All stars are assumed to have the same mass, $M_*$, and the same interaction radius, $R_*$.  For this analysis, we assume $M_* = 0.5~M_{\\odot}$ as an appropriate median stellar mass from young cluster initial mass functions reported in the literature \\citep[e.g., ][]{muen02,luhm03}, and $v_{disp} = 1~km~s^{-1}$ based on Maxwellian velocity dispersions derived from C$^{18}$O linewidths for these clusters (see Table~\\ref{littable}).  Figure~\\ref{colltime} shows the derived collisional timescales as a function of stellar density for two choices of $R_*$, $10^3$~AU for protostellar envelopes \\citep{ma01} and large T~Tauri disks and $10^2$~AU for classical T~Tauri disks \\citep{mo96}.  The timescales considered must be less than or on the order of $10^5$ to $10^6$~yr to allow for a significant frequency of interactions to occur over the length of a forming star's protostellar collapse phase \\citep{kh95} or the approximate length of a young cluster's active star--forming phase \\citep{ps00}, respectively.  For example, at a density of $10^4$~pc$^{-3}$, protostars and T~Tauri stars with large disks are likely to have a close encounter on the order of their envelope or disk radius on a timescale of $\\sim$$10^5$~yr, while classical T~Tauri disk radius interactions occur much less often, on a timescale of $\\sim$$10^7$~yr.  The latter timescale suggests that $\\sim$10$\\%$ of stars at densities of $10^4$~pc$^{-3}$ would have a close encounter on the order of $10^2$~AU in $\\sim$1~Myr, if the high stellar density environment can indeed last that long.  Utilizing our stellar density maps and cluster membership and field star contamination analyses, we can estimate the prevailing stellar density environment around the stars in the cluster regions relative to the interaction criteria derived above.  First, the associated stellar volume density estimate is determined for each star within the cluster region boundary using the nearest measurement in our density maps.  Note that these are nearest neighbor derived volume density estimates, thus they may be overestimated. We count the number of stars in the cluster regions with densities greater than $10^4$~pc$^{-3}$.  The mean number of field stars falling within the high density areas of the cluster region is estimated directly from the $N_{CM}$ field star models. See Table~\\ref{densitytable} for results for each cluster.  In the two highly embedded clusters presented, GGD~12-15 and IRAS~20050+2720, $72\\%$ and $91\\%$ of the inferred cluster region members respectively are in locations with stellar densities above $10^4$~pc$^{-3}$. This result suggests that there is ample opportunity for protostellar envelope and large T~Tauri disk interactions to occur in these two clusters even over relatively short timescales ($<10^5$~yr).  Furthermore, while classical T~Tauri disk scale interactions are less common, they still may occur in as many as $10\\%$ of the stars in these two clusters over the course of $\\sim$1~Myr. Indeed, 1~Myr is probably an upper limit on the high stellar density phase lifetime of these clusters, as significant gas expulsion is already apparent. Regardless, this estimate is consistent with the suggestion of \\citet{acp03} that up to $30\\%$ of T~Tauri disks must be dispersed within 1~Myr, possibly due to a combination of close encounters with other stars and more gradual evolution and accretion in the smallest disks. In contrast, less than $24\\%$ of the stars in the cluster core boundary of NGC~7129 are at densities greater than $10^4$~pc$^{-3}$.  Clearly, there is little chance for further disk-affecting interactions in this cluster, although this is not entirely unexpected.  \\citet{gute04} report the disk fraction in the cluster region is $54\\%$, suggesting that the population is already somewhat evolved.  Furthermore, spectroscopic evidence suggests that the low mass stars in NGC~7129 are between 1.5-2.0~Myr old \\citep{hill95}.  If we assume cluster asymmetry and high average extinction are signs of extreme youth ($<1$~Myr) in the young clusters presented, then the high density environments we observe in GGD~12-15 and IRAS~20050+2720 provide evidence supporting the idea that protostellar envelope and large ($10^3$~AU) T~Tauri disks are likely to be affected by close encounters with other members in the youngest clusters, and this may lead to significant tidal stripping or conversely, accelerated accretion, early in their evolution.  If the molecular gas mass of the natal cloud is indeed adequate to maintain the high density configurations we have observed over a 1~Myr timescale, a small but measurable fraction of the classical T~Tauri disks have the chance to be significantly truncated or disrupted entirely.  Once gas expulsion and subsequent dynamical expansion occurs, as we argue has happened in NGC~7129, stellar densities decrease to the level at which the disks are allowed to evolve and eventually disperse by means of more gradual activity, such as photoevaporation \\citep{acp03}, residual accretion, or planet formation. Clearly, determining the timescale of bulk gas expulsion is of fundamental importance to understanding the broad implications disk-affecting close encounters may have for disk fraction studies, planet formation frequency and timescale, and the stellar mass spectrum."
    },
    "0410/astro-ph0410616_arXiv.txt": {
        "abstract": "\\grb\\ was a well-localized gamma-ray burst (GRB) detected by the All-Sky Monitor (ASM) on the Rossi X-Ray Timing Explorer (\\rxte) for which no afterglow was identified despite follow-up \\asca\\ and \\ro\\ pointings and optical imaging to limiting magnitude $R > 23$. While an X-ray source, \\axj, was detected just outside the ASM error box, it was never associated with the GRB because it was not clearly fading and because no optical afterglow was ever found.  We recently obtained an upper limit for this source with \\cha\\ that is at least factor of 100 fainter than the \\asca\\ detection.  We also made deep optical observations of the \\axj\\ position, which is blank to limits $V > 25.2$ and $I > 24.0$. In view of these extreme limits we conclude that \\axj\\ was indeed the afterglow of \\grb, which corresponds to an optically ``dark'' GRB. \\axj\\ can therefore be ruled out as the counterpart of the persistent EGRET source \\source. The early light curves from BATSE and the \\rxte\\ ASM show spectral softening between multiple peaks of prompt emission. We propose that \\grb\\ might be a case in which the properties of an X-ray flash (XRF) and a ``normal'' GRB coincide in a single event. ",
        "introduction": "\\label{sec:intro}  One of the intriguing results from five years of GRB localizations by \\beppo\\ is that roughly 60$\\%$ of well-localized GRBs lack an optical transient despite intensive searches \\citep[e.g.,][]{rei01,djo01}. Some of these ``dark'' GRBs could simply be due to a failure to image deeply or quickly enough \\citep{fox03,li03,lamb}. However, in certain cases the optical afterglow may have been missed either because it is obscured by dust in the host galaxy, or because it is located at high redshift ($z\\gax5$).  In the first few months of the ``afterglow era'', which began with the localization of the X-ray afterglow of GRB 970228 \\citep{cos97}, the Burst and Transient Source Experiment (BATSE) detected a GRB that falls in the category of ``dark''. The bright event detected on UT 1997 August 15.50491 and labeled \\grb\\ had a total $\\gamma$-ray fluence $\\approx 5.8 \\times 10^{-5}$ erg cm$^{-2}$, placing it in the top 15$\\%$ of the BATSE fluence distribution. Nearly simultaneous detection by the \\rxte\\ ASM refined the position of \\grb\\ to a small error box \\citep{smi97,smi99}. The localization by the \\rxte\\ ASM was followed several days later by \\asca\\ \\citep{mur97} and \\ro\\ \\citep{gre97} pointings.  While a bright X-ray source \\axj\\ was detected just outside the ASM error box, it was never associated with the GRB because it was not clearly fading and because prompt optical observations failed to reveal an optical transient to limiting magnitude $R > 23$ \\citep{har97}.  In a subsequent review of the evidence we hypothesized nevertheless that \\axj\\ was the afterglow of \\grb, and proposed that this could be tested \\citep{mir03a}. In this paper, we present new \\cha\\ and optical observations of this source, which, together with an analysis of the \\asca\\ and \\ro\\ data, indicate that \\grb\\ was one of the earliest and most luminous ``dark'' bursts in the afterglow era [see \\citet{dep03} for a complete list].  In addition, we discuss the unusual softening over the burst's multiple peaks, which suggests that the intrinsic properties GRB 970815 varied over the duration of the event. Finally, we mention the implications for the counterpart of the steady unidentified EGRET source \\source\\ \\citep{muk02}, whose error ellipse includes the position of \\grb.  \\begin{figure*} \\begin{center} \\epsscale{1.0} \\plottwo{f1a.eps}{f1b.eps} \\caption{ ({\\it a}) \\asca\\ SIS CCD image of the field of GRB 970815 at 3.2--4.4 days after the burst, with the \\rxte\\ ASM error box ({\\it solid line}) and BATSE/{\\it Ulysses} annulus ({\\it dashed lines}) from \\citet{smi02} superposed.  ({\\it b}) \\ro\\ HRI image at 5.5--7.2 days after the burst.  Locations of \\ro\\ HRI point sources are indicated by crosses.  The marginal \\ro\\ source RX~J1608.8+8131 \\citep{gre97} is probably not real (see text). \\label{box} } \\vskip -0.2in \\end{center} \\end{figure*} ",
        "conclusions": "In summary, our verification of the transient nature of \\axj\\ and the lack of an optical counterpart for it compel the conclusion that \\grb\\ was an optically ``dark'' GRB, quite possibly the first one in the afterglow era. Its light curve can be fitted by a power-law decay of index $\\alpha = -1.4$ between the \\asca\\ and \\ro\\ observations, with a spectrum of photon index $\\Gamma = 1.64 \\pm 0.35$. Analysis of the \\rxte\\ ASM observation leads to the conclusion that at least some GRBs exhibit properties that are similar to XRFs after the cessation of ``normal'' gamma-ray activity. Such a detection suggests that variations in the intrinsic properties of the burst might account partly for the observed distribution of XRFs and GRBs.  This finding warrants a fresh examination of archival optical/IR data, as well as follow-up optical, IR, and sub-mm observations to search more deeply for a host galaxy and determine if it is obscured by dust or located at high redshift. Even if ambiguous within the \\ro\\ HRI error circle, identification of the host may be supported by spectroscopic detection of strong Ly$\\alpha$ emission, a common signature of large star formation rates in GRB host galaxies \\citep{fynbo}.  \\medskip"
    },
    "0410/hep-th0410017_arXiv.txt": {
        "abstract": "In cosmic holography, the fundamental quantity is the degrees of freedom on a horizon surface rather than the material contents within the volume.  That is, the horizon area and hence cosmological expansion rate $H$ is related to the entropy.  Using this as a guide, we examine the consequences of an effective dark {\\it entropy} on the past and future dynamics of the universe.  The results, including nonmonotonic expansion behavior, are interesting apart from the theoretical motivation. ",
        "introduction": "\\label{sec.intro}  Acceleration of the cosmic expansion has profound implications for not only the dynamics and contents of the universe, but for fundamental physics on both very high energy physics and low energy physics scales.  Explanations of the acceleration have been sought within new particles and fields, extra dimensions, extensions of gravitation theory, and string theory.  More frequently, purely ad hoc modifications of the Friedmann equations for the expansion dynamics have been proposed.  To avoid purely phenomenological, unmotivated mutations of the Friedmann expansion equation, one might be tempted to concentrate on quantities directly related to the physics. One example is the geometric dark energy proposed by \\cite{lingrav} involving the Ricci scalar curvature; by the Principle of Equivalence one might expect the spacetime curvature to be intimately connected with the acceleration.  Another possibility is the comoving horizon scale $(aH)^{-1}$, a central quantity in early universe inflation.  The time change of this quantity explicitly indicates whether the universe is accelerating or decelerating, and the horizon scale is closely tied to the flatness, structure, and causality issues of inflationary cosmology.  Here we draw our inspiration from holography and consider the horizon area $H^{-2}$ as a key quantity.  Recall that attempts to resuscitate a positive cosmological constant $\\Lambda$ as an explanation for the acceleration run into fine tuning and coincidence problems \\cite{carroll}.  Indeed the scale of any vacuum energy density seems generically much larger than the $(\\sim10^{-3}\\,eV)^4$ required to explain the observations.  One way out of this is if the vacuum energy is not a sum over modes up to some cutoff energy, but rather bounded by spacetime geometry.  Such a constraint could follow from the holographic principle limiting the number of possible quantum states \\cite{thomas}.  We do not here assume that but use that holography views the horizon surface as fundamental and relates it to an entropy, and hence information, bound \\cite{boussormp}.  Now, rather than adding terms to the Hubble expansion $H^2$ corresponding to new energy components in the volume, e.g.\\ quintessence, we add terms to its reciprocal, the horizon area. From holography this denotes adding degrees of freedom, and effectively corresponds to new {\\it entropy\\/} components.  In \\S\\ref{sec.dyn} we investigate the dynamical implications of this ansatz, including the past and future expansion behaviors. \\S\\ref{sec.cos} investigates tests of the scenario through cosmological probes such as distance-redshift and growth factor measurements. \\S\\ref{sec.holo} discusses further aspects of the entropy interpretation and its place within holographic cosmology. ",
        "conclusions": "In searching for new physics to explain the acceleration of the expansion of the universe, one can adopt 1) new, and often not well defined, components of the energy density, 2) new forms of the gravitation equations, or 3) new geometric structures of spacetime.  Here we postulated the last route, exploring some implications of the horizon area and its associated entropy.  The picture appears consistent and a real seed of motivation exists from the principle of holography. While the specific ansatz may not be significantly less ad hoc than quintessence models or modifications of the Friedmann equations of expansion it does lead to a rich variety of expansion histories of the universe. The implications include features of expansion rates and equations of state nonmonotonic with scale factor -- an important broadening of physical possibilities. Furthermore we have shown that the parameter space can be constrained by cosmological observations.  But in addition to fundamental physics connections, one of the most significant virtues of the dark entropy scenario is that it is distinct from dark energy, including the cosmological constant $\\Lambda$. That is, there is no limit of the parameters, no corner of phase space, that approaches the cosmological constant.  This is in stark contrast to almost all viable dark energy potentials or modifications of gravity. This property makes dark entropy, aside from other motivations, an important benchmark model for comparing to observations and understanding the fundamental physics."
    },
    "0410/astro-ph0410170_arXiv.txt": {
        "abstract": "We report the detection of several new circumstellar disks seen in silhouette against background nebular light in the outskirts of the Orion nebula and the neighboring H~{\\sc ii} region M43.  These were detected as part of our H$\\alpha$ survey of Orion with the Advanced Camera for Surveys on-board the {\\it Hubble Space Telescope}.  Several of the disks show bipolar reflection nebulae, microjets, or pronounced temporal variability of their central stars.  The relatively large fraction of bipolar reflection nebulae or microjets in our sample may be a selection effect caused by the faint nebular background far from the Trapezium.  Two disks in our sample are large and particularly noteworthy: A nearly edge-on disk, d216-0939, is located several arcminutes northwest of M43 and resembles the famous HH~30 disk/jet system in Taurus.  It drives the 0.15 pc long bipolar outflow HH 667, and exhibits a remarkable asymmetric reflection nebula caused by the tilt of the flared disk.  With a diameter of $\\sim$2$\\farcs$6 (1200 AU), d216-0939 is as large as the giant edge-on silhouette disk d114-426 in the core of the Orion Nebula.  The large disk d253-1536 is located in a binary system embedded within an externally-ionized giant proplyd in M43.  The disk exhibits distortions which we attribute to tidal interactions with the companion star.  The bipolar jet HH 668 emerges from the proplyd ionization front in a direction orthogonal to the disk, and can be traced to the young star embedded within it.  A bow shock lies 54$\\arcsec$ south of this binary system along the outflow axis.  Proper motions over a 1.4 yr baseline confirm that these emission knots are indeed moving away from d253-1536, with speeds as high as $\\sim$330 km s$^{-1}$ in the HH~668 microjet, and slower motion farther from the star. ",
        "introduction": "The {\\it Hubble Space Telescope} ({\\it HST}) has produced remarkable images of circumstellar disks surrounding young stars.  By far the largest population of known circumstellar disks is located in the Orion nebula at a distance of 460 pc (see Bally et al.\\ 2000), where the disks have typical diameters of 50-1000 AU.  O'Dell, Wen, \\& Hu (1993) found several disks seen in silhouette against the background nebular light, and dozens of young stellar objects surrounded by teardrop- or tadpole-shaped ionization fronts.  These are protoplanetary disks seen purely in silhouette, or externally-ionized photo-ablating protoplanetary disks (``proplyds''), both rendered visible by their location in or near an H~{\\sc ii} region (O'Dell et al.\\ 1993, 1996; McCullough et al.\\ 1995; Bally et al.\\ 1998, 2000). McCaughrean \\& O'Dell (1996) found six additional disks seen purely in silhouette against the bright nebular emission from the Orion nebula. More detailed {\\it HST} imaging studies of the nebular core revealed a total of 15 silhouette disks (Bally, McCaughrean, \\& O'Dell 2000; McCaughrean et al.\\ 1998) and nearly 200 bright proplyds (O'Dell et al.\\ 1993; O'Dell \\& Wong 1996; Bally et al. 1998, 2000).  Some silhouette disks are seen embedded inside bright proplyds (Bally et al.\\ 2000), where in some cases the disks glow in [O~{\\sc i}] and H$_2$ emission (Chen et al.\\ 1998; St\\\"{o}rzer \\& Hollenbach 1998). In addition, several one-sided ``microjets'' are associated with proplyds (Bally et al.\\ 2000; Reipurth \\& Bally 2001), indicating that the embedded stars are still surrounded by active accretion disks.  So far, only the inner portion of the Orion nebula -- a region that contains about 400 of the nearly 2,000 members of the extended Trapezium cluster of low mass stars -- has been completely surveyed with {\\it HST}.  However, O'Dell (2001) presented HST images of several objects outside the nebular core, demonstrating that additional young stars, disks, and jets remain to be discovered in the outer regions of the Orion nebula.  With the new Wide Field Camera of the Advanced Camera for Surveys (ACS/WFC) installed on {\\it HST} during the 2002 servicing mission, we explored the outer Orion Nebula where radiation fields are less intense than in the nebular core.  The outskirts of the Orion nebula contain young stars bathed in radiation fields that may be more typical of the environment in which most stars in the sky are born. The survey samples a large portion of the Orion Nebula Cluster.  Our ACS/WFC survey also included the first comprehensive {\\it HST} imaging of the neighboring H~{\\sc ii} region M43, several acrminutes north of the Trapezium, powered by the single B0-type star NU Ori.  The new disks are seen despite the the fainter background nebular emission far from the Trapezium.  On the other hand, the faint nebular background probably makes it easier to see extended emission from bipolar reflection nebulae and microjets emerging from the silhouette disks --- phenomena which are apparently less common among pure silhouette disks in the inner Orion nebula (Bally et al.\\ 2000). After a general description of these new silhouette disks (\\S 3), we provide more detailed investigations of two particularly interesting disks and their associated jets (\\S4 and \\S 5).  One additional proplyd with a bipolar jet and silhouette disk, d181-826 and HH~540, is also seen in our survey in the outer parts of the Orion nebula, but has been discussed separately (Bally et al.\\ 2004) because of its remarkable morphology. ",
        "conclusions": "Our {\\it HST}/ACS imaging survey of the outer parts of Orion and M43 yielded the discovery of 9 new disks seen in silhouette against faint background light, and they show important new details for one previously discovered silhouette disk.  Compared to other silhouette disks, several members of this new sample are unusual in that they show evidence for bipolar reflection nebulae or microjets.  This may be a selection effect due to the fainter background levels far from the Trapezium. Two of the new disks in our sample are found in or near M43 and deserve special attention:  1.  d216-0939 is a nearly edge-on disk with a remarkable reflection nebula and a bipolar microjet that resembles the HH~30 disk/jet system in Taurus.  The disk is as large as the largest previously-known silhouette disk.  It also drives the 0.15 pc long bipolar jet HH~667.  2.  d253-1536 is also a very large silhouette disk, and is the only silhouette disk known so far to exhibit tidal distortions induced by a companion star in a binary system.  It is located inside the M43 H~{\\sc ii} region, and is associated with a bright proplyd ionization front.  The disk drives both a microjet and a larger 0.14 pc long bipolar jet called HH~668.  Images taken during Cycles 11 and 12 reveal detectable proper motions in the jet, with de-projected speeds as high as 330 km s$^{-1}$.  With an electron density of order 4$\\times$10$^3$ cm$^{-3}$ near the source, this implies a mass-loss rate of roughly 10$^{-9}$ M$_{\\sun}$ yr$^{-1}$."
    },
    "0410/astro-ph0410493_arXiv.txt": {
        "abstract": "We present the results of a \\ion{Mg}{2} $\\lambda\\lambda2796,2803$ absorption-line survey using QSO spectra from the Sloan Digital Sky Survey Early Data Release.  Over 1,300 doublets with rest equivalent widths greater than 0.3\\,\\AA\\ and redshifts $0.366 \\le z \\le 2.269$ were identified and measured.   We find that the $\\lambda2796$ rest equivalent width ($W_0^{\\lambda2796}$) distribution is described  very well by an exponential function $\\partial N/\\partial W_0^{\\lambda2796} = \\frac{N^*}{W^*}\\,e^{-\\frac{W_0}{W^*}}$, with $N^*=1.187 \\pm 0.052$ and $W^*=0.702 \\pm 0.017$\\AA.  Previously reported power law fits drastically over-predict the number of strong lines.  Extrapolating our exponential fit under-predicts the number of $W_0^{\\lambda2796} \\le 0.3$\\,\\AA\\ systems,  indicating a transition in $\\partial N/\\partial W_0^{\\lambda2796}$ near $W_0^{\\lambda2796} \\simeq 0.3$\\,\\AA.  A combination of two exponentials reproduces the observed distribution well, suggesting that \\ion{Mg}{2} absorbers are the superposition of at least two physically distinct populations of absorbing clouds.   We also derive a new redshift parameterization for the number density of $W_0^{\\lambda2796} \\ge 0.3$\\,\\AA\\ lines: $N^*=1.001\\pm0.132\\,(1+z)^{0.226\\pm0.170}$ and $W^*=0.443\\pm0.032\\,(1+z)^{0.634\\pm0.097}$\\,\\AA.  We find that the distribution steepens with decreasing redshift, with $W^*$ decreasing from $0.80\\pm0.04$\\,\\AA\\ at $z = 1.6$ to $0.59\\pm0.02$ \\,\\AA\\ at $z=0.7$.   The incidence of moderately strong ($0.4\\,\\mathrm{\\AA} \\lesssim W_0^{\\lambda2796} \\lesssim 2$\\,\\AA) \\ion{Mg}{2} $\\lambda2796$ lines does not show evidence for evolution with redshift.   However, lines stronger than $\\approx 2$\\,\\AA\\ show a decrease relative to the no-evolution prediction with decreasing redshift for $z \\lesssim 1$. The evolution is stronger for increasingly stronger lines.   Since $W_0$ in saturated absorption lines is an indicator of the velocity spread of the absorbing clouds, we interpret this as an evolution in the kinematic properties of galaxies from moderate to low redshift. Monte Carlo simulations do not reveal any strong systematic  effects or biases in our results.  While more recent SDSS QSO spectra  offer the opportunity to increase the \\ion{Mg}{2} absorber sample by another order of magnitude, systematic errors in line identification and measurement will begin to dominate in the determination of absorber property statistics. ",
        "introduction": "Intervening QSO absorption lines provide an opportunity to study the evolution of  galaxies selected via their gas cross-sections, independent of their stellar luminosities, from the earliest era of galaxy formation up to the present epoch.   The chance alignment of galactic gas with the line of sight to a distant QSO allows the unique opportunity to study the gaseous phase of galaxies, without the biases involved in luminosity limited studies.  In particular, low-ion metal lines and  strong \\ion{H}{1} lines can be used to sample the low-ionization and neutral gas  bound in galactic systems.  Surveys for intervening \\ion{Mg}{2} absorption systems are particularly useful, since the \\ion{Mg}{2} doublet at rest wavelengths $\\lambda\\lambda2796,2803$ allows ground based  surveys to track them down to relatively low redshift $(z=0.11)$.  Strong \\ion{Mg}{2} systems are good tracers of large columns of neutral hydrogen gas (Rao \\& Turnshek 2000)  and are useful probes of the velocity structure of the neutral gas components of galaxies, including damped Ly$\\alpha$ (DLA) systems.  Furthermore, unsaturated low-ionization metal lines associated with strong \\ion{Mg}{2} systems can be used to deduce neutral phase metal abundances (Pettini et al. 1999;  Nestor et al. 2003; Prochaska et al. 2003).  Many surveys have studied the statistical properties of \\ion{Mg}{2} absorption systems.  Lanzetta, Turnshek, \\& Wolfe (1987) presented the first significant survey and provided the benchmark results for the statistics of \\ion{Mg}{2} absorbers at relatively high redshift ($1.25 < z < 2.15$).  They found marginal evidence for evolution in the number density of absorbers, but no significant evidence for evolution of the $\\lambda2796$ rest equivalent width ($W_0^{\\lambda2796}$) distribution, which they fit equally well with an exponential and a power law.  Tytler et al. (1987) and Sargent, Steidel, \\&  Boksenberg (1988) provided data on \\ion{Mg}{2} systems at lower redshift $(0.2 \\lesssim z \\lesssim 1.5)$ and found that the comoving number density of absorbers\\footnote{We use partial derivatives in place of the traditional $dN/dz$ since, as we discuss in \\S3, the $z$ and $W_0^{\\lambda2796}$ dependencies of $N$ are not separable.} ($\\partial N/\\partial z$)  increases with redshift, consistent within the errors with no evolution.  Caulet (1989) compared the Lanzetta, Turnshek, \\& Wolfe results to results at lower redshift, finding more moderately strong systems and fewer weak systems at $\\left<z\\right>=1.6$ as compared to $\\left<z\\right>=0.5$.  The study by Steidel \\& Sargent (1992, hereafter SS92), which contained 107 doublets with rest equivalent widths $W_0^{\\lambda2796} \\ge 0.3$\\,\\AA\\ over the redshift range $0.23 \\le z \\le 2.06$, has served as the standard for the statistics of the distribution of \\ion{Mg}{2} $\\lambda2796$ rest equivalent  widths $\\partial N/\\partial W_0^{\\lambda2796}$,  redshift number density $\\partial N/\\partial z$, and redshift evolution for the past decade.  Their conclusions included the following: \\begin{enumerate} \\item{ The $W_0^{\\lambda2796}$ distribution can be fit equally well with either an  exponential or power law.} \\item{$\\partial N/\\partial z$ for systems with $W_0^{\\lambda2796} \\ge 0.3$\\,\\AA\\ increases with redshift in a manner consistent with  no evolution.} \\item{The redshift number density $\\partial N/\\partial z$ for the strongest lines does show  evidence for evolution, with $\\partial N/\\partial z$ decreasing  from the no evolution prediction with decreasing redshift.} \\item{The mean $W_0^{\\lambda2796}$ increases with redshift.} \\item{There is no evidence for correlation between the doublet ratio $W_0^{\\lambda2796} / W_0^{\\lambda2803}$ and redshift.} \\end{enumerate} Churchill et al. (1999, hereafter CRCV99),  presented a study of 30 ``weak'' ($W_0^{\\lambda2796} \\le 0.3$\\,\\AA) systems,  and determined $\\partial N/\\partial W_0^{\\lambda2796}$ and $\\partial N/\\partial z$ for the weak extreme of the \\ion{Mg}{2} $W_0^{\\lambda2796}$ population.  They favored a power law over an exponential for the parameterization of $\\partial N/\\partial W_0^{\\lambda2796}$.  The Sloan Digital Sky Survey (SDSS; York et al. 2000) provides an opportunity to improve the statistics of \\ion{Mg}{2} and other low-ion absorption systems by increasing the number of measured systems by orders of magnitude.  The statistics and evolution of the absorber number density, which is the product of the space density and the absorption cross-section, and the distribution of line strengths, which correlate with the number of kinematic subsystems and absorber velocity dispersion,  are important factors in understanding the physical nature of \\ion{Mg}{2} systems and their evolution. Furthermore, the strong connection to DLA and sub-DLA systems (Rao \\& Turnshek 2000; Rao, Turnshek, \\& Nestor 2004, in preparation) makes the parameterization of the \\ion{Mg}{2} properties key to the study of large \\ion{H}{1} column density systems.  Finally, only large surveys provide statistically significant numbers of rare, ultra-strong systems for complementary imaging and high-resolution kinematic studies.  In this paper, we present the results from a \\ion{Mg}{2} $\\lambda\\lambda2796,2803$ survey using QSO spectra from the SDSS Early Data Release (EDR; Stoughton et al. 2002).  The results of  our analysis, using only a small fraction of the final SDSS database, is of interest because with the large number of \\ion{Mg}{2} doublets identified (over 1,300 with $W_0^{\\lambda2796} \\ge 0.3$\\,\\AA), systematic errors already compete with Poissonian errors in the measured statistics.  In \\S2 we describe the SDSS EDR data set, our doublet finding algorithm, and tests for systematic biases.  In \\S3 we present the results of our analyses, describe the absorber statistics and  their parameterizations, and consider sources of systematic errors.  In \\S4 we discuss our interpretations.  We present our conclusions in \\S5. ",
        "conclusions": "We have identified over 1,300 \\ion{Mg}{2} absorption systems with $W_0^{\\lambda2796} \\ge 0.3$\\,\\AA\\ over the redshift range $0.367 \\le z \\le 2.269$ in the SDSS EDR QSO spectra.  The size of the sample is such that statistical errors are comparable to systematic effects and biases.  We used simulations to improve our line-finding algorithm until the systematics could no longer be isolated from the noise at a level where we could further improve our parameterizations.   Using the combined redshift and $W_0^{\\lambda2796}$ data, we offered a new redshift parameterization for the distribution of systems.  In conclusion, we have shown: \\begin{enumerate} \\item{The rest equivalent width distribution for intervening \\ion{Mg}{2} $\\lambda2796$ absorption lines detected in the SDSS EDR QSO spectra with $W_0^{\\lambda2796} \\ge 0.3$\\,\\AA\\  is very well described by an exponential, with $N^*=1.187 \\pm 0.052$ and $W^* = 0.702 \\pm 0.017$\\,\\AA.  Power law parameterizations drastically over-predict the number of strong lines, and our exponential under-predicts previously reported values for the number density of weaker lines.} \\item{When compared with the number density of  $W_0^{\\lambda2796} < 0.3$\\,\\AA\\ lines from other studies, our results show that neither an exponential nor a power law accurately represents the full range of $W_0^{\\lambda2796}$.  Simulations of our catalog support this finding.  We propose a combination of two exponential distributions, where the weak lines are described by  the parameters $N^*_{wk} \\approx 1.7$ and $W^*_{wk} \\approx 0.1$\\,\\AA\\ and the moderate and strong lines by the parameters $N^*_{str} \\approx 0.9$ and $W^*_{str} \\approx 0.8$\\,\\AA.} \\item{The rest equivalent width distribution steepens with decreasing redshift, with $W^*$ decreasing from $0.80\\pm0.04$\\,\\AA\\ at $z = 1.6$ to $0.59\\pm0.02$\\,\\AA\\ at $z=0.7$.} \\item{For lines with $W_0^{\\lambda2796} \\lesssim 2$\\,\\AA, there is no significant evolution detected in $\\partial N/\\partial z$.} \\item{For lines with $W_0^{\\lambda2796} \\gtrsim 2$\\,\\AA, evolution is detected in $\\partial N/\\partial z$, with a decrease from the no-evolution prediction of $\\approx 45\\%$ from $z=1.8$ to 0.6.  The evolution is stronger for stronger lines and redshifts $z \\lesssim 1$.} \\item{The number density of \\ion{Mg}{2} absorption lines with $W_0^{\\lambda2796} \\ge 0.3$\\,\\AA\\ is well parameterized by $\\frac{\\partial ^2N}{\\partial z\\,\\partial W_0}=\\frac{N^*(z)}{W^*(z)}\\,e^{-\\frac{W_0}{W^*(z)}}$, with $N^*(z)=1.001\\pm0.132\\,(1+z)^{0.226\\pm0.170}$ and $W^*(z)=0.443\\pm0.032\\,(1+z)^{0.634\\pm0.097}$\\,\\AA.} \\end{enumerate} Several lines of evidence suggest that the clouds giving rise to \\ion{Mg}{2} absorption comprise multiple physically-distinct populations.  The results from high resolution work and the apparent transition in $\\partial N/\\partial W_0^{\\lambda2796}$ are perhaps the strongest evidences, but the apparent contribution of enriched Ly$\\alpha$ forest lines, the inability of simple physical models to reproduce in full the kinematic data and the menagerie of galaxy types, luminosities, environments and impact parameters that contribute to the absorber galaxy population are also consistent with this picture.  The situation will be made even clearer with the analysis of the full SDSS database.  With over an order of magnitude more data available, there will be enough high signal-to-noise  ratio spectra for studies to reach lines weaker than 0.3\\,\\AA\\ providing coverage across the $\\partial N/\\partial W_0^{\\lambda2796}$  transition in a single survey.  Although the large number of systems  will require large, improved Monte Carlo simulations, the improved statistics will allow for finer analysis of the strength-dependent redshift evolution. Additionally, work currently  in progress will extend our analysis down to redshifts $z\\approx0.2$ and $W_0^{\\lambda2796} \\gtrsim0.1$ (Nestor et al., in preparation)."
    },
    "0410/astro-ph0410346_arXiv.txt": {
        "abstract": "{Two years ago, the OGLE-III survey (Optical Gravitational Lensing Experiment) announced the detection of 54 short period multi-transiting objects in the Galactic bulge (Udalski et al., \\cite{udalski1}, \\cite{udalski2}). Some of these objects were considered to be potential hot Jupiters. In order to determine the true nature of these objects and to characterize their actual mass, we conducted a radial velocity follow-up of 18 of the smallest transiting candidates. We describe here our procedure and report the characterization of 8~low mass star transiting companions, 2~grazing eclipsing binaries, 2~triple systems, 1~confirmed exoplanet (OGLE-TR-56b), 1~possible exoplanet (OGLE-TR-10b), 1~clear false positive and 3~unsolved cases. The variety of cases encountered in our follow-up covers a large part of the possible scenarii occuring in the search for planetary transits. As a by-product, our program yields precise masses and radii of low mass stars. ",
        "introduction": "Since 1995 the search for planets by radial velocity surveys has lead to the detection of more than 120 planetary candidates. The diversity of orbital characteristics, the mass distribution of the planets (actually only the $m\\sin i$) and its link with brown dwarfs and low-mass stars as well as the characteristics of host stars prompted a reexamination of planetary formation theory (e.g., Udry et al. \\cite{udry}, Santos et al. \\cite{santos03}, Eggenberger et al. \\cite{eggen}). The most unexpected fact was the existence of extrasolar giant planets (EGPs) in very short orbits. Extra mechanisms, not envisioned by the study of our Solar system, have been suggested to explain these objects, like the migration of planets in the proto-planetary disk and gravitational interactions (e.g., Goldreich \\& Tremaine \\cite{goldreich}, Lin et al. \\cite{lin}).  Monitoring of photometric transits caused by an EGP passing on the disk of its hosting star and obscuring part of its surface provides an opportunity to determine its actual size. When combined with spectroscopic observations, it leads to the unambiguous characterization of the two fundamental parameters (mass and radius) used for internal structure studies of EGPs. The discovery of HD209458 by both Doppler measurements (Mazeh et al. \\cite{mazeh}) and photometric transit (Charbonneau et al. \\cite{charbonneau}; Henry et al. \\cite{henry}) led to the first complete characterization of an EGP, illustrating the real complementarity of the two methods. These last years many extensive ground-based photometric programs have been initiated to detect transits by short period EGPs (Horne \\cite{horne}). The OGLE-III survey (Optical Gravitational Lensing Experiment) announced recently the detection of 137 short-period multi-transiting objects (Udalski et al., \\cite{udalski1}, \\cite{udalski2}, \\cite{udalski3}, \\cite{udalski4}). The estimated radius of these objects range from 0.5 Jupiter radius to 0.5 solar radius and their orbital periods range from 0.8 to 8 days. The smallest objects could be suspected to be EGPs, but considering only the radius measured by OGLE, one can not conclude on the planetary nature of the objects per se. They could as well be brown-dwarf or low-mass stars since in the low mass regime the radius is independent of the mass (Guillot, \\cite{guillot}). No information on the mass of these companions is given by the transit measurements. Doppler follow-up of these candidates is the only way to confirm the planetary, brown-dwarf or low-mass star nature of the companions. Planetary transit detection suffers also some ambiguity related to the configuration of the system. The radial velocity measurement is therefore very important in that respect to discriminate true central transits from other cases due to, for example, grazing eclipsing binaries, blended system and stellar activity. The spectroscopy of the central star, which is a by-product of radial velocity measurement, is necessary to constrain the star radius and therefore the companion radius. The measurement of the true mass of a companion by the radial velocity orbit, coupled with the measurement of its radius, lead to a direct measurement of its mean density, an essential parameter for the study of internal structure of EGPs, brown dwarfs and low mass stars.  The difficulties of Doppler follow-up of OGLE candidates come from the faintness of the stars (with V magnitudes in the range 14-18). The fields, which are located in the galactic disk, are furthermore very crowed. To characterize a hot Jupiter, one needs radial velocity precision better than 100 {\\ms} and the capability to distinguish whether the system is blended or not by a third star like the triple system HD41004 (Santos et al., \\cite{santos03}).  Several teams are involved in the Doppler follow-up of OGLE candidates. Konacki et al. (\\cite{konacki1}) announced first that the companion of OGLE-TR-56 is a planet of 0.9 Jupiter mass. Dreizler et al. (\\cite{dreizler2}) gave an upper limit of 2.5 Jupiter mass for the companion of OGLE-TR-3. This object was however refuted by Konacki et al. (\\cite{konacki2}) who also gave informations on 3 other OGLE companions (OGLE-TR-10, 33 and 58). Additional measurements, conducted by Torres et al. (\\cite{torres}), lead to improve the mass determination of OGLE-TR-56b to 1.45 Jupiter mass. Recently, we announced the characterization of planets OGLE-TR-113b, OGLE-TR-132b (Bouchy et al. \\cite{bouchy}, Moutou et al. \\cite{moutou}) and OGLE-TR-111b (Pont et al. \\cite{pont}).  We present in this paper the Doppler follow-up observations of 18 OGLE multi-transiting companions (OGLE-TR-5, 6, 7, 8, 10, 12, 17, 18, 19, 33, 34, 35, 48, 49, 55, 56, 58 and 59) from the 54 detected in the Galactic bulge (Udalski et al., \\cite{udalski1}, \\cite{udalski2}). ",
        "conclusions": "\\label{discu}    \\begin{table*} \\caption{Summary table with the physical parameters $r$, $R$, $m$ and $M$, the orbital angle $i$ and the identification of the system.} \\label{tablezoo} \\begin{tabular}{r r r r r r l}\\hline \\hline Name& r & R & m & M & i & comments \\\\ & [R$_\\odot$] & [R$_\\odot$] & [M$_\\odot$] & [M$_\\odot$] & \\\\ \\hline 5 & 0.263 $\\pm$ 0.012    & 1.39 $\\pm$ 0.026 & 0.271 $\\pm$ 0.035 & 1.12 $\\pm$ 0.21 & 81-90 & G-M binary \\\\ 6 & 0.393 $\\pm$ 0.018    & 1.99 $\\pm$ 0.068 & 0.359 $\\pm$ 0.025 & 1.37 $\\pm$ 0.14 & 88-90 & F-M binary \\\\ 7 & 0.282 $\\pm$ 0.013    & 1.64 $\\pm$ 0.052 & 0.281 $\\pm$ 0.029 & 1.32 $\\pm$ 0.21 & 86-90 & F-M binary \\\\ 8 & 1.217 $\\pm$ 0.035    & 1.27 $\\pm$ 0.035 & 1.160 $\\pm$ 0.010 & 1.20 $\\pm$ 0.010 & 86 & G-G SB2 \\\\ 10 & 0.156 $\\pm$ 0.012   & 1.21 $\\pm$ 0.066 & 0.00063 $\\pm$ 0.00020 & 1.22 $\\pm$ 0.045 & 87-90  & possible planet\\\\ 12 & 0.440 [0.380-1.089] & 1.47 $\\pm$ 0.065 & 0.492 $\\pm$ 0.026 & 1.28 $\\pm$ 0.10 & 85 & F-M binary\\\\ 17 & 0.371 [0.325-0.956] & 1.52 $\\pm$ 0.110 & 0.482 $\\pm$ 0.040 & 1.06 $\\pm$ 0.13 &  86-88 & F-M binary\\\\ 18 & 0.390 $\\pm$ 0.040   & 1.95 $\\pm$ 0.042 & 0.387 $\\pm$ 0.049 & 1.18 $\\pm$ 0.22 & 79-86 & F-M binary \\\\ 19 &  -  &  -  &  -  &  -  & -  & unsolved case \\\\ 33 &  -  &  -  &  -  &  -  & -  & blend \\\\ 34 & 0.435 $\\pm$ 0.033   & 2.31 $\\pm$ 0.174 & 0.509 $\\pm$ 0.038 & 1.56 $\\pm$ 0.17 & 86-90 & F-M binary \\\\ 35 & 1.74 $\\pm$ 0.039    & 1.71 $\\pm$ 0.054 & 1.200 $\\pm$ 0.009 & 1.19 $\\pm$ 0.009 & 80-81 & F-F SB2 \\\\ 48 &  -  &  -  & - &  -  &  -  & unsolved case \\\\ 49 &  -  &  -  & - & -   &  -  & unsolved case \\\\ 55 & 0.209 [0.204-0.989] &  1.92 $\\pm$ 0.036   & 0.276 $\\pm$ 0.038 & 1.36 $\\pm$ 0.28 &  81-90 & F-M binary\\\\ 56 & 0.128 $\\pm$ 0.009   & 1.12 $\\pm$ 0.069 & 0.00113 $\\pm$ 0.00013 & 1.10 $\\pm$ 0.078 & 81-83  & confirmed planet \\\\ 58 &  -  & -   &  -  &  -  & -   & false positive \\\\ 59 &   - &  -  &  -  &  -  &  -  & blend \\\\ \\hline \\end{tabular} \\end{table*}  Table~\\ref{tablezoo} summarizes all the available information derived for our objects following our procedure described in section~\\ref{synth}. The variety of cases encountered in our sample of eighteen objects is striking and covers a large part of the bestiary of possible contaminations in the search for planetary transits. Target OGLE-TR-5, 6, 7, 12, 17, 18, 34, and 55 have clear resolved orbits of eclipsing binaries with a large F/G primary and a small M transiting companion. OGLE-TR-8 and 35 have clear resolved orbits of equal-mass, grazing eclipsing binaries. OGLE-TR-33 and 59 have resolved orbits of eclipsing binaries in a hierarchical triple system. OGLE-TR-56 shows small radial velocity variations in agreement with Torres et al. (\\cite{torres}) which confirm the planetary nature of the transiting companion. OGLE-TR-10 shows small radial velocity variations which could be due to planetary companion. OGLE-TR-19, 48 and 49 are unfortunately not yet solved but we strongly suspect false positive transit detections. OGLE-TR-58 shows no radial velocity variations and its light curve presents clear indication that the detected photometric signal is not a bona fide transit. Note that in some cases (OGLE-TR-12, 17 and 59), the initial period identified from the light curve was not correct and that in many cases our secondary radius are very different from the initial value of Udalski et al. (\\cite{udalski1}, \\cite{udalski2}).  Our study has yielded precise radii and masses for a certain number of low mass star companions. The mass-radius relation for these objects is given in the Fig.~\\ref{MRRwide}. We note that no brown-dwarves were detected in our sample in agreement with the so-called brown dwarf desert for the short period companions. No stellar companions were detected in the mass domain 0.6-1.0~M$_\\odot$ due to the fact that we selected in priority the smallest candidates of the OGLE survey. OGLE-TR12, 17 and 55 (white points) needs additional photometric measurements in order to properly constrain the impact parameter $b$ of the transits. For the other 5 low mass star transiting companions, the precision on the radius and mass determination are respectively on the range 4.5-10\\% and 7-13\\%. This precision is not at the level needed to provide a crucial test of stellar physics (e.g., Andersen~\\cite{andersen}). However the empirical Mass-Radius relation remains poorly constrained due to the lack of observations of M-type eclipsing binaries. These 5 new candidates significantly increase the number of known M-type eclipsing binaries and give new observational constraints on models. For comparison we added on Fig.~\\ref{MRRwide} the three known M-type eclisping binaries (Metcalfe et al.~\\cite{metcalfe}, Torres \\& Ribas~\\cite{torres02}, Ribas~\\cite{ribas}). Although characterized with a significatively lower accuracy, our 5 low mass stars seems to follow the same departure from models.  The OGLE fields are very crowded, and some of the targets are expected to be contaminated by background stars (or foreground fainter stars). Extrapolating the density of bright stars to fainter magnitudes indicates that there may be on average 1.4 contaminants stars per object down to $\\Delta mag =6$. Bright contaminant would be detected by the spectroscopy, but faint contaminants can go unnoticed and contribute a few percent to the light curve. This makes the photometric transit depth shallower, leading to an underestimation of $\\overline{r}$. Note that this effect is seeing dependent.  The effect of different assumptions on limb darkening on the derivation of the parameters from the light curve were verified using OGLE-TR-6 (central transit) and OGLE-TR-55 (grazing eclipse). Changing u1+u2 by 0.2 leads to a difference of the order of 2\\% on $\\overline{r}$ and on $V_T/R$, while removing the limb darkening entirely changes $\\overline{r}$ and $V_T/R$ by $\\sim$ 8\\%. We also test the effect of changing the limb darkening law in the derivation of the rotation velocity from the CCF. Using coefficients  $u1\\! + \\! u2\\! =\\! 0.3$ instead of $u1\\! + \\! u2\\! =\\! 0.6$ modifies $V_{rot} \\sin i$ by about 3\\%.     \\begin{figure*} \\resizebox{15cm}{!}{\\includegraphics{bouchy_f11.eps}} \\caption{The mass-radius relation for all stellar objects in our sample, primary and secondary. Open circles correspond to candidates OGLE-TR-12, 17 and 55 with a wide range of possible radius. Triangles correspond to the three known M-type eclipsing binaries. The dashed curve up to 0.6 M$_\\odot$ corresponds to the 1 Gyr theoretical isochrones of Baraffe et al. \\cite{baraffe}. Above 0.6 M$_\\odot$ the 3 dashed curves show the Padova model mass-radius relations for solar metallicity for three age values.} \\label{MRRwide} \\end{figure*}  This program illustrates the capability of ground-based spectrographs like FLAMES, UVES and HARPS to follow the faint transiting candidates found by photometric surveys. It demonstrates the usefulness of such a Doppler follow-up in order to discriminate among a large sample of possible contaminations in the search for planetary systems. We used in average only 2.5 hours of observing time per object, thanks essentially to the very high efficiency of the FLAMES multi-fibers facility. Our analysis shows that a large part of the transiting candidates could be rejected even a priori through a fine-tuned light curve analysis (confidence factor, sinusoidal variations and transit shape). It is also clear that more accurate measurements in both photometry and radial velocity will be very useful in order to provide stronger constraints on the mass and radius of transiting companions, especially for the suspected planetary system OGLE-TR-10, for the unsolved cases OGLE-TR-19 and 49 and for the unconstrained radius of the companions of OGLE-TR-12, 17 and 55."
    },
    "0410/astro-ph0410036_arXiv.txt": {
        "abstract": "We present initial results of WIYN spectroscopic observations of selected objects detected in the \\emph{Chandra} Multiwavelength Plane (ChaMPlane) Survey in fields towards the Galactic anti-center.  ChaMPlane is designed to identify low luminosity X-ray sources, both accretion-powered and stellar coronal, in the Galaxy.  It also includes a wider-fields optical imaging Survey conducted with the NOAO Mosaic cameras to identify optical counterparts as well as \\ha-selected objects in the $\\sim5\\times$ larger field. We report spectra for 1048 objects in galactic anti-center (i.e.\\ northern) fields, resulting in 609 type determinations.  These include 5 new cataclysmic variables, 4 Be stars, 14 lithium-absorption stars, 180 stellar coronal sources (primarily dMe stars), and 29 new quasars.  Bright optical counterparts of \\emph{Chandra} sources in this sample are most frequently dMe stars whereas a majority of the faintest (R $\\sim21$) spectroscopically classified \\emph{Chandra} source counterparts are quasars.  The bulk of \\ha-selected sources appears to be roughly evenly divided between dMe stars and M stars at all magnitudes. ",
        "introduction": "The ChaMPlane Survey \\citep{gri2003,gri2004} is designed to detect and optically identify low luminosity X-ray point sources in the Galactic plane. ChaMPlane is conducted in two phases: the ChaMPlane X-ray survey of serendipitous sources in \\emph{Chandra} X-ray Observatory archival deep ($>$20 ksec) Galactic plane ($|b|<12^\\circ$) fields, and the ChaMPlane Optical Survey of deep $V, R, I,$ and $H\\alpha$ images of fields taken at CTIO or KPNO with the 4m telescope Mosaic $(36'\\times36')$ cameras centered on the \\emph{Chandra} fields.  Detailed descriptions of the X-ray detections and database are given by \\citet{hon2004}, and the optical imaging, astrometry, and photometry with Mosaic are described by \\citet{zha2004}.  The ChaMPlane X-ray Survey takes advantage of the unparalleled angular resolution and $\\sim1''$ absolute astrometry of the \\emph{Chandra} X-ray Observatory to match the serendipitous X-ray sources in the field with faint ($R < 24$) optical counterparts. In all but the most crowded fields, nearly all matches are unambiguous.  The point-source sensitivity for a typical 30 ksec ChaMPlane exposure allows detection of sources with $L_X \\gtrsim 10^{31}~{\\rm erg}\\,{\\rm s}^{-1}$ at 8 kpc, which enables detection of cataclysmic variables (CVs), one of the prime ChaMPlane objectives, throughout a significant fraction of the Galaxy.  \\citet{gri2003} and \\citet{zha2003} have reported early results from the ChaMPlane X-ray Survey and followup optical photometry.  The ChaMPlane Optical Survey provides optical photometric followup of the \\emph{Chandra} serendipitous sources, providing \\har, \\ri, and \\vi\\ colors for the X-ray sources.  By definition, stars without \\ha\\ emission should have an average \\har\\ color index of $\\sim$0.0, while CVs and other accretion-powered systems frequently have significant \\ha\\ excesses, which register as negative \\har\\ colors.  This enables us to select objects that are \\ha\\ and/or X-ray bright for followup spectroscopy and classification.  The ChaMPlane Optical Survey will include all \\emph{Chandra} fields (observed with either the ACIS-I or ACIS-S cameras, with $16' \\times 16'$ and $8' \\times 8'$ fields, respectively) in the ChaMPlane X-ray Survey. The highest priority for the spectroscopic followup is to obtain spectra on each X-ray source optical counterpart with possible \\ha-emission; the second priority is to obtain spectra on {\\it all} \\emph{Chandra} source optical counterpart candidates, regardless of \\ha-emission status.  Since the Mosaic camera has a field of view about 4.5 times larger in area than the ACIS-I, the Optical Survey will also detect many \\ha-emission sources with unknown X-ray emission status.  Thus, as a third priority, spectroscopy is also performed on these objects, since Hydra can simultaneously observe 60--70 objects, which (for the anti-center fields at least) enables us to observe all \\ha-emission candidates in the ChaMPlane Optical Survey field in just one or two Hydra setups.  The ChaMPlane Optical imaging program is described in detail by \\citet{zha2004}. Normally, three (for $V$ and $I$) or five (for $R$ and \\ha) dithered exposures are taken and combined to remove cosmic rays and other image contamination.  Total exposure times are set to enable 5--7\\% precision photometry to 24th magnitude in all bands.  Short exposures in each band are also taken to give unsaturated data for the brighter objects.  Mosaic data reduction and analysis is done using standard IRAF\\footnote{IRAF in distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.} \\emph{mscred} tasks, and photometry with DAOPHOT.  In this paper, we report initial results for followup spectroscopy carried out with the WIYN\\footnote{The WIYN Observatory is a joint facility of the University of Wisconsin, Indiana University, Yale University, and the National Optical Astronomy Observatories.} Hydra multi-object spectrometer. This paper is organized as follows: \\S2 describes the observing and analysis precedures, with spectroscopic results presented in \\S3.  We summarize our results in \\S4. In \\citet{rog2004}, we discuss some of the implications of our anti-center field results for the space density of CVs, using a Monte-Carlo model of the CV distribution within the Galaxy. \\citet{lay2004} discuss the combined X-ray and optical (both photometric and spectroscopic) constraints on the total \\emph{Chandra} source populations in the anti-center fields, as well as X-ray properties of individual classes of objects. ",
        "conclusions": "On examination of the overall data set, some trends are evident in the object classifications. All of our five CV candidates are located in the 80\\% of each field imaged only with the Mosaic camera, and thus only show up on the \\ha\\ lists.  The two known LMXBs showed up, as expected, on both the \\ha\\ and X-ray target lists. The Be stars discovered so far are all at the very bright end of our target objects, owing to the concentration of star forming activity within the solar circle.  The five new CVs and two known LMXBs are distributed across the surveyed magnitude range. Most quasars are optically faint X-ray objects, as would be expected from the strong increase of quasar counts with redshift. Coronal emission stars make up the bulk of the brighter X-ray targets, with some nearby normal (non-emission) stars at the very brightest end of our magnitude range.  Both \\ha\\ lists contain primarily late-type stars, with about $45-50\\%$ (in the \\ha-definite list, with \\har$<-0.3$) and about $30-45\\%$ (in the \\ha-marginal list, with $-0.3<$\\har$<-0.2$) appearing as coronal emission objects regardless of magnitude. The predominance of late-type stars in our \\ha-selected sample is not surprising, given the locations of the strong TiO absorption bands in the continuum of the M-star spectrum at about \\ha\\ causing them to have an \\har\\ color closer to our threshold value than earlier-type stars.  In summary, during the first three years of followup spectroscopy of northern fields for the ChaMPlane project, we have found five new CV candidates, four new Be star candidates, 14 lithium-absorption stars, 29 new quasars, and many new coronal emission objects.  While at this time the total number of CVs detected is small, we have shown that our program of initial Mosaic photometry and followup WIYN-Hydra spectroscopy detects CVs in a systematic fashion.  The larger samples expected from the full ChaMPlane survey will allow a a volume-limited sample determination of the CV space density.  Our serendipitous discoveries of Li stars, AGN, and substantial numbers of coronal emission stars will all lend themselves to extensive further investigation."
    },
    "0410/astro-ph0410200_arXiv.txt": {
        "abstract": "Outflows in the form of fast winds from the accretion disk appear in nearly all high accretion rate cataclysmic variables (CVs), both novalikes (NLs) and dwarf novae in outburst (DN). The primary signatures of CV winds are broad, blueshifted absorption features and P~Cygni profiles that appear in the ionized metal lines of lower inclination systems and the broad emission lines that appear in eclipsing systems. Much progress has been made in our understanding of the structure and behavior of CV outflows based on data obtained with X-ray, EUV, and UV telescopes and the kinematical models developed to fit the observations.  Our current picture of CV outflows is that they are bipolar, rotating winds driven off the accretion disk primarily by radiative line driving from disk, boundary layer, and white dwarf radiation.  Other mechanical forces may play a role in the formation of CV winds. The winds are highly turbulent and show complex vertical and azimuthal structure. Future work will focus on determining the physical, thermal and ionization structure of the winds and probing the mechanisms that govern wind formation and behavior. ",
        "introduction": "\\subsection{The First Observations of CV Winds}  The study of outflows in cataclysmic variables (CVs) was born in the IUE era.  IUE ultraviolet spectra of luminous CVs --- novalikes (NLs) and dwarf novae (DN) in outburst --- showed high ionization UV transitions, particularly of CIV $\\lambda\\lambda$1548,1552~\\AA, SiIV $\\lambda\\lambda$1393,1402~\\AA\\ and NV $\\lambda\\lambda$1238,1242~\\AA, appearing with very broad, blueshifted absorption components and P Cygni profiles, shifting to pure emission in high orbital inclination systems \\citep{heap1978,krautter1981,klare1982,greenstein1982,szkody1982,cordova1982,hassall1983}. Examples of wind-dominated FUV spectra of CVs at different inclinations are given in Figure~1, which shows observations taken from the Far Ultraviolet Spectroscopic Explorer (FUSE) data archive. It was immediately noted that the line profiles were similar to those seen in early-type and luminous stars such as O stars and Wolf-Rayet stars, in which the lines were (and are) believed to originate in expanding winds driven by radiation pressure \\citep{conti1978}. As a result, a wind origin for the UV lines of high accretion rate CVs was adopted early on.  \\begin{figure}[!ht] \\plotone{froning_f1.ps} \\caption{Examples of CV FUV spectra for different system inclination angles.  The spectra are time-averaged plots of observations in the FUSE data archive binned to 0.1~\\AA. The upper left panel shows the spectrum of the DN SS~Cyg in outburst ($i = 30^{\\circ}$).  The upper right panel shows the spectrum of the NL IX~Vel ($i = 60^{\\circ}$). The lower left panel shows the spectrum of the NL, UX~UMa ($i = 71^{\\circ}$), and the lower right panel shows the spectrum of the DN, WZ~Sge, at outburst peak ($i = 75^{\\circ}$). Analysis of the UX~UMa and WZ~Sge observations can be found in Froning et al.\\ 2003 and Long et al.\\ 2003, respectively.} \\end{figure}  IUE spectroscopy revealed several properties of the CV winds. The maximum blue edge velocities of the absorption troughs ($\\simeq5000$~km~s$^{-1}$), which give a lower limit to the terminal velocity of the wind, were of order the escape velocity from the white dwarf (WD) accretor, suggesting that the outflows originate close to the WD \\citep{cordova1982}. The deepest absorption in the lines occured not at the blue edges of the lines, as is the case for early-type stars, but at lower velocities ($\\simeq2000$~km~s$^{-1}$), suggesting that the wind is more slowly accelerated in CVs than in luminous stars \\citep{mauche1987}. In eclipsing systems, the lines, unlike the continuum, were not occulted, indicating that the line-emitting region was extended relative to the bulk of the disk, with sizes of order the size of the occulting donor star \\citep{holm1982,mauche1994}.  A comparison of the strengths of the lines with model lines from spherically-symmetric, radiation pressure driven winds (designed for luminous stars) gave mass-loss rates in the wind from $10^{-11}$ -- $10^{-12} M_{\\odot}$~yr$^{-1}$, or $10^{-2}$ -- $10^{-4}$ of the mass accretion rate in the disk \\citep{krautter1981,klare1982,greenstein1982,cordova1982}.  \\subsection{Early CV Wind Models}  From the start, models for CV winds have been based on the wind models for luminous stars:  resonant line scattering from absorption of photons from the central source transfers momentum to ions; as the wind expands outwards, the local resonance is redshifted, perpetuating the line scattering and the driving of the outflow. (For a more detailed discussion of the physics of CV wind models, see Proga, this volume.) The first models assumed a radial outflow with a constant ionization structure emerging from the center of the disk and the WD \\citep{drew1985,drew1987,mauche1987}. The models were used to compare theoretical line profiles to the observed lines, in particular CIV, the strongest wind feature in the UV spectrum.  The results revealed problems with the picture of a radial wind. In high inclination CVs, the blueshifted absorption component of the wind lines disappears as the disk becomes more edge-on, but the radial models continue to show absorption at high inclination. \\citet{drew1987} cited this as evidence for bipolarity, rather than spherical symmetry, in the outflow, although it was noted by \\citet{mauche1987} that radial winds could still be present if the bipolarity was introduced by limb-darkening effects in the accretion disk. Mauche \\& Raymond pointed out another significant problem with radial CV winds, however: radial winds ``see'' so much of the continuum from the WD and boundary layer (BL) that the wind is ionized beyond CIV unless it is highly optically thick, which requires wind mass loss rates of $\\geq$1/3 of the mass accretion rate. Such a massive wind cannot be radiatively driven, produces UV line fluxes an order of magnitude larger than observed values, and would absorb the soft X-ray continuum from the BL, which is inconsistent with observed X-ray fluxes. On the basis of their models, the authors concluded that radial outflows were unlikely in CVs.  As a result, the next generation of CV wind models adopted biconical, axisymmetric wind structures, with the winds being launched from the accretion disk. \\citet{shlosman1993} developed a kinematical model of a rotating, biconical disk wind, taking into account radiation from the WD, the BL, and the accretion disk when calculating the local ionization structure of the (constant temperature) wind. Radiation transfer and line driving of the wind were treated using the standard wind modeling theory for luminous stars. By comparing theoretical line profiles from their models with those of radial winds, they showed that biconical disk winds provide a better match to the profiles of CIV wind lines at varying inclinations and do not suffer from the excessive absorption of BL photons and subsequent over-ionization of the wind as seen in radial wind models.   \\citet{vitello1993} compared the biconical disk model line profiles to observed IUE spectra of CVs. They were able to match the CIV profiles of the low inclination system, RW~Sex, and the eclipsing NL, RW~Tri, with wind mass-loss rates of order 1 -- 15\\% of the mass accretion rate. Their models gave a typical scale for the CIV line of 50 -- 100~R$_{WD}$ in vertical height above the accretion disk. At this point, the theoretical and observational evidence both pointed to the presence of biconical accretion disk winds driven by radiation pressure in high accretion rate CVs. ",
        "conclusions": "Outflows, in the form of moderately collimated winds, are closely tied to disk accretion in luminous CVs.  A combination of observations and modeling have shown that CV winds are bipolar, highly turbulent, and structurally complex.  Radiative line-driving models for the source of the wind explain many of the observations, but additional, mechanical mechanisms may play a role in the structure and formation of the wind. Future observations and modeling will concentrate on determining the vertical and azimuthal distribution of the wind, its thermal and ionization structure, and the mechanisms that influence wind formation and behavior."
    },
    "0410/hep-ph0410390_arXiv.txt": {
        "abstract": "It is shown that newly performed experimental studies by the Belle Collaboration are excellently incorporated in the unified picture for hadron spectra developed early. From our analysis it follows that the measured values for the masses of the BELLE states exactly coincide with the calculated masses of the states living in the corresponding KK towers. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410292_arXiv.txt": {
        "abstract": "We present time-resolved, ultraviolet (UV) spectroscopy of the SW Sex star DW~UMa in the high state. We confirm that, shortward of 1500~\\AA, the high-state, UV continuum level is lower than the white dwarf (WD)-dominated low-state level. We also do not see the WD contact phases in the high state eclipse light curves. These results confirm our earlier finding that the WD in this system is hidden from view in the high state. Based on this, we caution that eclipse mapping of high-inclination SW Sex stars in the high state may yield incorrect or misleading results. In the context of DW~UMa, we demonstrate explicitly that distance estimates obtained by recent eclipse mapping studies cannot be reconciled with the WD-dominated low-state spectrum. We also show that the fluxes of the UV emission lines in the high state drop near orbital phase 0.5. This is the first detection of a UV counterpart to the class-defining phase 0.5 absorption seen in the optical emission lines of SW~Sex stars. ",
        "introduction": "\\label{introduction}  The SW Sex stars are a sub-class of nova-like cataclysmic variable stars (CVs). As such, they are interacting binary systems containing a Roche-lobe filling secondary star that is losing mass to an accretion disk around a white dwarf (WD) primary. The SW Sex classification was first suggested by Thorstensen et al. (1991), building on earlier work by Downes et al. (1986), Honeycutt, Schlegel \\& Kaitchuck (1986), Shafter, Hessman \\& Zhang (1988), and Szkody \\& Pich\\'{e} (1990). The key common features of the four founding members of the class -- SW Sex, DW UMa, V1315 Aql and PX And -- were transient, central absorption features in the Balmer and He~{\\sc i} lines near orbital phase 0.5. Also, the radial velocity curves derived from these lines generally showed significant phase lags with respect to the orbital motion of the WD. Other common spectroscopic properties within the class are single-peaked optical emission lines at orbital phases well away from 0.5 (rather than the double-peaked lines expected from edge-on accretion disks); weak eclipses of the Balmer and He I lines; and He~{\\sc ii} features that do appear to track the WD motion and are more deeply eclipsed than the low-ionization lines.  The continuum eclipses of high-inclination SW Sex stars are also unusual. They are considerably more V-shaped than those of \"ordinary\" eclipsing, non-magnetic nova-like CVs (such as UX UMa). If the eclipsed light is assumed to be produced by a geometrically thin, optically thick disk, this immediately implies that the outer disk contributes relatively more light in SW Sex stars than in other CVs. Unsurprisingly, eclipse mapping experiments have therefore invariably inferred very flat surface brightness (and hence effective temperature) distributions across the disks of SW Sex stars (e.g. Rutten, van Paradijs \\& Tinbergen 1992; Baptista, Steiner \\& Horne 1996; Baptista et al. 2000; Groot, Rutten \\& van Paradijs 2001). This contrasts starkly with the theoretically expected temperature distribution, which drops as $R^{-3/4}$~with radius.  Against this background came the first time-resolved, ultraviolet (UV) observations of a bona-fide SW Sex star --- DW UMa -- with the {\\em Hubble Space Telescope} (Knigge et al. 2000; Araujo-Betancor et al. 2003). These fortuitously occurred during a low state of the system, during which the optical light was suppressed by approximately 3 mags relative to the normal high state. Such low states are thought to be caused by a temporary reduction or cessation of the mass supply from the donor star. The UV observations were broadly consistent with this scenario, inasmuch as the low state UV spectrum appeared to be dominated by the hot ($T_{eff} \\simeq 50,000$~K) WD primary. However, a surprise was that the far-UV flux level was higher in the low-state HST data than in high-state UV observations obtained with the {\\em International Ultraviolet Explorer} (IUE). This implies that the WD in DW~UMa is hidden from view in the high state. We therefore suggested that the disk in this high-inclination system ($i \\simeq 82^{o}$; Araujo-Betancor et al. 2003) is self-occulting in the high state, i.e. that the disk rim obscures our view of the WD and the central disk regions.  If the WD and inner accretion flow in high-inclination SW Sex stars are permanently occulted, then this needs to be accounted for in any realistic model for these systems. After all, the majority of confirmed SW~Sex stars are eclipsing \\footnote{See DWH's {\\em Big List of SW Sex Stars}~at {\\url http://spider.ipac.caltech.edu/staff/hoard/biglist.html}.}, which indicates that geometric effects influence our ability to even recognize members of the class. It is therefore important to confirm the result of Knigge et al. (2000), not least to rule out cross-calibration errors between the high-state IUE data and the low state-HST observations.  Here, we present a first look at time-resolved, HST observations of DW UMa in the high state, which were obtained using exactly the same observational set-up as the low-state data. The results confirm conclusively that the WD in this system is hidden from view in the high state and show, for the first time, that the UV lines take part in the phase 0.5 absorption events that are the key optical characteristic of the SW~Sex class. ",
        "conclusions": "Our UV observations establish conclusively that the WD in DW UMa is hidden from view in the normal high state. We still believe that the simplest way to explain this finding is to postulate the existence of a self-occulting accretion disk in this system. However, independent of any specific theoretical framework, it is clear that the accretion flow cannot even approximately be described as 2-dimensional, and that any viable model or description of the system must account somehow for the observed occultation. It is therefore more difficult to use the eclipses to gain insight into the nature of the high-state accretion flow (but not impossible: see Hellier \\& Mason [1989] for an example in the context of low-mass X-ray binaries). By contrast, all existing high-state eclipse mapping studies of DW~UMa (and other SW~Sex stars) have assumed that the disks in these systems are fully visible. \\footnote{In their eclipse mapping study, Stanishev et al. (2004) comment on the occultation problem and suggest that the optically thick base of an accretion disk wind may hide the WD in DW~UMa. But this scenario is just a specific incarnation of a self-occulting disk atmosphere with unknown geometry, and thus still violates the assumption of full disk visibility. It is also worth recalling here the work of Rutten (1998), who shows that the application of a standard (flat-disk) eclipse mapping code to a flared disk can give reasonable results, {\\em as long as the disk is not self-occulting}. Once the inner disk is no longer visible, the method produces artificially flat radial temperature distributions.} In our view, the immediate corollary of this is that results inferred from these studies should be regarded with caution.  We suspect that the short distance estimates for DW UMa suggested in recent eclipse mapping studies ($d \\simeq 300$~pc; B{\\'{\\i}}r{\\' o} 2000; Stanishev et al. 2004) are a manifestation of this problem. Such a short distance is inconsistent with the estimate d~=590~$\\pm$~100~pc obtained from fitting the low-state spectrum with a WD model atmosphere and also with the estimate d~=~930~$\\pm$~160~pc obtained from Bailey's method (Araujo-Betancor et al. 2003). We have considered whether the WD-dominated low-state spectrum can in some way be reconciled with a distance of 300~pc. We thus show in Figure~1 three $\\log{g} = 8$~DA WD models (calculated as in G\\\"{a}nsicke, Beuermann \\& de Martino [1995]) for different effective temperatures ($T_{eff}$). The models have been scaled to match the observed low-state spectrum at 1480~\\AA, have not been convolved with the instrumental line spread function and have also not been rotationally broadened. Thus the width of Lyman~$\\alpha$ in the models is a lower limit.  The $T_{eff} = 50,000$~K model was chosen to represent the best-fitting WD model found by Araujo-Betancor et al. (2003). This model fits the data fairly well away from metal lines and requires a distance (d = 540~pc) in line with our previous estimate (this assumes a fully visible WD and the best-bet WD radius found by Araujo-Betancor et al.). By contrast, the eclipse-mapping distance of d = 300~pc requires a WD model with $T_{eff} \\simeq 28,000$~K. This model produces much too wide a Lyman~$\\alpha$~feature and can be ruled out immediately. If only half the WD were visible in the low state (e.g. because the lower half is blocked by a disk), d = 300~pc corresponds to $T_{eff} \\simeq 37,000$~K. This model still overpredicts the width of Lyman~$\\alpha$~significantly and can also be ruled out.\\footnote{We actually do not consider this to be a physically realistic model: if an optically thick disk were present in the low state, it would be strongly irradiated and therefore hot. It should then emit copious amounts of UV radiation, but no disk contribution is seen.}  Is it plausible that the flat-disk assumption in eclipse mapping studies might cause distances to be underestimated? Roughly speaking, these estimates are based on the color and brightness of the pixels in a disk map, and assuming these to radiate as blackbodies (BBs) or stellar atmospheres (SAs). If the disk is assumed to be flat, the flux from each pixel then scales as $\\cos{i} / d^2$ (modulo a correction for limb-darkening); this allows $d$~to be estimated. However, most other accretion structures (e.g. a disk rim) would be {\\em less} strongly foreshortened when averaged over the visible projected area of the structure. The flat-disk scaling will yield erroneous results in such cases. To see this, consider an accretion flow element that is seen more face-on than the pixel representing it in a flat-disk eclipse map. The temperature of this element/pixel is roughly fixed by its color. If a $\\cos{i}$ foreshortening factor is incorrectly applied to this element/pixel by the eclipse mapping algorithm, the distance inferred from it will be an underestimate. After all, the excessively foreshortened disk map pixel must still produce the observed flux. This is only possible if it is brought closer than the face-on accretion flow element it represents. Thus the flat disk assumption might indeed cause distances to be underestimated systematically. When this happens, the color-magnitude diagrams constructed from disk map pixels will probably also show large scatter about the expected BB/SA relations (since different accretion flow elements are likely to have different foreshortening factors). This would appear to be the case in DW~UMa (see Figure~5 in B{\\'{\\i}}r{\\' o} 2000).  As a final note on the distance issue, we acknowledge that the distance estimates based on Bailey's method and the WD-dominated  UV spectrum are themselves only marginally consistent with each other (see discussion in Araujo-Betancor et al. 2003). However, these estimates are much more easily reconciled with each other than with the short eclipse mapping distance. We have already shown explicitly that a WD at 300~pc produces a much broader Lyman alpha feature than is observed. But this distance would also require an unrealistically late spectral type for the secondary (to avoid over-predicting the low-state, mid-eclipse K-band flux).  Turning to our detection of a UV counterpart to the phase 0.5 absorption events seen in optical spectra of SW~Sex stars, we note that this could not necessarily have been expected: the formation mechanism of these high-ionization (and  mostly resonance) lines might be expected to differ significantly from that of the optical low-ionization lines where the effect is normally seen. Strictly speaking, we cannot yet claim that the UV lines are affected by genuine {\\em absorption} near phase 0.5: none of the lines exhibit dips below continuum level at this phase, so the data are also consistent with an intrinsic weakening of the lines. However, the Balmer lines of SW~Sex stars generally also do not show absorption below continuum, although this effect has been seen in the He~{\\sc i} lines of several systems (e.g. Szkody \\& Pich\\'{e} 1990). On balance, we are therefore inclined to ascribe the weakening of the UV lines to the same absorbing medium that is affecting the optical lines. Given the wide range of transitions affected, this probably implies non-uniform physical conditions in the absorbing material.   \\vspace*{0.1cm} This work made use of Peter van Hoof's online atomic line list. \\footnote{\\url{http://www.pa.uky.edu/\\textasciitilde peter/atomic}} Financial support for KSL, SAB, PS and DWH was provided by NASA through grant numbers GO-9791 and GO-7362 from the Space Telescope Science Institute (STScI), which is operated by AURA, Inc.,under NASA contract NAS5-26555. RIH is supported by NASA through Hubble Fellowship grant HF-01150.01-A awarded by STScI. BTG was supported by a PPARC Advanced Fellowship."
    },
    "0410/astro-ph0410547_arXiv.txt": {
        "abstract": "We have conducted high resolution continuum observations at 1.4 and 2.4 GHz with the Australia Telescope Compact Array, towards the four southern Of stars: CD-47\\,4551, HD\\,93129A, HD\\,124314, and HD\\,150136. All stars have been detected at the two frequencies. HD\\,93129A -- the only O2 I star catalogued so far, and in a double system --, has also been observed at 17.8 and 24.5 GHz. Its radio spectrum, complemented with previous observations at higher frequencies, is analyzed here. The interpretation yields the estimate of its mass loss rate, and a non-thermal spectral index of radiation coming from a putative colliding wind region. The synchrotron and corresponding inverse Compton luminosities are derived. ",
        "introduction": "Early-type stars (O to $\\sim$B2) develop strong stellar winds that strengthen while they are leaving the main sequence towards the Wolf-Rayet phase. The mass loss can be as high as 10$^{-5}$ M$_{\\odot}$\\,yr$^{-1}$ at velocities of thousands of km s$^{-1}$. The winds are optically thick, radiate through the free-free mechanism, and produce an excess of flux density from the radio to the infrared range. The thermal nature of the radiation can be confirmed computing a spectral index of  about 0.6 -- 0.8 \\citep{wb1975}. Once this radiation regime is established, the detection of the wind region, and consecutive measurement of the flux density allow a straightforward determination of the mass loss rate \\citep{ll1993}. However, emission characterized by a non-thermal spectral index has been reported repeatedly from early-type stars \\citep[e.g.][]{bac1989,bck2001}. Given the fact that the majority of these stars are not single, this emission is probably produced at a colliding wind region (CWR) of a system formed by at least two early-type stars with winds.  The  particles accelerated at the CWR can be also involved in high energy processes. Many of these stars have been detected in X-rays \\citep[e.g.][]{po1987}, and are proposed as counterparts of unidentified gamma-ray sources \\citep{rbt1999}. Multifrequency studies -from radio to gamma rays- are fundamental to completely describe the emitting regions, and help in a comprehensive picture of the whole stellar wind phenomenon.  In the frame of a southern radio survey on early-type stars, using the Australia Telescope Compact Array (ATCA) at maximum angular resolution, we have detected four non-thermal emitters at 3 and 6 cm \\citep{bck2001,bk2004}: CD-47 4551, HD 93129A, HD 124314, and HD 150136. In order to  better define the radio spectrum, we have performed ATCA observations of these sources, at 1.4 and 2.4 GHz, and here the results are presented. ",
        "conclusions": "Non-thermal contribution to the radio emission has been detected for the stars CD-47\\,4551, HD\\,93129, HD\\,124314, and HD\\,150136. The spectral fitting for the radio emission from the O2 If* HD\\,93129A enabled to extract the thermal flux, and determine a value for the mass loss rate of this peculiar system, of $3.6\\,\\times\\,10^{-5}$ M$_\\odot$ yr$^{-1}$. The non-thermal radiation, coming probably from a colliding wind region between the stellar components, is characterized by a flux density spectral index of $-1.3\\pm0.2$, and a magnetic field of $\\sim$ 20 mGauss."
    },
    "0410/astro-ph0410367_arXiv.txt": {
        "abstract": "{In this paper we consider the problem of the limits concerning the physical information that can be extracted from the analysis of one or more time series (light curves) typical of astrophysical objects. On the basis of theoretical considerations and numerical simulations, we show that with no a priori physical model there are not so many possibilities to obtain interpretable results. For this reason, the practice to develop more and more sophisticated statistical methods of time series analysis is not very productive. Only techniques of data analysis developed in a specific physical context can be expected to provide useful results. The field of {\\it stochastic dynamics} appears to be an useful framework for such an approach. In particular, it is shown that modelling the experimental time series by means of the stochastic differential equations (SDE) represents a valuable tool of analysis. For example, the use of SDE permits to make the analysis of a continuous signal independent from the frequency sampling with which the experimental time series have been obtained. In this respect, an efficient approach based on the extended Kalman-filter technique is presented. Freely downloadable software is made available. ",
        "introduction": "The study of the light curves of astrophysical objects has always been an important tool for astronomers. The reason is simple: an effective way to get insight on the structure of a given physical system is to study its evolution over time. Some examples are the reconstruction of the structure of the binary star systems, the understanding of the nature of the pulsars, and the determination of the sizes of the central regions of the active galactic nuclei. However, in spite of these remarkable successes, in many other situations the analysis of the light curves has not proved to be so useful. The reason can be understood by taking into account that often the time evolution of a dynamical system, describable in terms of a set of generic physical quantities ({\\it state-variables}) $\\xb(t)$ \\footnote{Hereafter, vector quantities will be denoted in boldface.}, is governed by a n-dimensional system of differential equations ({\\it state-equation}) with the general form \\begin{equation} \\label{eq:evolves} \\xdb(t) = \\fb[\\xb(t), \\ssb(t), t], \\end{equation} where $t$ is the time coordinate, symbol ``~$\\boldsymbol{\\dot{}}~$'' means derivation with respect to $t$, and $\\fb[\\cdot]$ is a $n$-dimensional (possibly non-linear) function. The $m$-dimensional vector $\\ssb(t)$ represents independent processes whose time evolution does not depend on $\\xb(t)$ such as, for example, the processes that take place in regions external to the system of interest. In general, the quantities $\\xb(t)$ are not directly observable and the experimental time series $\\{ \\yb_{t_k} \\}_{k=0}^N$ are obtained through the {\\it measurement-equation} \\begin{equation} \\yb_{t_k} = \\hb[\\xb(t_k), t_k] + \\eb_{t_k}. \\end{equation} Here, $\\hb[\\cdot]$ is a $l$-dimensional (possibly non-linear) function, $\\{ t_k \\}_{k=0}^N$ is the set of sampling time instants, and $\\{ \\eb_{t_k} \\}_{k=0}^N$ represents the measurement errors. Usually, the number $l$ of available time series is smaller than $n$, and often $l=1$. This means that the observed signals provide information only on a projection of the dynamics of the system of interest. Curiously, instead of developing new methodologies for the analysis of the observed signals in a given physical context, in the past much effort has been spent in an attempt to devise more and more sophisticated techniques for the {\\it statistical} characterization of $\\yb(t_k)$ (e.g. AR and ARMA modelling, maximum entropy power spectra, ...). The expectation of these efforts was that such a characterization could be able to provide hints on the functional form of $\\fb[.]$. However, the results have very often been disappointing since the experimental time series do not contain all the information necessary for such a task.  This does not mean that the classic statistical analysis of the time series is unproductive. However, it has to represent only a starting point, otherwise there is the risk that the studies on the time evolution of the astrophysical objects could merely consist in a collection of data and in the elaboration of some generic statistical measures with no direct physical meaning.  On the basis of this argument, \\citet{vio92} stress the necessity to start directly from model~(\\ref{eq:evolves}), as provided by a given theoretical model, and to use time series as a test for such model. More in particular, from the consideration that many astrophysical objects show evolutions that are unpredictable over time, they suggest to modify Eq.~(\\ref{eq:evolves}) to \\begin{equation} \\label{eq:evolves1} \\xdb(t) = \\fb[\\xb(t), \\uub(t), \\wb(t), t, \\thetab], \\end{equation} with $\\uub(t)$ and $\\wb(t)$ representing deterministic and random processes, respectively, and then to solve them via a numerical approach. In other words, the functional form of $\\fb[ \\cdot ]$ is assumed known out of a set $\\thetab$ of parameters. In this way, once fixed the values of $\\thetab$, it is possible to obtain  ``{\\it synthetic}'' light curves that can be compared with the observed ones. The reason to add the random process $\\wb(t)$, typically a continuous Gaussian white noise\\footnote{NB. Hereafter, with the term noise we will indicate only the random processes perturbing the dynamics of a given physical system and not the contamination due to the measurement errors.}, is a consequence of the fact that this term represents the interaction of the physical system of interest with its surroundings and/or the action of complex processes that cannot be directly included into the model (e.g. gas turbulence). In general, such processes are characterized by an huge number of degrees of freedom and therefore they can be assumed to have a stochastic nature. In practice, this means to study the time evolution of a given physical system in the context of the so called {\\it stochastic dynamics}, i.e., through the modelling of the observed time series by means of stochastic differential equations (SDE).  Although {\\it stochastic dynamics} is an approach widely used for the study and simulation of realistic scenarios in many fields of applied science and engineering as, for example, fluid dynamics, structural and mechanical engineering, avionics, material properties, financial sciences \\ldots \\citep{gha91, klo97, gar99}, in Astronomy it is not so known. However, as we will show on the basis of some numerical experiments, even the action of quite weak noise sources on simple nonlinear dynamical systems can produce deep modifications of their behaviour over time. This means that in real scenarios the noise component must be considered as intrinsic to the physics of the systems and not only a secondary factor. Consequently, in many situations, {\\it stochastic dynamics} could represent the only possibility to use the experimental time series in an effective way.  In \\citet{vio92} it is suggested that the value of $\\thetab$ has to be derived from physical considerations. Here, we provide some tools that permits to estimate $\\thetab$ directly from the data.  In Sect.~\\ref{sec:why} some arguments are presented that support the necessity of stochastic modelling in the study of the physical systems, and in Sec.~\\ref{sec:astronomy} an example in the astronomical context is provided. In Sec.~\\ref{sec:discrete} it is shown that modelling the time series through discrete models is unsuited for most of physical systems. Hence, in Sec.~\\ref{sec:continuum} an approach based on SDE is suggested and some tools are provided for using it in modelling the time series. An example of application is given in Sec.~\\ref{sec:worked}. Finally, the conclusions and some final comments are presented in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The time series usually available in astronomy are able to characterize only a subset of the system of equations that describe the dynamics of the physical system under study. For this reason, although in principle it is always possible to determine a statistical model able to reproduce the experimental data, without any a priori physical model there are not so many possibilities to obtain a reliable reconstruction of the physical scenario investigated. In general, this means that an approach to the analysis of time series exclusively based on the experimental data will provide inconclusive results, and that the practice to search for more and more sophisticated statistical techniques is not very productive. In many situations, the only possibility to get some physical insights is to carry out the analysis in a well defined physical context. In this respect, {\\it stochastic dynamics}, i.e. modelling the time series with stochastic differential equations, appears to be a promising tool. A benefit of working within such a framework is that one is forced to provide a mathematical/physical formalization of the starting hypotheses adopted in the analysis of the signals (e.g. linearity, non-linearity, type of non-linearity, ...). In this way, there is no risk of misunderstandings concerning the interpretations of the final results. Additionally, a more direct relation between complex physical models and the limited observational material is made possible permitting a more efficient interaction between data and theory.  A code, implementing some efficient numerical tools for modelling the experimental time series with stochastic differential equations, is freely downloadable from {\\it http://www.imm.dtu.dk/ctsm/}."
    },
    "0410/astro-ph0410198_arXiv.txt": {
        "abstract": "We present the results of the analysis of light curves of V2051 Oph through an outburst with eclipse mapping techniques. ",
        "introduction": "\\hskip 0.55truecmDwarf novae show recurrent outbursts (of 2-5 mags on timescales from weeks to months) powered by a sudden increase in mass inflow in the accretion disc around the white dwarf.  Currently, there are two competing models to explain the cause of the sudden increase in mass accretion. In the mass transfer instability model (MTIM), the outburst is the time dependent response of a viscous accretion disc to a burst of matter transferred from the secondary star. In the disc instability model (DIM), matter is transferred at a constante rate to a low viscosity disc and accumulates in an annulus until a critical configuration switches the disc to a high viscosity regime and the gas diffuses rapidly inwards and onto the white dwarf. Tracking the evolution of a dwarf novae accretion disc along an outburst cycle with eclipse mapping techniques (EMT) provides a valuable opportunity to test these models againt observations. ",
        "conclusions": ""
    },
    "0410/astro-ph0410151_arXiv.txt": {
        "abstract": "\\noindent  In this Paper we report on optical spectroscopic observations of the low-mass X-ray binary 2S~0921--630 obtained with the Very Large Telescope. We found sinusoidal radial velocity variations of the companion star with a semi--amplitude of 99.1$\\pm$3.1 km s$^{-1}$ modulated on a period of 9.006$\\pm$0.007 days, consistent with the orbital period found before for this source, and a systemic velocity of 44.4$\\pm$2.4 km s$^{-1}$.  Due to X--ray irradiation the centre--of--light measured by the absorption lines from the companion star is likely shifted with respect to the centre--of--mass. We try to correct for this using the so--called K--correction.  Conservatively applying the maximum correction possible and using the previously measured rotational velocity of the companion star, we find a lower limit to the mass of the compact object in 2S~0921--630 of $M_{X} \\sin^3 i > 1.90\\pm0.25\\,M_\\odot$ (1 $\\sigma$ errors). The inclination in this system is well constrained since partial eclipses have been observed in X-rays and optical bands. For inclinations between 60$^\\circ < i< 90^\\circ$ we find $1.90\\pm0.25< M_{X} < 2.9\\pm0.4\\,M_\\odot$. However, using this maximum K--correction we find that the ratio between the mass of the companion star and that of the compact object, q, is $1.32\\pm0.37$ implying super--Eddington mass transfer rates; however, evidence for that has not been found in 2S~0921--630.  We conclude that the compact object in 2S~0921--630 is either a (massive) neutron star or a low--mass~black~hole. ",
        "introduction": "\\label{intro} The equation of state (EoS) of neutron star matter is intimately related to the physics of the strong interactions between fundamental particles and therefore it is of great relevance in high--energy and particle physics. Theories on the EoS of neutron--star matter at supra--nuclear densities provide a firm upper limit on the mass for each EoS. The stiffer the EoS, the higher the mass limit. Neutron stars with masses well above $1.4\\,M_\\odot$ cannot exist for so--called soft equations of state, in which matter at high densities is relatively compressible (e.g., due to meson condensation or a transition between the hadron and quark--gluon phases; cf.~\\citealt{2001ApJ...550..426L}). Therefore, measuring a high mass for even one neutron star would imply the firm rejection of many proposed soft EoSs (e.g.~see discussion by \\citealt{1995xrb..book...58V}).  Neutron star masses of more than $\\sim 3\\,M_{\\sun}$ are excluded assuming (among other things) that general relativity is the correct theory of gravity and that the velocity of sound is less than the velocity of light (\\citealt{1973ApJ...179..277N}). Therefore, a lower limit on the mass of the compact object of 3 $M_{\\odot}$ or more implies a black hole compact object.  The most accurately measured neutron star masses we have are from precise radio timing measurements of double neutron--star binaries. These empirical masses cluster around $1.35\\,M_\\odot$ (\\citealt{1999ApJ...512..288T}). This is close to the mass of 1.32 $M_{\\odot}$ predicted from theoretical model calculations of Type Ib supernovae (\\citealt{1996ApJ...457..834T}). However, these double neutron star binaries are formed in high--mass X--ray binaries and due to the short life of the massive companion only a limited amount of matter can be accreted. In contrast, in low--mass X--ray binaries (LMXBs), the rapid rotation and low magnetic field of the neutron stars are considered evidence for much larger amounts of accreted matter (see \\citealt{bhatta1995} for a review).  Indeed, from binary evolution models, LMXBs could accrete up to 0.7\\,$M_\\odot$ (e.g.~\\citealt{1995A&A...297L..41V}). If so, the masses of their neutron stars could be $>\\!2\\,M_\\odot$. Most of the constraints from timing observations of the successors of these LMXB primaries, the millisecond radio pulsars, are weak (\\citealt{1999ApJ...512..288T}), as the almost purely circular orbits of these systems often prevents one from measuring most of the post--Newtonian parameters that underlie the precise masses for the double--neutron star binaries. With masses of $1.57^{+0.25}_{-0.2}\\,M_\\odot$ and 2.2$\\pm0.4\\,M_\\odot$, the msec pulsars PSR~B1855$+$09 and PSR~J0751+1807 are exceptions to this (\\citealt{2003rapu.conf...75N}; \\citealt{2004aspin.nice},~95\\%~conf.).  Dynamical mass estimates or dynamically determined limits on the mass of the neutron star in LMXBs are available for the two transient sources Cen~X--4 and XTE~J2123--058, the Z--source Cyg~X--2, and the pulsar 2S~1822--371 (Cen~X--4; $M_{NS}=1.3\\pm$0.6 $M_\\odot$ \\citealt{1993MNRAS.265..655S}; XTE~J2123--058; $M_{NS}=1.55\\pm$0.31 $M_\\odot$ \\citealt{2002MNRAS.329...29C}; Cyg~X--2; $M_{NS}=1.78\\pm$0.23 $M_\\odot$ \\citealt{1998ApJ...493L..39C}; \\citealt{1999MNRAS.305..132O}; 2S~1822--371; $M_{NS}>1.14\\pm$0.06 $M_\\odot$; \\citealt{2003MNRAS.339..663J}; \\citealt{2003ApJ...590.1041C}). The problem is that many of the companion stars are small, late type stars which at a distance of more than a few kpc are too faint to observe, especially since the integration times are limited to typically 1/20th of the orbital period in order to avoid Doppler smearing of the spectral lines due to the binary motion. Furthermore, the binary inclination is often poorly constrained and effects of irradiation shift the centre-of-light with respect to the centre-of-mass.  Some of the aforementioned problems will be less severe when observing long orbital period systems for which the inclination is well constrained from the fact that eclipses are observed: 2S~0921--630/V395~Car is such a high inclination LMXB with a long orbital period (P$_{orb}\\approx 9.0$ days; e.g.~\\citealt {1982ApJ...256..605C}, \\citealt{1983MNRAS.205..403B}, and \\citealt{1987MNRAS.226..423M}). The inclination of 2S~0921--630 is relatively well constrained and must be high since partial eclipses of the compact object and accretion disc have been observed in both X--rays and optical wave bands (e.g.~\\citealt{1981SSRv...30..279B}; \\citealt{1981A&A....94L...3C}; \\citealt{1987MNRAS.226..423M}). 2S~0921--630 is thought to be a halo object in orbit with a K0--1 III companion star; absorption lines of the companion have been detected  (\\citealt{1983MNRAS.205..403B}; \\citealt{1999A&A...344..101S}). The corresponding radial velocity curve would provide a limit to the mass of the companion star. In this Paper we report on spectroscopic observations of 2S~0921--630/V395~Car. ",
        "conclusions": "We determined the radial velocity curve of the companion star in 2S~0921--630/V395~Car using VLT/FORS2 spectroscopic observations. A sinusoidal fit to the radial velocity measurements with a semi--amplitude, K$_2$, of 99.1$\\pm$3.1 km s$^{-1}$ and with a systemic velocity, $\\gamma=$44.4$\\pm$2.4 km s$^{-1}$ represents the data well. In these error estimates we artificially increased the error bars on the individual velocity measurements by a factor 4 in order to obtain a reduced $\\chi^2$ of the fit close to 1. Before we discuss the implications of the K--velocity, we investigate possible reasons for the large reduced $\\chi^2$ obtained when using the formal statistical errors only.  Near quadrature at phase 0.25 the velocity measurements obtained on different nights differ well beyond the statistical errors.  From Figure~\\ref{compfrac} we can see that the companion star contribution to the total amount of light in the 5920--6520 \\AA~range varies as a function of orbital phase. The phasing of this variation suggests that the centre of light emitted by the K--star is likely to be shifted with respect to its centre of mass. An explanation is that the inner side of the companion star is heated by the (X--ray) irradiation coming from (near) the compact object, reducing the equivalent width of the K--star stellar lines in the range 5920--6520 \\AA~(see also e.g.~\\citealt{1988ApJ...324..411W}). This manifests itself in several ways:\\newline {\\it (i)} a lower fractional contribution to the total amount of light from the K star at e.g.~phase 0.5 when using the optimal subtraction technique (e.g.~Figure~\\ref{compfrac})\\newline {\\it (ii)} an increase in velocity (especially near quadrature) with respect to the velocity associated with the centre of mass of the companion star, since the line profiles will, as a result of the stellar rotation, be preferentially shifted to either a larger red- (phase 0.25) or blue-shift (phase 0.75). Hence, variations in the amount of (X--ray) irradiation between the observations could explain the variations in velocities near quadrature as observed in Figure~\\ref{radvel}. Since the observations were obtained with at least one day in between, with a baseline of several months, it seems likely that the irradiation changed between the observations.  Furthermore, the velocity measurements around phase 0.5 do not fall on the best--fit sinusoid. Those at phases 0.4--0.5 lay above the best--fit sinusoid, whereas those at phases 0.5--0.6 fall below the curve. This could well be explained by the Rossiter effect (cf.~\\citealt{2001icbs.book.....H}). Such deviations occur if the companion star is partially eclipsed by the accretion disc. Such a partial eclipse of the companion star can also help explain the low fraction of light coming from the companion star at phase 0.5 (see Figure~\\ref{compfrac}). The radial velocity measurements at phases just prior to phase 0.5, the partial eclipse of the companion star, are biased to higher, red shifted velocities since the part of the companion star that is eclipsed at those phases is mainly rotating towards the observer. Hence, the rotational broadening of the line will be biased towards the red. The reverse holds for phases just after the partial eclipse.  For the rotational broadening of the stellar lines we take $v \\sin i=64\\pm9$ km s$^{-1}$ as found by \\citet{1999A&A...344..101S}.  For a Roche lobe filling companion star the following relation holds: $\\frac{{\\rm v\\,sin}i}{{\\rm K}_2}=0.46[(1+q)^2q]^\\frac{1}{3}$ (e.g.~\\citealt{1988ApJ...324..411W}; q is defined here as M$_2$ divided by the mass of the compact object, M$_X$). Before we can determine q, and hence with K$_2$ the mass of the compact object, we have to take effects of the non--uniform absorption distribution on the stellar surface caused by irradiation into account. To this end, \\citet{1988ApJ...324..411W} developed a so--called K--correction. Following the procedure laid out in \\citet{1988ApJ...324..411W}, ${\\rm K_{2,corr}=K_{obs} - \\Delta K}$.  $K_{2,corr}$ is the corrected observed ($K_{obs}$) radial velocity semi--amplitude and $\\Delta K = f R_2 K_2/a_2 = f v \\sin i$. Here $a_2$ is the distance between the centre-of-mass of the binary and the centre-of-mass of the companion star and $f$ is a geometrical correction factor less than 1. In the extreme case that the hemisphere facing the compact object does not contribute to the observed stellar absorption lines at all and the other hemisphere has a uniform absorption line strength, $f=\\frac{4}{3\\pi}$ (\\citealt{1988ApJ...324..411W}). Since we still have a contribution of the K1 star of around 7 per cent at phase 0.5 we know that this is an overestimation of the K--correction in the case of 2S~0921--630. Hence, applying this upper limit on the K--correction, $\\Delta K \\leq 26.9\\pm3.8$ km s$^{-1}$, we get a firm lower limit of K$_{2,corr}\\geq 72.2\\pm4.9$ km s$^{-1}$. Plugging--in the numbers given above yields q$=$1.32$\\pm0.37$; note that in deriving this we used the equation for $R_2/a$ of \\citet{paczynski1971} which is valid for $0<q<0.8$. However, the difference with the equation of \\citet{1983ApJ...268..368E} is small for q$<$1.7 and does not change the result significantly given the large uncertainty in $v \\sin i$. The uncertainty in q is dominated by the large error on the rotational velocity. Thus, we find for the minimum mass of the compact object, $M_X \\sin^3 i > 1.90\\pm0.25\\,M_\\odot$. As mentioned in the Introduction, the inclination of 2S~0921--630 must be high. Furthermore, we have presented evidence for a partial eclipse of the companion star. For inclinations, $60^\\circ<i<90^\\circ$ (cf.~\\citealt{1983MNRAS.205..403B}; \\citealt{1987MNRAS.226..423M}) the mass of the compact object is $1.90\\pm0.25<M_X<2.9\\pm0.4\\,M_\\odot$.  Using $v \\sin i = 2 \\pi R_2 \\sin i/P$, we get for the radius of the companion star, $R_2$, $ R_2 \\sin i= 11.4\\pm1.6\\,R_\\odot$. Under the assumption that the companion star fills its Roche lobe, the mass of the donor star, $M_2$, depends only on its radius and the orbital period (cf.~\\citealt{paczynski1971}), which gives $M_2 \\sin^3 i= \\frac{(R_2\\sin i)^3}{0.234^3 \\times P_{orb}^2}$ (here the orbital period is measured in hours, and the mass and radius of the companion star are in solar units). This gives $M_2 \\sin^3 i=2.5\\pm1.1\\,M_\\odot$. This M$_2$ suggests an interesting evolutionary stage of the binary. M$_2$ is more massive than standard donor stars in LMXBs. Recently, there have been a number of evolutionary calculations (e.g.~\\citealt{2000ApJ...530L..93T}; \\citealt{2002ApJ...565.1107P}) showing that systems with donor masses as high as $\\sim$5 M$_\\odot$~can enter a relatively long phase of mass transfer to a neutron star after the donor has lost most of its mass. 2S~0921--630 would nicely fit in this scenario, evolving along the track of an initial 3--4 M$_\\odot$~donor. However, in this evolutionary scenario the current time-averaged mass-transfer rate is expected to be high, certainly super-Eddington.  This is at odds with the finding from studying X-ray lines as measured with {\\it Chandra} and XMM--{\\it Newton} that the observed X-ray luminosity of $\\sim10^{36}$ erg s$^{-1}$ (for a distance of $\\sim$7 kpc) is likely close to the intrinsic source luminosity (\\citealt{2003ApJ...583..861K}). Perhaps the current mass-transfer rate is much lower than the time-averaged transfer rate, or the donor mass and/or the mass-ratio are overestimated.  Above, we assumed that the companion star is in co--rotation with the binary orbit. Synchronization of the rotational period and the binary period occurs on timescales shorter than the circularization timescale (e.g.~\\citealt{1988ApJ...324L..71T}) and no signs for orbital eccentricity have been found (see Figure~\\ref{radvel}). Furthermore, using the companion star parameters derived above and equation 7 of \\citet{1988ApJ...324L..71T} we find that the synchronization timescale is $\\approxlt$100 yr for 2S~0921--630, justifying this assumption.  We conclude that the compact object in 2S~0921--630 is likely to be a massive neutron star or a low-mass black hole. If indeed the compact object in 2S~0921--630 is a massive neutron star the mass measurement would rule out soft equations of state.  Future high resolution optical observations should provide a much more accurate handle on the rotational velocity of the companion star, yielding a more accurate constraint on the mass of the compact object."
    },
    "0410/astro-ph0410221_arXiv.txt": {
        "abstract": "We show that loitering at high redshifts ($z \\gsim 6$) can easily arise in braneworld models of dark energy which, in addition to being spatially flat, also accelerate at late times. Loitering is characterized by the fact that the Hubble parameter dips in value over a narrow redshift range which we shall refer to as the `loitering epoch'. During loitering, density perturbations are expected to grow rapidly. In addition, since the expansion of the universe slows down, its age near loitering dramatically increases. An early epoch of loitering is expected to boost the formation of high redshift gravitationally bound systems such as $10^9 M_\\odot$ black holes at $z \\sim 6$ and lower-mass black holes and/or Population III stars at $z > 10$, whose existence could be problematic within the LCDM scenario. Loitering models also help to reduce the redshift of reionization from its currently (high) value of $z_{\\rm reion} \\simeq 17$ in LCDM cosmology, thus alleviating a significant source of tension between observations of the high-redshift universe and theoretical model building. Currently a loitering universe accelerates with an effective equation of state $w < -1$ thus mimicking phantom dark energy. Unlike phantom, however, the late-time expansion of the universe in our model is singularity free, and a universe that loitered in the past will approach a LCDM model asymptotically in the distant future. ",
        "introduction": "Standard Big Bang cosmology, as epitomized in its most recent avatar, LCDM, is in excellent agreement with a host of cosmological observations including galaxy clustering, fluctuations in the CMB, and the current accelerating epoch. Yet it appears that recent observations at modest redshifts ($6\\lsim z \\lsim 20$) may have some surprises in store for LCDM.  \\begin{itemize}  \\item[(i)]  In less than a decade of observations, the number of known high redshift QSO's has shown an almost twenty-fold increase\\,! Indeed, over $400$ QSO's with redshifts $z > 4$ are known at present, and the seven highest redshift quasars have $z > 5.7$ \\cite{richards03}. If quasars shine by virtue of an accreting black hole at their centers, then all these QSO's must host $\\gsim 10^9 M_\\odot$ black holes. Whether such highly massive black holes can successfully form in a LCDM universe which is less than a billion years old at $z \\sim 6$ remains an open question, but most theorists seem to agree that theoretical models of the growth of black holes, whether by accretion or through BH--BH mergers, are under some tension to explain the observations \\cite{richards03,bh}.  \\item[(ii)]  In addition to the presence of large supermassive black holes at $z \\sim 6$, there is indirect evidence to suggest that a population of less massive black holes and/or Population~III stars was already in place by $z \\gsim 17$ and may have been responsible for ionizing the universe at lower redshifts.\\footnote{WMAP observations give $\\tau = 0.17 \\pm 0.06$ for the optical depth which translates into $z = 17 \\pm 5$ for the reionization redshift in LCDM cosmology \\cite{wmap}.} Whether the LCDM model can form structure early and efficiently enough to successfully reionize the universe by $z \\sim 17$ is a moot point \\cite{reionization}. In any case, both (i) and (ii) provoke the concerned cosmologist to look for alternative models, which, while preserving the manifold strengths and successes of LCDM, will also be able to provide a compelling resolution to the issues raised above. In this paper, we show that one such model --- a braneworld universe which loiters at an early epoch --- may provide an attractive alternative to LCDM.  \\end{itemize} ",
        "conclusions": "We have demonstrated that a loitering universe is possible to construct within the framework of braneworld models of dark energy. An important aspect of braneworld loitering is that, in contrast to the conventional loitering scenarios that demand a closed universe, loitering on the brane can easily occur in a spatially flat cosmological model. A key role in making the brane loiter is the presence of (negative) dark radiation --- a generic five-dimensional effect associated with the projection of the bulk gravitational degrees of freedom onto the brane. Our universe can loiter at large redshifts ($z \\gsim 6$) while accelerating at the present epoch.\\footnote{Although both the degree of loitering and the loitering redshift are free parameters in our model whose values can be determined by matching to observations, the loitering braneworld nevertheless does not claim to resolve the `cosmic coincidence' conundrum associated with the current value of the effective cosmological constant, which plagues most dark-energy models and usually requires some degree of fine tuning of cosmological parameters \\cite{sahni04,ss00}.} During loitering, the value of the Hubble parameter decreases steadily before increasing again. As a result, the age of the loitering braneworld is larger than that of a LCDM universe at a given redshift. This feature may help spur the formation of $\\sim 10^9M_\\odot$ black holes at redshifts $\\gsim 6$ whose presence (within high redshift QSO's) could be problematic for standard LCDM cosmology \\cite{richards03}.\\footnote{From Fig.~\\ref{fig:age}, we find that the age of a loitering universe at $z \\sim 6$ can be several times that in LCDM cosmology, which is less than a billion years old at that redshift.} Loitering is also expected to increase the growth rate of density inhomogeneities and could, in principle, be used to reconcile structure formation models which predict a lower amplitude of initial `seed' fluctuations with the observed anisotropies in the cosmic microwave background (see \\cite{dodelson} and references therein). In addition, an early loitering phase could lower the redshift of reionization from its currently high value of $z \\simeq 17$ for the LCDM model \\cite{wmap}. Finally, we would like to draw attention to the fact that earlier work on braneworlds has emphasized departure from the standard Friedmannian behaviour either in the distant past ($z \\gsim 10^{9}$) \\cite{BDL} or else, in the current epoch and remote future ($z \\lsim 2$) \\cite{DGP,DDG,SS,padilla}. In this paper, we have shown that a braneworld can also show interesting significant departures from the conventional behaviour at {\\em intermediate\\/} redshifts $z \\gsim 6$. It is meaningful to ask ourselves whether this feature of braneworld cosmology is a unique aspect of the higher-dimensional action (\\ref{action}) or whether such properties are shared by a larger class of modified gravity and string-inspired models. Perhaps future work will throw light on this question.   {\\em Acknowledgments\\/}: It is a pleasure to thank Kandaswamy Subramanian, Subhabrata Majumdar and Carlo Contaldi for stimulating discussions and Alexei Starobinsky for valuable comments on an early version of the paper. The authors acknowledge support from the Indo-Ukrainian program of cooperation in science and technology sponsored by the Department of Science and Technology of India and Ministry of Education and Science of Ukraine.  \\bigskip  \\begin{appendix}"
    },
    "0410/astro-ph0410017_arXiv.txt": {
        "abstract": "We report on $XMM$-$Newton$ observations of G32.45+0.1 and G38.55+0.0. These were discovered as diffuse hard X-ray sources with the $ASCA$ Galactic plane survey, but the limited spatial resolution of $ASCA$ could not conclude whether these are truly diffuse or a group of unresolved point-sources. $XMM$-$Newton$, with higher spatial resolution than $ASCA$, confirmed that G32.45+0.1 has a diffuse shell-like structure with a radius of $\\sim 4^{\\prime}$. The spectrum shows featureless continuum, hence can be fitted with a power-law model of $\\Gamma \\sim 2.2$ with an absorption of  $N_{\\rm H}\\sim 5.2\\times 10^{22}~{\\rm cm}^{-2}$. From this $N_{\\rm H}$ value, we estimate the distance to G32.45+0.1 to be $\\sim 17$~kpc, then the luminosity (in the 0.5--10.0~keV band) and radius of the shell are $\\sim 9.5\\times 10^{34}$~ergs~s$^{-1}$ and $\\sim $20~pc, respectively. The radio complex sources in the NRAO/VLA Sky Survey  (NVSS; 1.4~GHz) are globally associated to the X-ray shell of G32.45+0.1. Therefore G32.45+0.1 is likely to be a synchrotron dominant shell-like SNR. No significant diffuse structure was found in  the $XMM$-$Newton$ image of another $ASCA$ diffuse source G38.55+0.0. The upper limit of the observed flux (0.5--10.0~keV) is  estimated to be $9.0\\times 10^{-13}$~ergs~cm$^{-2}$~s$^{-1}$ being consistent with the $ASCA$ result.  \\footnote[1]{Department of Physics, Graduate School of Science, Kyoto University, Sakyo-ku, Kyoto 606-8502} \\footnote[2]{RIKEN (The Institute of Physical and Chemical Research), 2-1 Hirosawa, Wako, Saitama 351-0198} \\footnote[3]{Faculty of Humanities and Social Sciences, Iwate University, 3-18-34 Ueda, Morioka, Iwate 020-8550} ",
        "introduction": "Since cosmic rays were discovered by Hess in 1912, the source and the mechanism of acceleration have been unknown. The spectrum of cosmic rays shows a single power-law up to the knee energy ($\\sim 10^{15.5}$~eV), where the gyro radius of electrons in typical interstellar magnetic field is much smaller than the Galactic radius. Therefore  cosmic rays below the knee energy is likely  to be Galactic origin. The discoveries of synchrotron X-rays and inverse Compton TeV $\\gamma $-rays from the supernova remnant (SNR) SN~1006 indicated that electrons are accelerated close to the knee energy (Koyama et al. 1995; Tanimori et al. 1998). The $Chandra$ results of the small-scale structure is successfully explained by a diffusive shock acceleration model (DSA); the first-order Fermi mechanism is working at the shock front of SNRs (Bamba el al. 2003b; Yamazaki et al. 2004).  Besides SN~1006, synchrotron X-ray emissions were discovered from other SNRs, G347.5-0.3 (RX~J1713.7$-$3946; Koyama et al. 1997; Slane et al. 1997), Cas~A (Vink et al. 2000), Tycho's SNR (Hwang et al. 2002), RX~J0852.0$-$4622 (Slane et al. 2001), RCW~86 (Borkowski et al. 2001a), AX~J1843.8$-$0352 (Bamba et al. 2001; Ueno et al. 2003), 30~Dor~C (Bamba et al. 2004). Although the synchrotron X-ray emissions were observed from several SNRs, the total number and the X-ray fluxes so far discovered are insufficient to account for all cosmic rays in our Galaxy; if SNRs are the main accelerator of the cosmic rays, we can expect more non-thermal SNRs. Accordingly, we searched for the SNR candidates in the data of the $ASCA$ Galactic plane survey which covered $|l|\\leq 45^{\\circ}$ $|b|\\leq 0.^{\\circ}4$ on the Galactic plane (Yamauchi et al. 2002). About half a dozen candidates were found in this survey area; follow-up deep exposure observations with $ASCA$ were made on three candidates, G11.0+0.0, G25.5+0.0, and G26.6$-$0.0, and these were suggested  to be  non-thermal SNRs (Bamba et al. 2003a). G28.6$-$0.1 (AX~J1843.8$-$0352) is deeply observed with both $ASCA$ and $Chandra$ and is established to be a synchrotron X-ray emitting shell-like SNR (Bamba et al. 2001; Ueno et al. 2003).  Two other candidates G32.45+0.1 and G38.55+0.0 are seen as diffuse-like hard sources  in the $ASCA$ images, and the spectra show a power-law like feature with $\\Gamma $=1--3. The spatial sizes ($\\sim$ a few arcmin) are, however, marginal compared with the $ASCA$ point spread function (PSF) of ($\\sim$ 1~arcmin) to determine whether these are diffuse or a group of unresolved point-sources. We, therefore, made $XMM$-$Newton$ (Aschenbach et al. 2000) observations on  G32.45+0.1 and G38.55+0.0 with higher spatial resolution (PSF$\\sim $5~arcsec) and larger effective area than those of $ASCA$. ",
        "conclusions": "\\subsection{G32.45+0.1}  We found a clear shell structure with a radius of $\\sim 4^{\\prime}$ in the $XMM$-$Newton$ hard band (2.0--7.0~keV) image. The X-ray spectrum is well-fitted with either a power-law or a thin thermal plasma model. The thermal model requires the temperature of 3.0--12~keV which is similar to (or higher than) that of typical young SNRs such as Cas~A and Tycho's SNR ($k$T $\\sim $2--4~keV; e.g., Vink et al. 1996, Gotthelf et al. 2001 (Cas~A); Fink et al. 1994 (Tycho's SNR)~). This model, however, also requires an uncomfortably low metal abundance ($<$ 0.35~solar abundance), compared with Cas~A and Tycho's SNR ($\\sim $2~solar abundance). Thus the spectrum of G32.45+0.1 is likely to be a power-law. The best-fit parameter of $N_{\\rm H}$=5.2~(3.9--7.5)$\\times 10^{22}~{\\rm cm}^{-2}$ gives rough estimation of the source distance to be 17~kpc with an assumption that the density in the Galactic plane is 1~H~cm$^{-3}$. Then the X-ray luminosity (0.5--10.0~keV) and the radius are estimated to be $\\sim 9.5\\times 10^{34}$~ergs~s$^{-1}$ and $\\sim 20$~pc, respectively. These values are not largely different from those of the non-thermal component of SN~1006 ( $\\sim 2.0\\times 10^{34}$~ergs~s$^{-1}$ and $\\sim $9.5~pc ; e.g., Dyer et al. 2004).  We found a radio shell at the position of G32.45+0.1 from the NRAO/VLA Sky Survey (NVSS) data at 1.4~GHz (Condon et al. 1998) overplotted with contours in figure \\ref{fig:G32_2color_img}, although the authors regarded  the radio shell as 7 individual (point-like) sources as listed in the NVSS catalogue. The  flux sum is $\\sim 1.7\\times 10^{-1}$~Jy. From the shell-like morphology in both the radio and X-ray bands and the X-ray power-law model of photon index of $\\Gamma $=2.2~(1.8--3.0), G32.45+0.1 is highly possible to be a synchrotron X-ray emitting shell-type SNR. For a discussion of a wide band spectrum, we fitted the X-ray spectrum with an SRCUT model (Reynolds and Keohane 1999). The radio index has not yet been determined, but simple acceleration model predicts the energy index to be 0.5. In fact, the energy index ($\\alpha $) at 1~GHz of SN~1006, the typical synchrotron X-ray SNR, is reported to be $\\sim $0.57 (Allen et al. 2001). Thus we tried fitting assuming $\\alpha $ to be 0.5 or 0.6. The best-fit parameters are shown in table \\ref{tab:srcut}.   \\begin{table}[h] \\begin{center} \\caption{The best-fit parameters by the SRCUT model$^a$} \\label{tab:srcut} \\begin{tabular}{lcc} \\hline \\hline Parameters & $\\alpha = 0.5^b$ & $\\alpha = 0.6^b$ \\\\ \\hline $N_{\\rm H} [\\times 10^{22}~{\\rm cm}^{-2}]$ & 5.2 (4.5--6.2) & 5.2 (4.4--6.1) \\\\ Cut-off frequency [$\\times 10^{17}$~Hz] & 3.4 (0.79--91) & 5.4 ($<$ 35) \\\\ Norm$^c$ & 0.019 (0.017--0.022) & 0.11 (0.095--0.12)  \\\\ Reduced $\\chi ^2$ (d.o.f.)  & 1.14 (46) & 1.14 (46) \\\\ \\hline \\multicolumn{3}{l}{$^a$Parentheses indicate 90\\% confidence intervals.} \\\\ \\multicolumn{3}{l}{$^b$The fixed value of $\\alpha $.} \\\\ \\multicolumn{3}{l}{$^c$The flux density at 1~GHz~[Jy].} \\end{tabular} \\end{center} \\end{table}   Here the absorption is consistent with the fitting result with the power-law model. Using the values of the cut-off frequency, we can estimate the maximum energy of electrons accelerated by the SNR shock according to Reynolds and Keohane (1999). Assuming the cut-off frequency is $5\\times 10^{17}$~Hz, the maximum electron energy ($E_{\\rm max}$) is estimated to be \\[ E_{\\rm max} \\sim 60 \\left( \\frac{10\\mu {\\rm G}}{B} \\right) ^{\\frac{1}{2}} {\\rm TeV.} \\] This result is consistent with the standard picture of the non-thermal SNRs. For example, $E_{\\rm max}$ and $B$ of SN~1006 are 20--70~TeV and 3.5--85~$\\mu $G, respectively (Bamba et al. 2003b).   In either $\\alpha =$ 0.5 or 0.6, the predicted flux density at 1.4~GHz is smaller than NVSS result. This apparent inconsistency may be attributable to a non-uniform shell structure. In fact,  the radio brightest region  show no X-ray shell (see figure \\ref{fig:G32_2color_img}). This can be  interpreted that the X-ray emission is dominated in relatively  low $B$ region, because the synchrotron energy loss rate is proportional to $B^2$, and hence the energy loss at the X-ray emitting electron energy may be negligible. The radio shell, on the other hand, may  come from a higher $B$ region, where energy loss of high energy electron is significant, and hence X-ray is weak. More quantitative study requires spatial resolved X-ray and radio spectroscopy, which is beyond the scope of this paper.    \\subsection{G38.55+0.0}  On the contrary to the $ASCA$ results, no significant diffuse emission  is detected in the $XMM$-$Newton$  image of the G38.55+0.0 region. The absorbed surface brightness~(2.0--10.0~keV) of the GRXE is $\\sim 3.5\\times 10^{-11}~{\\rm ergs}~{\\rm cm}^{-2}~{\\rm s}^{-1}~{\\rm deg}^{-2}$ at $l=38.^{\\circ}55$ (Sugizaki et al. 2001). Assuming a thin thermal plasma of 7~keV for GRXE (Kaneda et al. 1997), the $XMM$-$Newton$ count  is estimated to be $\\sim 1.0$~counts~s$^{-1}$~deg$^{-2}$. The unabsorbed surface brightness of the CXB component (in 2.0--10.0~keV), on the other hand, is $\\sim 1.8\\times 10^{-11}~{\\rm ergs}~{\\rm cm}^{-2}~{\\rm s}^{-1}~{\\rm deg}^{-2}$ (Ishisaki 1996). With the Galactic absorption of $N_{\\rm H}=3.0\\times 10^{22}$~cm$^{-2}$, the count rate is estimated to be $\\sim 1.0$~counts~s$^{-1}$~deg$^{-2}$. Then the sum (GRXE+CXB) is $\\sim 2.0$~counts~s$^{-1}$~deg$^{-2}$. This value is consistent with  the best-fit result of the projected profile in table \\ref{tab:upper_limit} (2.3~counts~s$^{-1}$~deg$^{-2}$). Thus our flux estimation based on the model fitting to the projected profile is reliable. We then obtain the absorbed flux upper limit of the putative ``diffuse'' emission to be $6.4\\times 10^{-13}$~ergs~cm$^{-2}$~s$^{-1}$ (0.5--10 keV band), and the absorbed flux of the resolved point-source to be $2.6\\times 10^{-13}$~ergs~cm$^{-2}$~s$^{-1}$ in the same energy band. Then the sum of these fluxes is  $9.0\\times 10^{-13}$~ergs~cm$^{-2}$~s$^{-1}$, which is consistent with the $ASCA$ observed flux of G38.55+0.0. We note that no radio counterpart at the position of G38.55+0.0 has been reported.  Accordingly, whether G38.55+0.0 is diffuse source or a new SNR is still an open issue."
    },
    "0410/astro-ph0410529_arXiv.txt": {
        "abstract": "We report the discovery, from our preliminary survey of the Praesepe open cluster, of two new spectroscopically confirmed white dwarf candidate members.  We derive the effective temperatures and surface gravities of WD0837+218 and WD0837+185 (LB5959) to be $17845^{+555}_{-565}$K  and log g $= 8.48^{+0.07}_{-0.08}$ and $14170^{+1380}_{-1590}$K and log g $=8.46^{+0.15}_{-0.16}$ respectively. Using theoretical evolutionary tracks we estimate the masses and cooling ages of these white dwarfs to be $0.92\\pm0.05$M$_{\\odot}$ and $280^{+40}_{-30}$Myrs and $0.90\\pm0.10$M$_{\\odot}$ and $500^{+170}_{-100}$Myrs respectively. Adopting reasonable values for the cluster age we infer the progenitors of WD0837+218 and WD0837+185 had masses of $2.6\\le$ M $\\le $ M$_{\\rm crit}$M$_{\\odot}$ and $2.4\\le$ M $\\le 3.5$ M$_{\\odot}$ respectively, where M$_{\\rm crit}$ is the maximum mass of a white dwarf progenitor. We briefly discuss these findings in the context of the observed deficit of white dwarfs in open clusters and the initial mass-final-mass relationship. ",
        "introduction": "The common age, metallicity and distance of their members make galactic open star clusters favourable environments in which to examine fundamental issues in stellar and galactic astrophysics e.g. the shape of the initial mass function (IMF) or the form of the initial mass-final mass relationship (e.g. Weidemann 1987). The modestly rich and well studied  Praesepe (NGC2632) cluster at a distance of 177pc, as determined from Hipparcos astrometric measurements (Mermilliod et al. 1997), appears particularly suited to such investigations. Its members share a distinct proper motion so it is comparatively straightforward to discriminate them from the vast majority of field objects along this line of sight. For example, Hambly et al. (1995) performed an astrometric survey of 19 sq. degrees centered on the cluster and found the proper motions of members tightly clumped around $\\mu_{\\alpha} cos~\\delta = -30$ mas yr$^{-1}$ and $\\mu_{\\delta} = -8 $mas yr$^{-1}$. A more recent Hipparcos based study of Praesepe finds mean values of $\\mu_{\\alpha} cos~\\delta = -35.7$ mas yr$^{-1}$ and $\\mu_{\\delta} = -12.7 $mas yr$^{-1}$ (van Leeuwen 1999).  \\begin{table*} \\begin{minipage}{175mm} \\begin{center} \\caption{Details of the six new candidate WD members unearthed (top) and the four previously known WD members recovered (bottom) by our preliminary survey of Praesepe.} \\label{mass1} \\begin{tabular}{llcccccccr} \\hline Designation & Other ID  & RA & DEC & O & E & B$_{\\rm J}$ & R$_{\\rm F}$ &$\\mu_{\\alpha} cos~\\delta$ & $\\mu_{\\delta}$ \\\\ & & \\multicolumn{2}{|c|}{J2000.0} & \\multicolumn{2}{|c|}{USNO-B} &\\multicolumn{2}{|c|}{SuperCOSMOS} &\\multicolumn{2}{|c|}{mas yr$^{-1}$} \\\\ \\hline candidate 1 &  -         &  08 36 10.01 & 19 38 19.1 & 17.60 & 18.17  & 18.44& 18.30& -36$\\pm2$  & -2$\\pm3$  \\\\ WD0837+185 & LB5959     &  08 40 13.30 & 18 43 26.4 & 17.81 & 18.20  &18.28 & 18.24&  -32$\\pm0$  & -12$\\pm7$ \\\\ WD0837+218 &  -         &  08 40 31.47 & 21 40 43.1 & 17.87 & 18.64  &18.04 &18.17 &   -30$\\pm5$ & -6$\\pm5$  \\\\ candidate 4 &  -         &  08 42 58.03 & 18 54 35.5 & 17.86 & 18.44  & 18.34& 18.43&   -34$\\pm3$ & -8$\\pm2$  \\\\ candidate 5 &  -         &  08 43 22.00 & 20 43 30.5 & 18.34 & 18.39  & 18.44& 18.54&   -30$\\pm4$ & -8$\\pm2$  \\\\ candidate 6 & LB8648     &  08 46 01.91 & 18 30 48.5 & 18.04 & 18.25  & 18.11& 18.25&   -36$\\pm5$ & -14$\\pm5$ \\\\ \\hline WD0836+201 & LB393, EG61      &  08 39 45.57 & 20 00 16.0 & 17.54 & 18.24  & 17.98 & 18.15 &   -32$\\pm2$ & -10$\\pm4$ \\\\ WD0836+199 & LB1847, EG60     &  08 39 47.20 & 19 46 12.1 & 17.64 & 18.31  &18.73 &18.25 &   -38$\\pm3$ & -10$\\pm2$ \\\\ WD0837+199 & LB390, EG59      &  08 40 28.09 & 19 43 34.8 & 17.28 & 17.78  &17.48 &17.59 &   -34$\\pm2$ & -2$\\pm3$ \\\\ WD0840+200 & LB1876     &  08 42 52.32 & 19 51 11.3 & 17.27 & 18.00  &17.75 &18.02 &   -30$\\pm3$ & -10$\\pm4$  \\\\ \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table*}  A spectroscopic study of F type members indicates the cluster is slightly metal rich with respect to the Sun ([Fe/H]=0.038, [C/H]=0.01; Freil \\& Boesgaard 1992). This is consistent with the conclusions reached by previous investigations of this type (e.g. Boesgaard \\& Budge 1988). However, there is still uncertainty as to the age of the cluster, with estimates ranging from 0.4-2 Gyrs (e.g. Allen 1973; Mathieu \\& Mazeh 1988). Those determinations based on isochrone fitting generally support an age of between 0.7-1.1 Gyrs (e.g. Anthony-Twarog 1982; Mazzei \\& Pigatto 1988), although Claver et al. (2001), hereafter C01, favour a value closer to that of the Hyades (625Myrs), on grounds that the two clusters have similar metalicity and, kinematically, Praesepe is part of the Hyades moving group (Eggen 1960).  To date, five white dwarf members of Praesepe have been identified: LB390, LB5893, LB393, LB1847 and LB1876 (Luyten 1962; Eggen \\& Greenstein 1965; Anthony-Twarog 1982, 1984; C01). This is fewer than the 7-20 observable degenerates predicted from the extrapolation of the present day cluster luminosity function, allowing for reasonable assumptions about the form of the IMF, the maximum progenitor mass (M$_{\\rm crit}$) and the binary fraction (Williams 2004, hereafter W04). Several explanations have been put forward to account for this shortfall and the deficit of white dwarfs observed in other open clusters such as the Hyades. For example, if  M$_{\\rm crit}\\sim4$M$_{\\odot}$ there would have been fewer white dwarf progenitors in the first place (e.g. Tinsley 1974). However, the presence of the white dwarf LB1497 in the Pleiades with an estimated progenitor mass M$\\simgreat6$M$_{\\odot}$ argues against this (C01). It has recently been shown that asymmetry at a level of only 1\\% in the post main sequence mass loss process is sufficient to lead to the rapid loss of a significant fraction of the white dwarf population from an open cluster (Fellhauer et al. 2003). Alternatively, for Praesepe at least, it may be that no investigation to date has included a sufficient fraction of the total area the cluster projects on the sky. The surveys of Anthony-Twarog (1982, 1984) and C01 have both concentrated on the central $\\sim2$ sq. degrees of the cluster but Adams et al. (2002) determine the tidal radius to be $\\sim5^{\\circ}$ (see Figure 1).  There have been a number of studies of the five previously known Praesepe white dwarfs. Reid (1996) have used high resolution spectroscopy of the H$-\\alpha$ line cores to derive gravitational redshift based mass estimates of 0.42, 0.91 and 0.67M$_{\\odot}$ for LB390, LB5893 and LB393 respectively. C01 fit synthetic line profiles to moderate resolution, high S/N spectra of the H$-\\eta$ to H$-\\beta$ members of the Balmer series in each white dwarf, to measure effective temperatures (T$_{\\rm eff}$) and surface gravities (log g). Subsequently, they derived the mass of each, in the order listed in the above paragraph, to be 0.82, 0.91, 0.62, 0.82 and 0.75M$_{\\odot}$. They also noted the anomolously low mass Reid determined for EG59 likely stemmed from his neglect of the Zeeman splitting of the H$-\\alpha$ line core caused by a magnetic field with a strength of $\\sim3$MG. Reid and C01 have estimated the mass of the progenitor star of each degenerate, comparing the difference between the age of the cluster and the cooling time for the white dwarf to the predictions of stellar evolutionary models. For four of the white dwarfs studied the progenitor mass is consistent with the existence of a monotonically increasing relationship between the initial mass and the final mass (see Figure 11 of C01). However, in both investigations LB5893 is found to be ``too young'' for its comparatively high mass. C01 speculate that it may be the outcome of binary evolution, perhaps a double degenerate merger. Alternatively, taking into account both this white dwarf and his mass estimate for LB390, Reid suggests that a simple relationship between initial mass and final mass may not exist.  To move towards a resolution of these issues we are embarking on a comprehensive search for additional white dwarf members of Praesepe. Here we report the discovery, from a preliminary version of this survey, of two new white dwarf candidate cluster members. For each object we present an optical spectrum, determine T$_{\\rm eff}$ and log g and by comparing these measurements to evolutionary models estimate mass and cooling time. We conclude by briefly discussing our findings in the context of the reported deficit of white dwarfs in Praesepe and the initial mass-final mass relationship. \\vspace{0.1cm} \\section[]{A preliminary search for Praesepe white dwarfs}  We have utilised the USNO-B1.0 catalogue to undertake a survey of a $5^{\\circ} \\times 5^{\\circ}$ region centred on the Praesepe open cluster ($\\alpha=08~40$ $\\delta= +19~40$, J2000.0). The USNO-B catalogue contains astrometric information and photographic magnitudes for over a billion objects, gleaned from digitally scanned photographic plates spanning a baseline of $\\sim50$ years. The internal astrometric accuracy and the dispersion in the photometry are estimated to be $\\sim0.2''$ and $\\sim0.3$ magnitudes respectively (for details see Monet et al. 2003).  In this preliminary effort we have extracted all sources with $19 \\ge $ O $\\ge17$, O-E $ \\le 0$ and with proper motions $-25\\ge \\mu_{\\alpha} cos~\\delta \\ge -45 $ mas yr$^{-1}$, $0 \\ge \\mu_{\\delta} \\ge -20$ mas yr$^{-1}$. This encompasses the magnitude range of known cluster white dwarfs and is virtually coincident with the astrometric range sampled by Hambly et al. (1995). Further, the survey should be near complete for O-E $ \\simgreat -1$ (Hambly et al. 1995). Subsequently, the POSS II J and F images of each candidate have been inspected to eliminate extended sources, blended objects and spurious detections originating in the diffraction spikes of bright stars. As an additional check, candidates have been cross referenced against the 2MASS Point Source Catalogue (Skrutskie et al. 1997), keeping only those which are either non-detections or have blue near-IR colours within the photometric errors. Finally, we have used photographic photometry measured by SuperCOSMOS to compare the location of our new candidates to the locus of cluster white dwarfs in the B$_{\\rm J}$, R$_{\\rm F}$ colour-magnitude diagram (Figure 2). The external accuracy of individual passband magnitudes in the SuperCOSMOS Sky Survey is quoted as 0.3 magnitudes, but this uncertainty is dominated by drifts in zeropoints as a function of magnitude and position on the sky. These systematic errors do not appear when using colours as they are the same in all passbands, so indices like B$_{\\rm J} - $ R$_{\\rm F}$ are accurate to $\\sim0.1$ magnitudes (Hambly et al. 2001). This is not the case in USNO-B where the scatter in colours is $>0.3$ magnitudes.  To enhance the white dwarf sequence which is rather loosely defined by the five known degenerate members (the photometry for LB5893 appears to have been adversely affected by the proximity of the bright stellar cluster member KW195) we have used suitable objects drawn from the 20pc sample of Holberg et al. (2002) with trigonometric parallax determinations and SuperCOSMOS B$_{\\rm J}$ and R$_{\\rm F}$ photometry, scaling these to the cluster distance of 177pc. Of the seven new candidates the locations of six are deemed consistent with them being white dwarf members of Praesepe, the remaining object lying $\\sim1$ magnitude below the sequence (see Figure 2). Details of these six candidates and the four previously known white dwarf members recovered here are given in the top and the bottom of Table 1 respectively. Although C01 recognise LB5893 to be an astrometric member, it is not recovered here as its proper motion is listed in the USNO-B1.0 catalogue as $\\mu_{\\alpha} cos~\\delta=-56$ mas yr$^{-1}$, $\\mu_{\\delta} =-14$ mas yr$^{-1}$. This astrometry also appears to have been affected by the proximity of KW 195.   \\begin{figure} \\vspace{200pt} \\special{psfile=m44map.eps voffset=5 hscale=98 vscale=98 hoffset=22} \\caption{A schematic plot of the Praesepe cluster showing stars down to V $\\approx9$ and the areas surveyed by Anthony-Twarog (1982, 1984; solid outline) and C01 (grey shading). The region included in this investigation is outlined (dashed grey line). All objects listed in Table 1 (open cricles) and the known white dwarf cluster members (open stars) are also overplotted. The locations of the two new spectroscopically confirmed white dwarf candidate members are highlighted (open triangles).} \\end{figure}   \\begin{figure} \\vspace{180pt} \\special{psfile=m44cmd.eps voffset=14 hscale=85 vscale=85 hoffset=7} \\caption{A B$_{\\rm J}$, B$_{\\rm J} - $ R$_{\\rm F}$ colour-magnitude diagram of the eleven objects remaining after cross referencing against the 2MASS PSC (open circles). The thick line represents a linear least squares fit to SuperCOSMOS photometry of the known degenerate members (open circles+stars) and objects drawn from the 20pc sample of Holberg et al. (2002) with trigonometric parallax determinations (filled triangles). The magnitudes of the latter have been scaled to correspond to a distance of 177pc. Note that LB5893 (filled square) has been excluded from the fit.} \\end{figure}   We have obtained optical spectra of two candidates in the sample (LB5959 and WD0837+218) using the William Herschel Telescope and the double armed ISIS spectrograph on 28/01/2001. Sky conditions were fair on this night with clear skies but with seeing $\\sim2$''. For the course of the run ISIS was configured with the 5700 dichroic and the EEV12 and TEK4 detectors on the blue and red arms respectively. The data were obtained using the R158B and R158R gratings and a slit width of 1'' to provide a spectral resolution of $\\approx6$\\AA. Total exposure times were 60 and 120 minutes for LB5959 and WD0837+218 respectively. The CCD frames were bias subtracted, flat fielded and cosmic ray hits removed using the IRAF routines CCDPROC and FIXPIX. Subsequently the spectra were extracted using the APEXTRACT package and wavelength calibrated by comparison with the CuAr+CuNe arc spectra. Our observed spectral standards (G191-B2B and Feige 67) were drawn from the catalogues of Oke (1974,1990) and used to remove the instrument signature and telluric features from the science spectra.  \\begin{figure} \\vspace{128pt} \\special{psfile=tefflogg.eps voffset=5 hscale=70 vscale=70 hoffset=5} \\caption{The results of our fitting of synthetic profiles (thin black lines) to the observed Balmer lines (thick grey lines).} \\end{figure} ",
        "conclusions": "\\begin{table*} \\begin{minipage}{175mm} \\begin{center} \\caption{Progenitor lifetimes and corresponding masses for various adopted cluster ages. } \\label{pmass} \\begin{tabular}{lcccccc} \\hline Progenitor & \\multicolumn{2}{|c|}{8.80 (log$_{10}$ Myrs)}    &  \\multicolumn{2}{|c|}{8.92 (log$_{10}$ Myrs)}   & \\multicolumn{2}{|c|}{9.04 (log$_{10}$Myrs)} \\\\ of WD & \\multicolumn{1}{|c|}{$\\tau_{\\rm prog}$ (Myrs)} & \\multicolumn{1}{|c|}{M$_{\\rm prog}$ (M$_{\\odot}$)}  &  \\multicolumn{1}{|c|}{$\\tau_{\\rm prog}$ (Myrs)} & \\multicolumn{1}{|c|}{M$_{\\rm prog}$ (M$_{\\odot}$)} & \\multicolumn{1}{|c|}{$\\tau_{\\rm prog}$ (Myrs)} & \\multicolumn{1}{|c|}{M$_{\\rm prog}$ (M$_{\\odot}$)} \\\\ \\hline 0837+185     & $130^{+100}_{-170}$ & $4.9^{\\simless\\rm{M}_{\\rm crit}}_{-1.0}$ & $330^{+100}_{-170}$ & $3.4^{+1.1}_{-0.3}$  & $600^{+100}_{-170}$ & $2.8^{+0.3}_{-0.2}$  \\\\ 0837+218  & $350^{+30}_{-40}$ & $3.3^{+0.2}_{-0.1}$ &  $550^{+30}_{-40}$ & $2.9^{+0.0}_{-0.1}$ & $820^{+30}_{-40}$ & $2.5^{+0.0}_{-0.1}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table*}    \\subsection{Cluster membership and the white dwarf deficit}  A steep powerlaw ($\\Gamma=2$) was the only shape of IMF of the four investigated by W04 consistent with his assumption of seven white dwarf members of Praesepe, for reasonable values of the maximum progenitor mass (6M$_{\\odot}\\le$ M$_{\\rm crit}\\le10$M$_{\\odot}$). The ``Naylor'' IMF, a broken powerlaw with index $\\Gamma=0.2$ (M$\\le1$M$_{\\odot}$) and $\\Gamma=1.8$ (M$>1$M$_{\\odot}$) was found to be consistent only for M$_{\\rm crit}=6$M$_{\\odot}$. However, in a subsequent paper, Williams, Bolte \\& Liebert (2004), show LB6037 and LB6072 to be QSOs, reducing the number of bona fide white dwarf members of Praesepe unearthed to date to only five. In this case, only the steep powerlaw form with M$_{\\rm crit}\\le8$M$_{\\odot}$ can be considered consistent with the observations.  Nevertheless, we have presented evidence here the number of Praesepe white dwarfs is at least seven. As a further check of the membership status of our two new spectroscopic candidates we constrain their distances using the measured T$_{\\rm eff}$s and log g s and radii derived from evolutionary tracks. From the ``thick H layer'' models of Wood (1995) we determine R$=0.0093\\pm0.0020$R$_{\\odot}$ and R$=0.0091\\pm0.0009$R $_{\\odot}$ for LB5959 and WD0837+218 respectively. Subsequently,we estimate the absolute visual magnitude of LB5959 to be M$_{\\rm V} = 12.0\\pm0.5$  and that of WD0837+218 to be M$_{\\rm V} = 11.7\\pm0.2$. Refering to the synthetic photometry of Bergeron et al. (1995; private comm.), we estimate B-V $\\approx +0.15$ and +0.05 for the cooler and hotter degenerate respectively. Thus from the SuperCOSMOS data, B$_{\\rm J}=18.3\\pm0.3$ and B$_{\\rm J}=18.0\\pm0.3$, we determine V magnitudes of $18.15\\pm0.3$ and $18.0\\pm0.3$. Hence, we estimate these white dwarfs to reside at $170^{+45}_{-40}$pc and $180^{+40}_{-30}$pc respectively.  Indeed, the favourable success rate of our preliminary survey in unearthing cluster white dwarfs suggests several of the remaining four new candidates will also prove to be degenerate members. However, as confirmation of this must await further spectroscopy, for now we consider the observed number of Praesepe white dwarfs, N, to lie in the range $7\\le$ N $\\le11$. Based on the Williams simulations we find the steep powerlaw form of the IMF is consistent with the observed number for any reasonable value of M$_{\\rm crit}$. Similarly, if N$\\ge8$ the ``Naylor'' form can be considered consistent  for any reasonable value of M$_{\\rm crit}$. We note, one requires to detect at least ten white dwarf members for the Salpeter form of the IMF to be regarded at best, in Williams scheme, as mildly inconsistent with the observed number (P$\\approx0.07$). As our preliminary results indicate white dwarf members are to be found beyond the well studied inner regions of the cluster (Figure 1 shows WD0837+218 lies at a projected separation of $\\sim2^{\\circ}$) more likely await discovery. A detailed survey extending out to at least the tidal radius should be undertaken before any firm conclusions are drawn regarding the form of the IMF.  Prior to this work only the ``Naylor'' form of the IMF was found to be consistent with the non-detection of cluster white dwarfs residing in unresolved binary systems (W04). However, there is no evidence to suggest that either of the two new confirmed white dwarfs resides in a binary. We find a white dwarf with T$_{\\rm eff}=17000$K and log g $=8.0$, typical of the Praesepe population, has M$_{\\rm V}$=11.0, M$_{\\rm B}$=11.0 and  M$_{\\rm U}$=10.2 (Bergeron et al. 1995) and a young disc M3 dwarf has M$_{\\rm V}$=10.7, M$_{\\rm B}$=12.3 and M$_{\\rm U}$=13.4 (Leggett et al. 1992). An unresolved binary consisting of these two objects has U-B$=-0.6$, B-V$=0.6$. It thus lies on the fringes of the W04 criteria for being ``observable'' (U-B $\\le0.0$, B-V $\\le0.6$), which were chosen to approximate the limits of the UBV surveys of C01 and Anthony-Twarog. White dwarfs with more massive main sequence companions are unlikely to have been identified by these surveys. Similarly, since both Luyten's and the current work utilised blue and red photographic plates, selecting objects with colours I.C.$\\le0.2$ ($\\approx$ B-V$\\simless0.5$) and O-E$\\le0$ respectively, these too are likely biased against finding white dwarfs residing in binaries with stars of spectral type earlier than mid-M.  Farihi et al (2003) report a deficit of objects of spectral type later than mid-M paired to field white dwarfs at separations of $\\sim$ few 100AU. Further, radial velocity surveys point towards a drop in the relative frequency of main sequence binaries with mass ratios M$_{2}/$M$_{1}\\simless0.2$, at separations $\\simless5$ AU (Halbwachs et al. 2003; Marcy \\& Butler 2000). Hence, it seems plausible, particularly as the progenitors of the Praesepe and Hyades white dwarfs had M$\\simgreat2.5$M$_{\\odot}$, that by including a population of zero age binaries consisting of randomly paired stars, the simulations of W04 overpredict the number of detectable white dwarfs in unresolved binaries with main sequence companions. This assumption was acknowledged by W04 as a possible shortcoming in their modelling. Our forthcoming GALEX survey of Praesepe will expand the parameter space in which we are able to search for white dwarfs to include those in unresolved systems with K, G, and F type companions. It will reveal important additional information on the form of the IMF and the binary fraction of this cluster.  \\subsection{White dwarf masses and the initial mass-final mass relationship}  We have estimated the masses of the two new white dwarfs by comparing their measured T$_{\\rm eff}$s and log g s to the predictions of evolutionary calculations (Wood 1995). As expected for their progenitor masses, M $\\simgreat2.5$M, these two objects have masses greater than the canonical white dwarf value of M $\\approx 0.6$M$_{\\odot}$ (see Table 2). We have also used the evolutionary tracks to estimate the cooling times of these objects, determining $500^{+170}_{-100}$Myrs and $280^{+40}_{-30}$Myrs for LB5959 and WD0837+218 respectively. As the age of the cluster is rather uncertain we have used three different estimates encompassing the likely value (log $\\tau_{\\rm cluster} = 8.80, 8.92 ~{\\rm and}~ 9.04$) to derive the progenitor lifetimes.  Using cubic splines to interpolate between the lifetimes calculated for stars of solar composition by Girardi et al. (2000), we constrain the masses of the progenitors of LB5959 and WD0837+218 to be $2.6\\le$ M $\\le $ M$_{\\rm crit}$ M$_{\\odot}$ and $2.4\\le$ M $\\le 3.5$ M$_{\\odot}$ respectively.  Examining the location of these objects in Figure 11 of C01 we find while LB5959 fits in comfortably with a monotonic relationship between initial mass and final-mass, like LB5893, WD0837+218 appears to be too hot and hence too young for its high mass. We note there are five blue straggler members of Praesepe (e.g. Andrievsky 1998). The evolution of these objects appears to have been delayed, either through binary interaction or another as yet unidentified mechanism. We suggest that LB5893 and WD0837+218 may be related to this population. Alternatively, as suggested by Reid (1996), perhaps, at least some stars, do not subscribe to a simple monotonic relationship between their mass and the mass of their resulting white dwarf remnant."
    },
    "0410/astro-ph0410003_arXiv.txt": {
        "abstract": "Individual giant radio pulses (GRPs) from the Crab pulsar last only  a few microseconds. However, during that time they rank among the brightest objects in the radio sky reaching peak flux densities of up to 1500 Jy even at high radio frequencies. Our observations show that GRPs can be found in all phases of ordinary radio emission including the two high frequency components (HFCs) visible only between 5 and 9~GHz \\citep{moff96}. This leads us to believe that there is no difference in the emission mechanism of the main pulse (MP), inter pulse (IP) and HFCs. High resolution dynamic spectra from our recent observations of giant pulses with the Effelsberg telescope at a center frequency of 8.35~GHz show distinct spectral maxima within our observational bandwidth of 500~MHz for individual pulses. Their narrow band components appear to be brighter at  higher frequencies (8.6~GHz) than at  lower ones (8.1~GHz). Moreover, there is an evidence for spectral evolution within and between those structures. High frequency features occur earlier than low frequency ones. Strong plasma turbulence might be a feasible mechanism for the creation of the high energy densities of $\\sim 6.7 \\times 10^4~\\rm erg~cm^{-3}$ and brightness temperatures  of $\\sim 10^{31}~\\rm K$. ",
        "introduction": "\\label{obs}  Observations with the Effelsberg 100-m radio telescope began on 25 November 2003 and ended on  28 November 2003. We used a  secondary focus cooled  HEMT receiver with a center frequency of 8.35 GHz, providing two circularly polarized IF signals with a  system temperature of  25~K on both channels. With a sky temperature of 8~K, and a contribution of 33~K from the Crab nebula, the effective system temperature was 70~K. Two detection systems were used. First, a polarimeter with 1.1~GHz bandwidth detected total power of the left-hand and right-hand circularly polarized signals (LHC and RHC respectively) as well as $\\cos \\angle (\\rm~LHC, RHC)$ and $\\sin \\angle (\\rm~LHC, RHC)$. These four signals were then recorded by the standard {\\bf E}ffelsberg {\\bf P}ulsar {\\bf O}bservation {\\bf S}ystem ({\\bf EPOS}). With the Crab pulsar's dispersion measure of $56.8~{\\rm pc~cm^{-3}}$, we had an un--dedispersed time resolution of $t_{sample} = 890~\\mu{\\rm s}$, while our sampling resolution was fixed at $\\sim 640~\\rm \\mu s$. EPOS was therefore set to continuously record data blocks containing 20 periods divided into 1020 phase bins for all four signals (Fig.~1).  The minimum detectable flux per bin was $\\Delta S_{min} = 0.117~{\\rm Jy}$. Considering the dispersion pulse broadening, the detection limit for a single GRP of duration $\\tau_{grp} = 3~\\mu{\\rm s}$ would be $\\Delta S_{min}\\,t_{sample}/\\tau_{grp} = 25~{\\rm Jy}$. Selected GRPs are shown on Fig.~2.  A wide band ultra-high resolution detection system similar to the one of \\citet{Han2003} was used as well. The two IF signals (100--600~MHz) carrying the LHC and RHC information were sampled and  recorded with a fast digital storage oscilloscope (LeCroy LC584AL). $4 \\times 10^{6}$ samples per channel were recorded at the rate of $2 \\times 10^{9}$ samples per second. The storage scope was used in a single shot mode. We used giant radio pulses to trigger data acquisition by the digital storage scope. The RHC signal was detected, amplified and low-pass filtered (3~kHz), then passed to a second scope (Tektronics 744). The second scope triggered the digital scope whenever pulses stronger than 3--4 times rms (75--100~Jy) were detected. The recorded waveforms were then transferred to disk. This required manual intervention and led to a dead time of about 100 seconds after each pulse. ",
        "conclusions": ""
    },
    "0410/astro-ph0410145_arXiv.txt": {
        "abstract": "We construct an analytic model for the mass distributions of superclusters and clusters in each supercluster. Our model is a modification of the Press-Schechter theory, and defines superclusters as the regions that have some overdensity smaller than those for usual virialized objects. We compare the mass functions with a catalog of superclusters in the Sloan Digital Sky Survey Early Data Release and found that they are in reasonable agreement with each other. ",
        "introduction": "It is now known that our universe is hierarchical. The largest virialized objects in the universe are clusters that are made of hundreds of galaxies. Clusters have extensively been used to test structure formation models in the universe, through their abundances and correlations\\cite{bahcall83,bahcall92}.  Superclusters are defined as an ensemble of clusters. They are still assembling, and are even larger structure than clusters; superclusters have extents of $\\sim 100{\\rm Mpc}$ or even larger. Thus superclusters offer us information on late evolution of universe, and also on the transition from linear to non-linear regime. Examples of the superclusters include A901/902\\cite{gray02}, MS0302+17\\cite{fabricant94,gavazzi04}, Cl1604+43\\cite{lubin00,gal04}, the Sculptor supercluster\\cite{schuecker91}, and the Shapley supercluster\\cite{quintana95}. Detailed investigations have shown that these superclusters are real physical systems where clusters are gravitationally assembling. Detailed studies of these superclusters from galaxy distributions and weak lensing also showed that the mass distribution of superclusters is in good agreement with the distribution of early-type galaxies in superclusters. The physical state of gas in superclusters is still unknown, but it is worth mentioning that soft X-ray has been detected in the Sculptor supercluster\\cite{zappacosta04}.  Although most the studies have been limited to those of individual superclusters so far, the emergence of wide-field surveys/telescopes, such as the Sloan Digital Sky Survey (SDSS) and the Subaru Suprime-cam, begins to make it possible to do {\\it statistical} studies of superclusters. Nevertheless, there is no analytic model of superclusters that predicts, e.g., the mass distribution. In this paper, we construct such an analytic model of superclusters. Basically, our model is a modification of the Press-Schechter theory\\cite{press74}, and predicts not only mass distributions of superclusters but also number distribution of clusters in each supercluster. We also compare our model with superclusters in the SDSS extracted by Einasto {\\it et al.}\\cite{einasto03}. Throughout the paper we assume a flat universe $\\Omega_M+\\Omega_\\Lambda=1$, and the Hubble constant $h\\equiv H_0/(100{\\rm km\\,s^{-1}Mpc^{-1}})=0.7$. ",
        "conclusions": "In this paper, we have constructed an analytic model of superclusters. Our model is based on the Press-Schechter theory, and defines superclusters as overdense regions that are on the course of the spherical collapse. Our model allows us to compute the mass function of superclusters, and also the number distributions of clusters in each supercluster.  To test the model, we have compared the model with the supercluster catalog in the SDSS EDR data. We have shown that the model shows in reasonable agreement with the observation. We have also examined parameter dependences of the mass function of superclusters, and found that the mass function of massive superclusters is sensitive to $\\sigma_8$. Thus superclusters will be a useful probe of $\\sigma_8$ if superclusters are selected more rigorously. On the other hand, the number of clusters in superclusters shows only weak dependences on parameters.  Our model would be useful not only in constraining cosmological parameters from the abundance of superclusters but also in many applications. One of such applications is hot gas in superclusters. Recently it has been claimed that soft X-ray\\cite{zappacosta04} and Sunyaev-Zel'dovich effects\\cite{myers04} in the supercluster regions may be detected, and this might suggest the existence of relatively hot gas in superclusters. If this is true, the statistical significance of the signals that are caused by the hot gas can be computed by using our model, once the we measure the temperature and density structure of hot gas. Other application is gravitational lensing; gravitational lensing measures the total mass along the line of sight, thus a supercluster outside a cluster may affect, e.g., the mass estimate of the cluster. Our model is useful since it allows us to estimate the distribution of the host superclusters for a given cluster. In practice, we found that the effect is negligible when we consider superclusters with $\\Delta_{\\rm sc}\\sim 2$; assuming a homogeneous sphere, convergence of a supercluster with $M=10^{16}h^{-1}M_\\odot$ can be estimated as \\begin{equation} \\kappa=\\frac{\\Sigma}{\\Sigma_{\\rm crit}}\\sim \\frac{9\\times 10^{12}(M/10^{16}h^{-1}M_\\odot)^{1/3}(\\Delta_{\\rm sc}/2)^{2/3}hM_\\odot{\\rm Mpc^{-2}}}{4\\times 10^{15}hM_\\odot{\\rm Mpc^{-2}}}\\sim 2\\times10^{-3}, \\end{equation} where we assumed a source at $z\\sim 1$ and a supercluster at $z\\sim 0.3$. However, it is of interest to consider the effect more carefully, because the effect is dependent on the density profile of superclusters as well as the definition of superclusters, $\\Delta_{\\rm sc}$. In particular, superclusters could have much more significant effect if superclusters have centrally concentrated density distributions.  Although our model looks successful at this moment, definitely we need more tests. The more direct way to test the model is to compare it with numerical simulations, since we know the initial condition and cosmological parameters in numerical simulations. Of course, comparison with observations is also of great importance in order to test our understanding of the universe. Now the SDSS data have been becoming public\\cite{abazajian03,abazajian04}, and the quite large survey region of the SDSS would allow us to make three-dimensional catalog of superclusters. The follow-up work is currently in progress, and the results will be presented elsewhere."
    },
    "0410/astro-ph0410690_arXiv.txt": {
        "abstract": "We discuss the possibility to detect O/Ne white dwarfs by evidence for Ne overabundances. The hottest known WD, H1504+65, could be the only single WD for which we might be able to proof its O/Ne nature directly. Apart from this, strong Ne abundances are known or suspected only from binary systems, namely from a few novae, and from a handful of LMXBs.  We try to verify the hypothesis that the latter might host strongly ablated O/Ne WD donors, by abundance analyses of the accretion disks in these systems. In any case, to conclude on O/Ne WDs just by strong Ne overabundances is problematic, because Ne enrichment also occurs by settling of this species into the core of C/O WDs. ",
        "introduction": "Stellar evolution theory predicts the existence of massive O/Ne white dwarfs as a result of carbon burning. WDs with masses exceeding 1\\,M$_\\odot$ should be such O/Ne WDs, but there are other ways to produce massive WDs, namely by binary merging or by mass accretion during close-binary evolution. The question is: Is there evidence for the existence of O/Ne WDs? A very useful overview about this topic has been given by Weidemann (2003) at the last WD workshop at Naples. We refer to his paper for a detailed discussion and references concerning stellar evolution scenarios, which we only briefly summarize here.  O/Ne WDs should have masses of 1--1.3\\,M$_\\odot$. Below a core mass of 1\\,M$_\\odot$ C burning cannot ignite. If the core mass exceeds 1.3\\,M$_\\odot$, an electron capture collapse will occur and probably result in a neutron star. There is a debate as to what initial-mass range this 1--1.3\\,M$_\\odot$ core-mass range corresponds. Depending on details in evolutionary models, the range of initial masses is between 9--11\\,M$_\\odot$ or between 7--9\\,M$_\\odot$ (for solar metallicity). In any case, in order to stop the O/Ne-core to grow (by C burning) beyond the 1.3\\,M$_\\odot$ limit, the star must loose mass, either by wind mass loss or by envelope removal in close binary systems.  The WD mass distribution does show the existence of such massive  WDs, however, their origin is not clear. It is possible that they are binary mergers.  If O/Ne WDs exist, then they have C/O envelopes and, thus, they are spectroscopically indiscernible from C/O WDs. Possible evidence, however, are observed neon overabundances in several cases: in neon novae, in some low-mass X-ray binaries, and in one individual object (H1504+65). We will discuss here the LMXBs and H1504+65. In the case of neon novae, only those very few objects with extreme neon overabundances (over $\\approx$30 times solar, i.e.\\ 3\\% mass fraction) suggest O/Ne WDs, because up to 3\\% $^{22}$Ne is expected in C/O WDs.  For the interpretation of Ne overabundances in novae or LMXBs it is important to consider gravitational settling of Ne into the core of a C/O WD. This problem is not yet solved theoretically. It could well be, that Ne overabundances observed in eroded WD cores are the result of this process and not the result of C burning. ",
        "conclusions": ""
    },
    "0410/astro-ph0410373_arXiv.txt": {
        "abstract": "N-Body simulations are a very important tool in the study of formation of large scale structures. Much of the progress in understanding the physics of high redshift universe and comparison with observations would not have been possible without N-Body simulations. Given the importance of this tool, it is essential to understand its limitations as ignoring the limitations can easily lead to interesting but unreliable results. In this paper we study the limitations arising out of the finite size of simulation volume. This finite size implies that modes larger than the size of the simulation volume are ignored and a truncated power spectrum is simulated. If the simulation volume is large enough then the mass in collapsed haloes expected from the full power spectrum and from the truncated power spectrum should match. We propose a quantitative measure based on this approach that allows us to compute the minimum box size for an N-Body simulation. We find that the required box size for simulations of $\\Lambda$CDM model at high redshifts is much larger than is typically used. We can also use this approach to quantify the effect of perturbations at large scales for power law models and we find that if we fix the scale of non-linearity, the required box size becomes very large as the index becomes small. The appropriate box size computed using this approach is also an appropriate choice for the transition scale when tools like {\\sl MAP} \\cite{1996ApJ...472...14T} that add the contribution of the missing power are used. ",
        "introduction": "Large scale structures like galaxies and clusters of galaxies are believed to have formed by gravitational amplification of small perturbations \\cite{1980lssu.book.....P,1999coph.book.....P,2002tagc.book.....P,2002PhR...367....1B}. Observations suggest that the initial density perturbations were present at all scales that have been probed by observations. An essential part of the study of formation of galaxies and other large scale structures is thus the evolution of density perturbations for such initial conditions. The basic equations for this have been known for a long time \\cite{1974A&A....32..391P} and these equations are easy to solve when the amplitude of perturbations is small. Once the amplitude of perturbations at relevant scales becomes large, i.e., $\\delta \\sim 1$, the perturbation becomes non-linear and the coupling with perturbations at other scales cannot be ignored. The equation for evolution of density perturbations cannot be solved for generic perturbations in the non-linear regime. N-Body simulations \\cite{1998ARA&A..36..599B} are often used to study the evolution in this regime, unless one requires only a limited amount of information and quasi-linear approximation schemes \\cite{1970A&A.....5...84Z,1989MNRAS.236..385G,1992MNRAS.259..437M,1993ApJ...418..570B,1994MNRAS.266..227B,1995PhR...262....1S,1996ApJ...471....1H,2002PhR...367....1B} or scaling relations \\cite{1977ApJS...34..425D,1991ApJ...374L...1H,1995MNRAS.276L..25J,2000ApJ...531...17Ka,1998ApJ...508L...5M,1994MNRAS.271..976N,1996ApJ...466..604P,1994MNRAS.267.1020P,1996MNRAS.278L..29P,1996MNRAS.280L..19P,2003MNRAS.341.1311S} suffice.  In N-Body simulations, we simulate a representative region of the universe. This representative region is a large but finite volume and periodic boundary conditions are often used -- this is necessary considering the the universe does not have a boundary. Typically the simulation volume is taken to be a cube. Effect of perturbations at scales smaller than the mass resolution of the simulation, and of perturbations at scales larger than the box is ignored. Indeed, even perturbations at scales comparable to the box are under sampled. It has been shown that for gravitational dynamics in an expanding universe, perturbations at small scales do not influence collapse of large scale perturbations in a significant manner \\cite{1974A&A....32..391P,1985ApJ...297..350P,1991MNRAS.253..295L,1997MNRAS.286.1023B,1998ApJ...497..499C} as far as the correlation function or power spectrum at large scales are concerned. Therefore we may assume that ignoring perturbations at scales much smaller than the scales of interest does not affect results of N-Body simulations. However, there may be other effects that are not completely understood at the quantitative level \\cite{2004astro.ph..8429B}.  Perturbations at scales much larger than the simulation volume can affect the results of N-Body simulations. Use of periodic boundary conditions implies that the average density in the simulation box is same as the average density in the universe, in other words we are assuming that there are no perturbations at the scale of the simulation volume (or at larger scales). Therefore the size of the simulation volume should be chosen so that the amplitude of fluctuations at that scale (and at larger scales) is ignorable. If the amplitude of perturbations at larger scales is not ignorable and the simulations do not take the contribution of these scales into account then clearly the simulation will not be a faithful representation of the model being simulated. It is not obvious as to when fluctuations at larger scales can be considered ignorable. In this paper we will propose one method of quantifying this concept.  If the amplitude of density perturbations at the box scale is small but not ignorable, simulations underestimate the correlation function though the number density of small mass haloes does not change by much \\cite{1994ApJ...436..467G,1994ApJ...436..491G}. In other words, the formation of small haloes is not disturbed but their distribution is affected by non-inclusion of long wave modes. The mass function of massive haloes changes significantly \\cite{1994ApJ...436..467G,1994ApJ...436..491G}. The void spectrum is also affected by finite size of the simulation volume if perturbations at large scales are not ignorable \\cite{1992ApJ...393..415K}.  Significance of perturbations at large scales has been discussed in detail and a method (MAP) has been devised for incorporating the effects of these perturbations \\cite{1996ApJ...472...14T}. These methods make use of the fact that if the box size is chosen to be large enough then the contribution of larger scales can be incorporated by adding displacements due to the larger scales independently of the evolution of the system in an N-Body simulation. But this again brings up the issue of what is a large enough scale in any given model such that these methods can be used to add the effect of larger scales without introducing errors. Large scales contribute to displacements and velocities, and variations in density due to these scales modify the rate of growth for small scales perturbations \\cite{1997MNRAS.286...38C}. The MAP algorithm \\cite{1996ApJ...472...14T} can be improved by adding corrections for this effect \\cite{1997MNRAS.286...38C}. The motivation for developing these tools was to enlarge the dynamic range over which results of N-Body simulations are valid by adding corrections that change the distribution of matter and velocities at scales comparable to the simulation volume.   \\begin{table} \\begin{center} \\begin{tabular}{||l|l|l|l|l||} \\hline Name & $N_p$ & $L_{box}$ & $\\epsilon$ & Cutoff \\\\ \\hline T\\_300\\_C\\_0 & $256^3$ & $300$h$^{-1}$Mpc & $0.47$h$^{-1}$Mpc & No \\\\ T\\_300\\_C\\_2 & $256^3$ & $300$h$^{-1}$Mpc & $0.47$h$^{-1}$Mpc & $k \\leq 2 k_f$ \\\\ T\\_300\\_C\\_3 & $256^3$ & $300$h$^{-1}$Mpc & $0.47$h$^{-1}$Mpc & $k \\leq 3 k_f$ \\\\ T\\_300\\_C\\_4 & $256^3$ & $300$h$^{-1}$Mpc & $0.47$h$^{-1}$Mpc & $k \\leq 4 k_f$ \\\\ \\hline \\end{tabular} \\caption{This table lists parameters of N-Body simulations we have used.  All the simulations were done using the TreePM code, configuration described in Bagla and Ray (2003).  We simulated the $\\Lambda$CDM model with $\\Omega_B = 0.05$, $\\Omega_{dm}=0.3$, $k=0$, $h=0.7$ and $n=1$. The first column lists name of the simulation, second column lists the number of particles in the simulation, the third column is the size of the box in physical units, fourth column lists the softening length for force and the last column lists the cutoff used in units of the fundamental mode of the box $k_f$.  Phases were kept the same for the simulations listed here in order to make the comparison meaningful.} \\end{center} \\end{table}  Our motivation in this work is to understand the effect of large scales on scales that are much smaller than the simulation volume. We run a series of numerical experiments, these are described in \\S{2}. We discuss the proposed method of quantifying the effect of large scales in \\S{3}. Here, we propose a measure based on mass functions. The minimum size of simulation volume for $\\Lambda$CDM model as a function of redshift is presented here. ",
        "conclusions": "We have studied the influence of long wave modes on gravitational clustering at small scales. We find that for the $\\Lambda$CDM model, scales larger than $100$h$^{-1}$Mpc affect the mass function of haloes and the distribution of matter at scales as small as a few Mpc. The effect of long wave modes not present in an N-Body simulation can be incorporated independently of the evolution at small scales for some quantities but making such corrections is not possible in general. In particular, it is not possible to make corrections if the contribution of large scales changes the mass of collapsed haloes by a significant amount. We can turn this argument around and check whether a given size of the simulation volume will give (close to) correct results for total collapsed mass in haloes or not. This can be done using simple analytical formulae outlined in \\S{3}. The fractional deviation of collapsed mass from its expected value if density fluctuations at all scales are taken into account is a good indicator of the influence of large scales. The fractional deviation ($\\epsilon$) can be checked at a given mass scale $M$ of physical interest. Given a choice of $M$, $\\epsilon$ and the redshift $z$ at which output of the simulation is to be studied, we can compute the recommended minimum size for the simulation volume. It is important to note that this scale is the minimum required and a larger simulation may be warranted by other requirements.  The mass in collapsed haloes converges faster than the correlation function of haloes, implying that an even larger simulation volume may be required. Corrections to positions and velocities due to large scales can be made using tools such as MAP \\cite{1996ApJ...472...14T,1997MNRAS.286...38C}. Thus we can choose the box size by requiring convergence of mass in haloes and then use MAP to get the correct distribution and velocity field at all scales. Of course, if MAP is not being used then a larger box size than the minimum indicated by convergence of mass function may be required. Of course, using a large box size can make it difficult to address interesting questions using N-Body simulations. We believe that it is better to be cautious rather than obtain unreliable answers, unless one can make a convincing case that the relevant physical quantities are not as sensitive as mass in collapsed haloes to the size of the simulation box."
    },
    "0410/hep-ph0410250_arXiv.txt": {
        "abstract": "% We analyze the phase diagram of a three-flavor Nambu-Jona-Lasinio model at finite temperature $T$ and chemical potentials $\\mu_u, \\mu_d, \\mu_s$. We study the competition of pion and kaon condensation and we propose a physical situation in which kaon condensation could be led only by light quark finite densities. ",
        "introduction": "\\label{int} In the last few years the study of QCD at finite temperature and densities has been a subject of great interest. In particular, the regime of high temperatures and moderate densities is relevant for the physics of heavy ion collisions: the presence of a tricritical point, suggested in \\cite{Barducci:1989wi} is significant to determine the transition from a QGP phase to a hadronic phase \\cite{Jakovac:2004pu}.  Moreover, the presence of an isospin chemical potential, namely of an asymmetry in the densities of light quarks, could complicate this picture: for values of $\\mu_I=(\\mu_u-\\mu_d)/2$ attainable in the experiments, it is expected that the critical lines associated to the light flavors, and therefore the ending points, are mutually shifted. The presence of two critical lines would make the transition smoother and would produce observable effects on physical quantities across the transition \\cite{Toublan:2003tt,Toublan:2004ks}.   On the other hand, the regime of low temperature and high densities is relevant for compact stars. The existence of stars whose core is so dense that quarks can be treated as fundamental degrees of freedom has been considered for instance in \\cite{Steiner:2000bi}: there, the temperature is small enough to safely consider the limit $T\\rightarrow0$. In this context, the physical scenario is very rich and interesting. For quark densities higher than $300\\div400$ MeV, the presence of superconductive phases has been established  \\cite{Rajagopal:2000wf}: there, the mass of the strange quark plays the role of distinguishing the pattern of symmetry breaking between a two flavor superconductive (2SC) and a three flavor color flavor locking (CFL) phase, in both cases the gap being of the order of a few tens of MeV. For this reason, an exhaustive study of these phenomena must take into account the strange quark sector. It is also worth to mention that taking into account weak equilibrium and color and electric neutrality many other interesting superconductive phases could exist \\cite{Alford:2000ze,Neumann:2002jm,Alford:2003fq,Shovkovy:2003uu}.  The study of finite density on lattice has well known problems and except for the attempt of Fodor and Katz in ref. \\cite{Fodor:2001pe} the only studied case is the one corresponding to low densities \\cite{deForcrand:2002ci}. For this reason, the only tool available at this moment for the study of the phase diagram in the range from medium to high densities is a class of effective models: the stability of the results obtained with different approaches must be checked to give predictivity to the analyses.     In any case, at zero temperature for values of $\\mu_I\\geq m_{\\pi}/2$ a pion condensed phase should appear, as has been first suggested by \\cite{Son:2000xc} in the framework of a chiral model: this phase is confirmed by lattice analyses \\cite{Kogut:2002tm}, directly investigable since for $\\mu_I\\neq0$ the fermion determinant is positive, within chiral models \\cite{Splittorff:2000mm} and within microscopical models \\cite{Toublan:2003tt,Klein:2003fy,Frank:2003ve,Barducci:2003un,Barducci:2004tt}. Our previous study in the NJL model \\cite{Barducci:2004tt} and that in ref. \\cite{Frank:2003ve}, together with the one in random matrix \\cite{Klein:2003fy}, showed the effect of both isospin and baryon chemical potentials on the structure of the QCD phase diagram.  By extending this analysis to the third flavor, a kaon condensed phase is expected in the region of high strange quark chemical potential $\\mu_s\\geq m_K$, as shown by ref \\cite{Kogut:2001id}. The latter analysis being realized in the context of chiral models at zero temperature, it would be interesting to study the competition between pion and kaon condensation in the framework of a microscopical model, for a general configuration in the space of thermodynamical parameters $(T,\\mu_u,\\mu_d,\\mu_s)$.   Moreover, it has been suggested a kaon condensed phase in the context of high density nuclear matter \\cite{Kaplan:1986yq} and, more recently, in particular in the CFL regime \\cite{Bedaque:2001je}: since we don't yet know whether the densities attainable in the core of a neutron star can be high enough to favor a superconductive or a baryon rich phase, we study here the possibility of kaon condensation in a regime where baryons are still present, and where the densities are associated only with non strange quark matter. Also, for the first time we perform the study of kaon condensation by using a microscopical model.   To summarize, in this paper we study pion and kaon condensation at finite temperature and quark densities, by using a three flavor non local NJL model, and we propose a physical scenario in which kaon condensation could be driven only by a light quark finite density. The possibility of a kaon condensation at $\\mu_s=0$ could be important in the context of neutron stars, where, by neglecting electroweak effects, the densities are associated only to light quarks. To this end, we generalize our analyses previously performed in the two flavor sector \\cite{Barducci:2004tt,Barducci:2003un}. In this paper, working at intermediate densities, we do not consider the possibility of diquark condensation, and we turn off the electroweak effects. ",
        "conclusions": "In this paper we have presented a calculation of the QCD phase diagram obtained by using a three-flavor NJL model. We have considered the plane ($\\mu_I,\\mu_Y$) to study the competition of kaon and pion condensation at zero and finite temperature and to establish a comparison with a previous analysis performed in a chiral model \\cite{Kogut:2001id}. The results of the microscopical model accurately reproduce those achieved in the effective model only for low densities, since in the high strangeness density regime approximate chiral symmetry restoration (relative to the strange sector) occurs, disfavoring kaon with respect to pion condensation. Moreover we have suggested the possibility of kaon condensation driven only by light quark densities: this situation could be of interest in the core of compact stars, where the densities are approximatively associated only to light flavors. Up to now, we have neglected the possibility of di-quark condensation: however, we are currently working on an extension of this analysis that takes into account the superconductive phases too, so as to reproduce the physical picture proposed in the context of effective models for the K0-CFL phase \\cite{Bedaque:2001je}, and to understand the behaviour of color superconductivity by varying arbitrarily the three quark chemical potentials \\cite{Bedaque:1999nu}. Eventually, the inclusion of the $\\beta$-equilibrium as well as charge and color neutrality, will give us a reasonable picture of cold quark matter within stars, and will make possible a comparison with experimental data."
    },
    "0410/astro-ph0410415_arXiv.txt": {
        "abstract": "{We study the presence of low intensity high velocity components, which we have termed {\\it wing} features in the integrated \\ha\\ emission line profiles of the \\hii\\ region populations of the spiral barred galaxies NGC~1530, NGC~3359 and NGC~6951. We find that more than a third of the \\hii\\ region line profiles in each galaxy show these components. The highest fraction is obtained in the galaxy whose line profiles show the best S:N, which suggests that wing features of this type may well exist in most, if not all, \\hii\\ region line profiles. Applying selection criteria to the wing features, we obtain a sample of \\hii\\ regions with clearly defined high velocity components in their profiles. Deconvolution of a representative sample of the line profiles eliminates any doubt that the wing features could possibly be due to instrumental effects. We present an analysis of the high velocity low intensity features fitting them with Gaussian functions; the emission measures, central ve\\-lo\\-ci\\-ties and velocity dispersions for the red and blue features take similar values. We interpret the features as signatures of expanding shells inside the \\hii\\ regions. Up to a shell radius of $\\rm R_{shell}\\sim 0.2~R_{reg}$, the stellar winds from the central ionizing stars appear to satisfy the energy and momentum requirements for the formation and driving the shell. Several examples of the most luminous \\hii\\ regions show that the shells appear to have larger radii; in these cases additional mechanisms may well be needed to explain the kinetic energies and momenta of the shells. ",
        "introduction": "The study of the internal kinematics of \\hii\\ regions is quite di\\-ffi\\-cult technically and has proved to be a far from easy task to achieve. The gas motions may be observed in the velocity profiles of the emission lines from the \\hii\\ regions, which show either a broad component, typically with supersonic width, or several components. Previous studies have shown evidence of low intensity broad components in the integrated spectra of the \\hii\\ regions, not only in Galactic regions but also in those in external galaxies. Arsenault \\& Roy (1986) found that a significant fraction of their extragalactic \\hii\\ region sample showed integrated line profiles which are better characterized by a Voigt profile rather than by a single Gaussian profile; while Chu \\& Kennicutt (1994) and Yang et al. (1996) found the same broad component for nearby extragalactic \\hii\\ regions. Up to now the interpretation of these extended components has not been stu\\-died for several reasons; the low intensity of the broad components makes it difficult to confirm their existence, and because of this, examples of \\hii\\ regions showing clearly this type of components have been scarce until now.  The kinematics of nearby extragalactic \\hii\\ regions, notably 30~Doradus in the LMC, and NGC~604 in M33, have been carefully studied and have shown the existence of rapid gas motions in their interiors. Smith \\& Weedman (1972) first studied the internal kinematics of 30~Doradus with Fabry-P\\'erot observations and found asymmetric profiles at different positions within the region; the asymmetries could be matched by fitting the line profile with several Gaussian components. Chu \\& Kennicutt (1994) studied the structural kinematics of 30~Doradus using several long--slit spectra to map the velocity field of the region. They concluded that the velocity field is dominated by a large number of expanding structures ranging in size and expansion velocities, most of them confined within the central zone of radius $\\rm R\\sim 65$~pc. In spite of the complex kinematic structure, the integrated velocity profile for the central part of the region, derived by adding the profiles for all the slits, shows a broad Gaussian high intensity component with faint extended wings.  Yang et al. (1996) made a similar study of NGC~604 combining CCD ima\\-ging with long--slit and Fabry-P\\'erot observations. They confirmed the existence of several expanding shells located in the central part of the region, as in 30~Doradus. The total \\ha\\ luminosity coming from these shells is $\\sim$25\\% of the total luminosity of the \\hii\\ region, and the integrated ve\\-lo\\-ci\\-ty profile of the region shows similar faint extended wings to the profile of 30~Doradus. Other stu\\-dies (Rosa \\& Solf (1984); Sabalisck et al. (1995)) have found the same kinematic structure of expanding shells for NGC~604.  An important question not addressed till now is whether this phenomenon can be generalized to all \\hii\\ regions. Up to the present time the spatial and spectral resolution of the observations made of \\hii\\ regions in galaxies outside the Local Group have not been sufficient to allow us to obtain an answer to this question. It is necessary to obtain a significant sample of \\hii\\ region line profiles to make general inferences. However, the similarities between 30~Doradus and NGC~604 suggest that, in general, \\hii\\ regions could well have comparable structures and internal kinematics.   In order to overcome these observational limitations, we present an analysis of the line profiles of the \\hii\\ region po\\-pu\\-la\\-tions of the barred spiral galaxies NGC~1530, NGC~3359 and NGC~6951 obtained using Fabry-P\\'erot interferometry. We confirm the existence of peaked high velocity low intensity features in a sig\\-ni\\-fi\\-cant fraction of the line profiles and analyze a selected sample of \\hii\\ regions to study further the nature of these extended wings. In providing a statistical study of the phenomenology of the wings we need to pay a very careful atten\\-tion to eliminate doubts about the nature of the wings. For selected examples, we deconvolved the line profiles to obtain the best fit to the input spectra and to verify the existence of the low intensity components.   We interpret the wing features as signatures of an ex\\-pan\\-ding shell within a given \\hii\\ region. The energetics of the shell and the energy input requirements to produce it are studied for each \\hii\\ region of a selected sample making use of representative physical parameters measured in the most luminous shells observed in the nearby \\hii\\ regions, 30~Doradus and NGC~604. Finally, we have studied quantitatively the hypothesis that the expanding shells are produced by the interaction of the stellar winds coming from the ionizing stars with the ISM medium. ",
        "conclusions": "\\begin{itemize} \\item We have studied the existence of low intensity high ve\\-lo\\-ci\\-ty components in the integrated line profiles of the \\hii\\ region populations of three barred spiral galaxies. We have analyzed the integrated line profiles and extracted a representative sample of spectra which show these components. Down to a limit of log~L$_{\\scriptsize\\ha}=38.0$~(\\ergs), we found that 26.77\\%, 42.0\\% and 31.63\\% of the catalogued \\hii\\ regions in NGC~3359, NGC~6951 and NGC~1530, respectively, show high velocity low intensity features in their integrated line profiles.  \\item The higher fraction is correlated with a higher S:N ratio in the line profiles, which suggests that at sufficiently high value of S:N all \\hii\\ regions may well show these wing features, which means that these high velocity features could well exist in most \\hii\\ region line profiles, and as such provide evi\\-den\\-ce of a widely distributed phenomenon.  \\item Based on the mean S:N of the two high velocity components, we have obtained the spectra from the set of line profiles with increasing apertures that best defines the wing features and used them as the representative spectra of each region which characterize the wing parameters.  \\item The analysis of the fit parameters of the high velocity low intensity features show that parameters of the red and blue component are in the same range of values. The mean va\\-lues of the Emission Measure for both components are 8.9\\% of the total EM of the \\hii\\ region and the velocity separation ranges from 40~\\kms\\ to 90~\\kms\\ for our sample of the most luminous \\hii\\ regions (which cover the high of the luminosity range), with a mean value for each component of $<\\rm v_{red}>=60.5$~\\kms\\ and $<\\rm v_{blue}>=59.02$~\\kms. The mean measured velocity dispersions for the red and blue components are $<\\rm\\sigma_{red}>=19.5$~\\kms\\ and $<\\rm \\sigma_{blue}>=17.0$~\\kms, though some line profiles are not fully resolved.  \\item We interpret the high velocity features as evidence of an ex\\-pan\\-ding shell inside the \\hii\\ region. The shell expansion velocity is obtained as the mean value of $\\rm v_{red}$ and $\\rm v_{blue}$ for each \\hii\\ region. Based on parameters for the highest luminosity shells observed in the nearby extragalactic \\hii\\ regions, 30~Doradus and NGC~604, we have estimated the physical parameters and the energetics of the shell of each \\hii\\ region in our sample.  \\item We have compared the shell kinetic energy with the kinetic ener\\-gy of the stellar winds coming from the ionizing stars inside the \\hii\\ region. For shells with radii $\\rm\\sim 0.2~R_{reg}$, the ratios we derive between the kinetic energy of the shell and the energies of the stellar winds are rather higher than the expected theoretical ratio 0.2, but the uncertainties in obtaining a reliable value for the wind energy input from the observed \\ha\\ luminosity, as well as those inherent in our derivation of the observational values, can well explain the differences. Thus, the ratios obtained imply that ste\\-llar winds alone could drive the expanding shells, if these are located within $\\rm\\sim 0.2~R_{reg}$.  \\item Some of the most luminous \\hii\\ regions of our sample show evidence that their expanding shells are located at radii higher than $\\rm\\sim 0.5~R_{reg}$. For these \\hii\\ regions the stellar wind mechanism alone may well not be able to explain the observed shell kinetic energies or momenta, implying that there are other mechanisms required to account for the observed properties.   \\end{itemize}"
    },
    "0410/astro-ph0410623_arXiv.txt": {
        "abstract": "{We present some results from an archival {\\it VLA} study of ultraluminous X-ray sources (ULXs). These unresolved non-nuclear X-ray sources have luminosities ($L_X \\geq 10^{39}$ ergs s$^{-1}$) which may require somewhat exotic explanations, such as intermediate mass black holes or super-Eddington accretion. Radio emission is a powerful way to investigate such sources, through radio morphology and the implications made by source energy and lifetimes derived from the radio. The three galaxies we present here suggest that no single model explains all ULXs, yet there is growing evidence that some ULXs are powered by intermediate mass black holes. } ",
        "introduction": "Sensitive, high-resolution X-ray observatories such as {\\it Chandra} are identifying numerous unusually strong X-ray sources in nearby galaxies. These sources have X-ray luminosities in excess of X-ray binaries (which have $L_X < \\sim10^{38}$ ergs s$^{-1}$), yet fall short of active galactic nuclei (which have $L_X > \\sim10^{42}$ ergs s$^{-1}$). What is most intriguing about this luminosity range is that it is above the Eddington limit for spherical accretion onto the end products of normal stellar evolution, yet such sources often do not lie near the nuclei of their host galaxies. Consequently, they are generally called ``intermediate X-ray objects'' (IXO) or more frequently ``ultraluminous X-ray sources'' (ULX). They are generally defined as non-nuclear sources with $L_X \\geq 10^{39}$ ergs s$^{-1}$.  What are these enigmatic sources? Multiple theories exist, including the strong possibility that no single model can explain all ULXs. First, ULXs may simply be chance superpositions of sources not physically associated with the apparent host galaxies (mainly background AGN). However, even in the days of {\\it ROSAT} it was realized that there are simply too many candidate ULXs to be explained by chance superpositions of background sources \\cite{fabbiano89}. Another less exotic potential progenitor is strong young supernova remnants (SNR), which can reach $L_X \\sim 10^{40}$ ergs s$^{-1}$. The main weakness of this explanation is that the X-ray emission from ULXs is generally too variable to be explained as one or more SNR. In fact, the X-ray spectra of ULXs are usually consistent with X-ray binaries (XRBs), including those with a black hole primary \\citep[e.g.,][]{makishima00}. The question then becomes ``What is the mass of the black hole?'' The importance of this question lies in the distinction between stellar mass black holes and intermediate mass black holes \\citep[IMBHs, of $10^2 - 10^4$ M$_\\odot$, e.g.][]{colbert99}. ULXs might arise from the former should their emission be beamed, either mechanically \\citep[i.e., anisotropic emission due to the accretion disk;][]{king01} or relativistically \\citep[akin to microquasars;][]{kording02}. This beaming would alleviate any concern over super-Eddington luminosities \\citep[additionally, models with super-Eddington emission from thin accetion disks have been proposed;][]{begelman02}. Obviously, should the accreting object be an IMBH there is no violation of the Eddington limit and the sources can radiate isotropically. This model is attractive as it could fill in the black hole mass function between the stellar mass black holes which arise from normal stellar evolution and the supermassive black holes black holes found at the centers of galaxies. The tradeoff is that we would then require an explanation for how such compact objects formed, although models do indicate they might form via mergers in compact stellar clusters \\citep{miller02}.  Radio observations represent a powerful yet often overlooked tool toward evaluating ULXs. The long wavelengths of radio emission are essentially unaffected by the dust extinction which can complicate optical searches, and there are numerous baseline expectations for the relative proportions of radio and X-ray flux for known objects \\citep[e.g.,][]{neff03}. At a very simple level, the radio morphology of any detected ULX strongly discriminates between beamed stellar mass XRBs and isotropic emission from an IMBH: unresolved radio emission is consistent with beaming towards the observer, whereas resolved emission suggests any beaming is not along the line of sight (assuming the radio and X-ray emission originate from the same source). Along these lines, \\citet{kaaret03} used an unresolved ATCA detection of the ULX in NGC~5408 to argue for a beamed progenitor. Lastly, radio fluxes and spectral indices can be used to evaluate the emission mechanism (e.g., bremsstrahlung, synchrotron) and source total energy and lifetime. These are important clues toward understanding the nature of ULXs.  We are involved in a radio study of known ULXs which relies on archival {\\it VLA} data. In particular, we rely on archival data which combines fairly lengthy integration times with higher resolution arrays (e.g., A array at 4.86~GHz). The study is largely blind in that we produce radio images of galaxies hosting ULXs and then compare these with the locations of X-ray point sources identified by {\\it Chandra} and {\\it XMM}. To date, our findings include a number of nuclear radio and X-ray detections (presumably low luminosity AGN) and several non-nuclear detections. We present here the results for three galaxies which exemplify the diversity of intrinsic non-nuclear source types: Holmberg~II, NGC~3877, and NGC~4631. ",
        "conclusions": "The results for these three galaxies are illustrative of some of the difficulties in understanding ULXs. In general, no simple picture can fit all four sources and this suggests that ULXs are a mixed bag of source types. The ULX in Holmberg~II is quite enticing, with extended radio emission arguing against beamed progenitors and suggesting IMBHs as the root of ULXs. However, the ULX in NGC~5408 has a comparable $R_X$ value yet is apparently compact based on the ATCA radio data described in \\citet{kaaret03}. We also have one clear SNR (SN1998S in NGC~3877), with the other source in NGC~3877 presenting some confusion. Although this second source is offset from the optical nucleus by a few arcseconds, its $R_X$ places it between expectations for low luminosity AGN and other radio-detected ULXs. Thus, it may represent a low luminosity AGN \\citep[similar to the pair of low luminosity AGN in the merger NGC~3256;][]{neff03} or a ULX similar to the one in Holmberg~II but whose X-ray observation happened to occur at a time of unusually low luminosity. Lastly, the radio-detected X-ray source in NGC~4631 appears to be a background AGN.  Finally, we stress that the variety of intrinsic source types underscores the utility of radio observations. The radio data greatly aid in the characterization of the candidate ULXs, and continuation of such studies should lead to significant numbers of the various source types and ultimately to an answer to the question of what powers these mysterious objects."
    },
    "0410/astro-ph0410309_arXiv.txt": {
        "abstract": "We study analytically and numerically the evolution of bound systems in universes with accelerating expansion where the acceleration either increases with time towards a Big Rip singularity (phantom cosmologies) or decreases with time (quintessence). We confirm the finding of Caldwell et. al. \\cite{Caldwell:2003vq} that bound structures get dissociated in phantom cosmologies but we demonstrate that this happens earlier than anticipated in Ref. \\cite{Caldwell:2003vq}. In particular we find that the `rip time' when a bound system gets unbounded is not the time when the repulsive phantom energy gravitational potential due to the average $(\\rho + 3p)$ balances the attractive gravitational potential of the mass M of the system. Instead, the `rip time' is the time when the minimum of the time dependent effective potential (including the centrifugal term) disappears. For the Milky Way galaxy this happens approximately 180Myrs before the Big Rip singularity instead of approximately 60Myrs indicated in \\cite{Caldwell:2003vq} for a phantom cosmology with w=-1.5. A numerical reconstruction of the dissociating bound orbits is presented. ",
        "introduction": "There is mounting observational evidence that the universe has entered a phase of accelerating expansion (the scale factor obeys ${\\ddot a}>0$) and that the total amount of clustered matter in the universe is not sufficient for its small average spatial curvature. This converging observational evidence comes from a diverse set of cosmological data which includes observations of type Ia supernovae \\cite{snobs}, large scale redshift surveys \\cite{lss} and measurements of the cosmic microwave background (CMB) temperature fluctuations spectrum \\cite{cmb}. The observed accelerating expansion and flatness of the universe, requires either a modified theory of gravity\\cite{modgrav} or, in the context of standard general relativity, the existence of a smooth energy component with negative pressure termed `dark energy'\\cite{dark energy}. This component is usually described by an equation of state parameter $w\\equiv{p\\over \\rho}$ (the ratio of the homogeneous dark energy pressure $p$ over the energy density $\\rho$). For cosmic acceleration, a value of $w<-{1\\over 3}$ is required as indicated by the Friedmann equation \\be {{\\ddot a}\\over a}=-{{4\\pi G}\\over 3}(\\rho +3p) \\label{fried}\\ee Current observational bounds \\cite{snobs,phant-obs2} on the value of the dark energy equation of state parameter $w(t_0)$ at the present time $t_0$ yield \\be -1.48<w(t_0)<-0.72 \\ee at the $95 \\%$ confidence level. The role of dark energy can be played by any physical field with positive energy and negative pressure which violates the strong energy condition $\\rho +3p>0$ ($w>-{1\\over 3}$).  Quintessence scalar fields\\cite{quintess} ($-1<w<-{1\\over 3}$) violate the strong energy condition but not the dominant energy condition $\\rho + p >0$. Their energy density scales down with the cosmic expansion and so does the cosmic acceleration rate. Phantom fields\\cite{Caldwell:1999ew,Johri:2003rh,phantom} ($w<-1$) violate the strong energy condition, the dominated energy condition and maybe physically unstable. However, they are also consistent with current cosmological data and according to recent studies\\cite{phant-obs2} they are favored over their quintessence counterparts. In contrast to quintessence fields, the energy density of phantom fields increases with time and so does the predicted expansion acceleration rate ${{\\ddot a}\\over a}$. This monotonically increasing acceleration rate of the expansion may be shown to lead to a novel kind of singularity which occurs at a finite future time and is characterized by divergences of the scale factor $a$, the Hubble parameter $H$ its derivative ${\\dot H}$ and the scalar curvature. This singularity has been called `Big Smash' \\cite{McInnes:2001zw} the first time it was discussed and `Big Rip' \\cite{Caldwell:2003vq} (hereafter CKW) in a more recent study. Even though there are mechanisms by which the `Big Rip' singularity could be avoided \\cite{Nojiri:2004ip} it remains an interesting possible fate of the universe.  An immediate consequence of the very rapid expansion rate as the Big Rip singularity is approached is the dissociation of bound systems due to the buildup of repulsive negative pressure in the interior of these systems. This observation was first made in \\cite{Caldwell:2003vq} where a qualitative study of the dissociation times for various bound systems was also made.  The quantitative study however of the evolution of a bound system in an expanding universe remains an issue of current research. In particular the question of whether the expansion of the universe affects gravitationally bound systems like clusters, galaxies or planetary systems has been addressed in several early \\cite{ES,BS,Cal,Peeb} and recent \\cite{Baker:2001yc,Cooperstock:1998ny,Bon} studies. The recurrent attention paid to this issue indicates that to this point a definitive answer is still lacking. The prevalent perception however is that the physics of systems which are small compared to the radius of curvature of the cosmological background is essentially unaffected by the expansion of the universe. Even though this perception is applicable in the case of phantom cosmology, it is not applicable as the Big Rip singularity is approached and curvature increases rapidly.  A qualitative approach to the problem of bound system dissociation in phantom cosmologies was made by CKW where it was assumed that a bound system becomes unbound when the source of the repulsive phantom energy gravitational potential for a bound system of size $R$ $(-{4\\over 3}\\pi (\\rho +3 p)R^3)$ balances the attractive source of the ordinary gravitational potential (the mass $M$ of the system). Even though this assumption is qualitatively useful as it provides a physical understanding as to why does a bound system dissociate due to phantom  energy repulsion, quantitatively it leads to incorrect estimates of the dissociation times because it ignores the effects of the centrifugal barrier in the effective potential of bound systems. A quantitative analysis should make use of the particle equations of motion in the local inertial frame  based on the geodesic deviation equation. Using such an equation, the time dependent effective potential of two body bound systems may be shown to consist of three terms, the static attractive gravitational mass term, the static repulsive centrifugal term and the time-dependent repulsive dark energy term. The stability of a bound system depends on the existence of a minimum for the above effective potential. In the case of quintessence $(w>-1)$ the minimum of the effective potential can not disappear at any future time because the repulsive time-dependent term decreases with time. For phantom energy however $(w<-1)$ the repulsive time-dependent term increases with time and at some critical time dominates over the other two terms and destroys the minimum of the effective potential dissociating at the same time the corresponding bound system. It should be stressed that this critical time when the effective potential minimum disappears is not the time when the repulsive dark energy term balances the attractive gravitational mass term as assumed in CKW. This will be demonstrated in detail, analytically and numerically in the following sections.  The structure of this paper is the following: In section 2 we use a metric that interpolates between the static Schwarzschild at small scales and a general time dependent Friedmann universe metric at large scales to derive the geodesics of a test particle in the Newtonian limit. We then focus on the particular Friedmann universe that contains a two component cosmic fluid (matter and dark energy) and derive the form of the scale factor for any value of constant $w$ (the equation of state parameter). Using this form of the scale factor in the previously derived geodesic equation we find the equation of motion for a two body bound system in an expanding quintessence or phantom cosmological setup. In section 3 we study this equation of motion and derive the evolution of the radius of two body bound systems with distance and mass scales corresponding to the Solar System, the Milky Way Galaxy and the Coma Cluster. We also derive analytically the dissociation time as a function of $w$ for a bound system of given mass and radius and test our result using numerical simulations of the above three types of bound systems (Solar System, Milky Way and Coma Cluster). Our analytical result for the time difference `Big Rip time - Dissociation time' differs from the corresponding result of CKW by a factor of 3.08 (we find the dissociation to occur earlier). The time evolution of the bound system effective potential is also shown. Our results are extended to the case of planar mass sources and to the case of quintessence $(w>-1)$. In the later case it is shown that the size change of the bound system due to the expansion is negligible. Finally in section 4 we conclude, summarize our main results and propose possible extensions of this work. ",
        "conclusions": "We have studied the evolution of bound systems in expanding backgrounds and focused on the case of accelerated expansion powered by phantom energy ($w<-1$). We have found the radial time dependence of bound systems in phantom cosmologies and determined the time when these systems dissociate due to the repulsive effects of phantom energy, as a function of the equation of state parameter $w$. A universal behavior was found for the dissociation time for different geometries of bound systems. We have also plotted the bound system trajectory around the time of dissociation and demonstrated that the bound systems explode radially outward after dissociation. Our results were compared with previous corresponding results in the literature and were found to be in qualitative but not in quantitative agreement.  In the present study we have assumed a constant equation of state parameter $w$. The extension of our results to the case of a redshift dependent $w$ ($w(z)$) is straightforward and consists a potentially interesting extension of this work. A potential improvement to the accuracy of our results may come by using a more accurate metric for the interpolation between the Schwarzschild and the Friedmann metric \\cite{Baker:2001yc}. Such an improvement would be more important for strongly bound high velocity systems.  The Mathematica\\cite{wolfram} file used for the production of the figures of the paper can be downloaded from \\cite{mathfile} or sent by e-mail upon request.  {\\bf Acknowledgements:} This work was supported by the European Research and Training Network HPRN-CT-2000-00152."
    },
    "0410/astro-ph0410079_arXiv.txt": {
        "abstract": "Short-term variability of X-ray continuum spectra has been reported for several Active Galactic Nuclei. Significant X-ray flux variations are observed within time scales down to $10^3-10^5$ seconds. We discuss short variability time scales in the frame of the X-ray flare model, which assumes the release of a large hard X-ray flux above a small portion of the accretion disk. The resulting observed X-ray spectrum is composed of the primary radiation and of a reprocessed Compton reflection component that we model with numerical radiative transfer simulations. The incident hard X-rays of the flare will heat up the atmosphere of the accretion disk and hence induce thermal expansion. Eventually, the flare source will be surrounded by an optically -thick medium, which should modify the observed spectra. ",
        "introduction": "In the last few years several Active Galactic Nuclei (AGN) showing short-term X-ray variability have been reported (see for instance Ark 564 and Ton S180 \\citep{Edelson2002}, I Zw 1 \\citep{Gallo2004}, MGC-6-30-15 \\citep{Vaughan2004,Ponti2004}). Modeling this behavior remains a complex task: owing to the short variability time scales observed, one has to assume small emission regions down to the size of a few $R_{Schw}$ (Schwarzschild radii). Such regions exist within the framework of the flare model. In this model compact hard X-ray sources resulting from magnetic reconnection evolve above the accretion disk.  In the following we present detailed numerical modeling of the reprocessed X-ray spectra induced by a magnetic flare. We investigate the influence of a realistic flare illumination as compared to an isotropic one and the dependency on the disk inclination. Using this model we attempt to derive the consequences for the time evolution of the observed spectra.  Then we briefly discuss the hydrostatic balance and the optical properties of a the disk atmosphere below a flare. ",
        "conclusions": ""
    },
    "0410/astro-ph0410286_arXiv.txt": {
        "abstract": "We give preliminary results from a spectroscopic study of composite spectrum hot subdwarfs (sd+late-type).  We obtained spectra of a sample of hot subdwarfs selected from the {\\em Catalogue of Spectroscopically Identified Hot Subdwarfs} {\\citep*{KHD}} on the basis of near-infrared photometry from the {\\em Two Micron All Sky Survey} (2MASS).  The sample consists of 20 photometric and spectroscopic single and 54 composite hot subdwarfs, 6 resolved (or barely resolved) visual doubles, and 5 objects with emission lines or broad absorption lines with emission cores.  Spectra of 84 ``standard'' (single late-type) stars with {\\em Hipparcos} parallaxes were also obtained for calibration.  These observations cover 4600--8900~{\\AA} with $\\sim$3~{\\AA} resolution.  We measured equivalent width-like indices around Mg~I~b, Na~I~D, the Ca~II infrared triplet, H$\\alpha$, and H$\\beta$.  Using the single late-type star observations combined with model energy distributions, we explore how the measured indices of a composite spectrum vary as the temperature and luminosity of the late-type companion are varied and as the temperature and radius of the hot subdwarf are varied.  We use the measured indices of the composite systems to estimate the temperature and gravity of the late-type star, taking into account the dilution of its spectral features by light from the hot subdwarf. ",
        "introduction": "Hot subdwarfs were selected from the {\\emph{Catalogue of Spectroscopically Identified Hot Subdwarfs}} {\\citep{KHD}}.  The observed sample consists of 20 single and 54 composite hot subdwarfs (primarily classified as sdB) that have $V \\la 13.5$.  Composite subdwarfs for this sample were defined to have $J\\!-\\!K_S \\ga$~$+0.05$ {\\citep{SW}}.  There are also 6 subdwarfs that have close ($<$10$''$) resolved companions, and 5 that show emission lines (4 met our composite color criterion).  Eighty-four stars (52 main sequence, 32 subgiant) with parallax measurements from {\\emph{Hipparcos}} were selected to provide empirical calibration.  They cover the Pop~I main sequence from early F to early M-type, and post-main sequence subgiants from F to K-type ($2.7\\la V\\la 10.0$, $-1.1$ $\\la \\mathrm{M}_V \\la 9.2$, and $0.35\\la \\bv\\la 1.6$).  Spectra were taken with the GoldCam spectrograph on the Kitt Peak National Observatory 2.1m telescope.  The wavelength interval 4600--8900~{\\AA} was covered at $\\sim$1.3~{\\AA/pixel}, using two spectrograph setups.  This wavelength range includes H$\\alpha$, H$\\beta$, H-Paschen 11$\\rightarrow$limit, Mg~I~b, Na~I~D, Ca~II infrared triplet (CaT), {He~I} 5875 and 6678~{\\AA}.\\@ Unfortunate observing conditions prevented accurate flux calibration on most nights, so we consider the normalized spectra only.  Example spectra are shown in Fig.~\\ref{Spec}.  \\begin{figure}[!ht] \\resizebox{\\textwidth}{!}{\\rotatebox{-90}{\\plotone{fig1.eps}}} \\caption{Example spectra for a single hot subdwarf (top), a composite hot subdwarf (middle), and a K1V standard (bottom) offset by constants.  It is evident that there is less dilution at longer wavelengths. \\label{Spec}} \\end{figure} ",
        "conclusions": "} We obtained spectra of composite colored hot subdwarfs and single cool standard stars with {\\emph{Hipparcos}} parallaxes.  Preliminary analysis indicates {\\em the majority of cool companions in composite hot subdwarf systems are consistent with the main sequence}.  Overall, subgiant companions do not explain the observed colors and EWs for the majority of composite hot subdwarfs; but they are not excluded in some cases if the hot subdwarf is relatively cool (i.e., $\\sim$20kK). Further analysis of additional features in individual spectra and refinement of the model grids is needed to distinguish between main sequence and subgiant in these cases.  An independent measure of the temperatures of the hot subdwarfs (from IUE spectra or GALEX fluxes) would further constrain uncertain cases."
    },
    "0410/astro-ph0410443_arXiv.txt": {
        "abstract": "There are many different phenomenological models describing the cosmological dark energy and accelerating Universe by choosing adjustable functions. In this paper we consider a specific model of scalar tachyon field which is derived from the NSR string field theory and study its cosmological applications. We find that in the effective field theory approximation the equation of state parameter $w<-1$, i.e. one has a phantom Universe. It is shown that due to nonlocal effects there is no quantum instability that the usual phantom models suffer from. Moreover due to a flip effect of the potential the Universe does not enter to a future singularity. ",
        "introduction": "It was suggested by Ia Supernova observations that the Universe is presently  accelerating\\cite{Perlm,Riess}. The basic sets of experiments now includes also the Cosmic Microwave Background (CMB) anisotropies,  X-ray data from galaxy clusters, large scale structure and age estimates of global clusters, for a review  see \\cite{Spergel,S-St,Sachni}. It is believed that a new particle and/or gravitational physics is required to explain the acceleration of the expansion of the Universe. The observations suggest that the bulk of energy density in the Universe is gravitationally repulsive and appears like an unknown form of energy (dark energy) with negative pressure. It is believed that 2/3 of the total density of the universe is in a form of dark energy.  There exist many different models of dark energy. It is convenient to describe them by using the equation of state parameter $w=p/\\rho$, where $p$ is a pressure and $\\rho$ is the energy density. The analysis of the current observation data  shows that $w$  lies in the range $-1.61<w<-0.78$ \\cite{knop,Spergel,Tegmark}.  A list of dark energy models includes (see \\cite{S-St}-\\cite{kitaj} and refs. therein) \\begin{itemize} \\item $w=-1$: the cosmological constant; \\item $w=const \\neq -1$: the cosmic strings, domain walls, etc.; \\item $w \\neq const$: quintessence scalar field, chaplygin gas, k-essence, Dirac-Born-Infeld(DBI) action, braneworlds, etc.; \\item $w < -1$: phantom models. \\end{itemize}  The most challenging for theoretical physics would be the case $w < -1$ \\cite{Caldwell}-\\cite{Melchiorri} (see \\cite{Frampton,Sachni,Padmanabhan-rev} for  reviews). In this case the weak energy condition ($\\rho>0, \\rho+p>0$) is violated and some strange phenomena as negative entropy and temperature appear. In many models the phantom Universe in finite time ends up in the singularity called the Big Rip \\cite{Caldwell03}).   In this paper we consider string field theories nonlocal tachyon where the condition $w < -1$ is realized. We show that due to a peculiar properties of nonlocal tachyon dynamics we can get a phantom universe without problems with unstability. A nonlocal dynamics for the tachyon field is obtained by truncation of the covariant string field equations. We study a nonlocal dynamics of a string tachyon in the cosmological Friedmann metric. The string theory that we have in mind is the NSR string theory compactified on a six dimensional compact manifold. Moreover we assume that all moduli are frozen and unstable non-BPS brane extends along the three large spatial dimensions. The tension of the 3-dimensional non-BPS brane acts as the 4-dimensional universe cosmological constant and the tachyon dynamics describes a deviation from a pure cosmological constant regime.  A  presence of  tachyons was considered as a main drawback of corresponding  strings theory. Bosonic string has a tachyon and its absence in superstrings was a main motivation to introduce superstrings. However few years ego tachyon has  found an application in the context of brane scenarios. The open bosonic string and GSO $-$ NSR string tachyons were used to describe D-brane decays. According to the Sen conjecture in the perturbative vacuum there are instable branes filling space-time and all these branes disappear in the true vacuum (for review see \\cite{review-sft},\\cite{Sen-g}). Rolling tachyon solutions describe transitions from perturbative vacuum to non-perturbative one.  Here we use the covariant open string field theory(SFT) approach \\cite{review-sft}  to tachyon dynamics (about the effective  DBI approach to tachyon dynamics see \\cite{Sen},\\cite{0301117},\\cite{Sen-g} and refs. therein). A crucial feature of  the rolling tachyon solution for non-BPS brane \\cite{AJK} is that it interpolates between two non-perturbative vacua with the same energy. We argue that at large time  the dynamics is governed   by an effective action that contains a ghost kinetic term, in spite of absence of ghosts in the nonlocal action. We emphasize that the ghost here is a result of an approximation to the exact nonlocal action. There is no ghost in the nonlocal action. So we get a phantom universe as an approximation to the true nonlocal dynamics. This is the reason why there are no pathologies in this scenario with the thermodynamical instabilities.  The rolling tachyon solution passes  the perturbative vacuum with non-zero velocity. Therefore, to reach the true vacuum starting from the perturbative one we have to supply  the tachyon with a large enough initial velocity . This is a rather non-standard situation from the local field theory point of view, where one does not have to push strongly the  tachyon to make it reach the true vacuum, it is need just an infinitesimal small perturbation. This effect occurs due to the presence of derivative terms in the interaction. These terms may be studied in an effective action approximation, which corresponds to keeping only few terms of an expansion of an nonlocal operator. We note  that to reach the nonperturbative vacuum one has to add to the action a brane tension which is larger that is required by the Sen hypothesis. This large  brane tension can be interpreted as an effect of the closed string excitations.  We find that in the effective field theory approximation the equation of state parameter $w<-1$, i.e. one has a phantom Universe. In the DBI approach  the equation of state parameter interpolates between -1 at early time and 0 at a later time. Within application of the DBI action to cosmology there are  problems with large density perturbations, reheating and caustics formation \\cite{Kofman,FKS}.  The paper is organizes as follows. In Sect.  2.1. the nonlocal tachyon action is written in the Friedmann background. In Sect. 2.2 we shortly remind a construction of the rolling tachyon solution for non-BPS brane \\cite{AJK}. In Sect. 2.3 we study the tachyon dynamics in the Friedmann metric in the effective action approximation. In this approximation one gets $w<-1$. In Sect. 2.4 we show that to reach a nonperturbative vacuum one has to add to the action a brane tension which is larger that is required by the Sen hypothesis. In Sect. 3 we show the stability of the nonlocal tachyon model in the true vacuum. ",
        "conclusions": "We have studied  the evolution of open GSO $-$ NSR string tachyon in the Friedmann space-time. The corresponding  solution in the flat space-time is known as a rolling tachyon and it describes the decay of the space filling D3 brane corresponding to the unstable perturbative vacuum to the local stable vacuum. We have performed calculations under the following approximations and assumptions: \\begin{itemize} \\item the level truncation and an approximation of a slow varying axillary field; \\item  a direct generalization of the  tachyon nonlocal action to the Friedmann space-time; \\item an effective local action approximation. \\end{itemize}  We have found that in the effective field theory approximation the equation of state parameter $w<-1$, i.e. one has a phantom Universe, but there is no problem with  quantum instability without this approximation. We have found  that to reach the nonperturbative vacuum one has to add to the action a brane tension which is larger that is required by the Sen hypothesis. This large  brane tension can be interpreted as an effect of the closed string excitations.   It would be interesting to try to find an analog of this solution to a full SFT theory (without level truncation), in particular with a framework of Vacuum String Field Theory. VSFT is a version of usual SFT that is supposed to describe the theory at the minimum of the tachyon potential. Corresponding solutions for bosonic VSFT have been found recently in  \\cite{LB}. It would be also interesting to see if one has to shift the tension of D-brane to get a rolling solution in the Friedmann metric.  $$~$$  {\\bf Acknowledgements}  I would like to thank L. Bonora, B. Dragovich, A. Koshelev, S. Mukhanov, S. Vernov and I. Volovich for useful discussions. This work is supported in part by RFBR grant 02-01-00695 and INTAS grant 03-51-6346. $$~$$"
    },
    "0410/astro-ph0410675_arXiv.txt": {
        "abstract": "We present the first full map of the H\\ensuremath{\\alpha}\\ emission from the inter\\-mediate-velocity cloud Complex L. Kinematically resolved emission components from the ionized gas reveal a close spatial and velocity correspondence to that of the neutral gas as mapped by 21 cm line observations. Using a simple Gaussian component fitting technique, we compare the details of the emission from the two phases of Complex L. The mean velocity of the H\\ensuremath{\\alpha}\\ and 21 cm emission are strongly correlated, but we find that the some of the ionized gas is systematically offset ($\\sim -4$ km~s\\ensuremath{^{-1}}) relative to the neutral gas. Line widths are qualitatively comparable, but both emission components are too low signal-to-noise to study any differences between the phases. Finally, as with our previous study toward Complex K, we find \\ensuremath{I_{\\mathrm{H}\\alpha}}\\ is not strongly correlated to \\ensuremath{N_{\\mathrm{H\\;I}}}\\ along the same lines of sight through Complex L. ",
        "introduction": "Intermediate- and high-velocity clouds (IVCs and HVCs) are routinely being studied in optical emission lines thanks to advances in detector and instrument design (\\citealp{Putman+03, Tufte+02}; \\citealp*{WVW02,HRT01,TRH98}). In particular, the efficiency of the Wisconsin H-Alpha Mapper (WHAM) allows a multi-wavelength study of the ionized component over the entire extent of these predominantly neutral structures. The relationship between the H\\ensuremath{\\alpha}\\ and 21 cm emitting gas should provide important clues into the source of ionization and conditions within the Galactic halo. Lyman continuum radiation is believed to escape the disk with sufficient flux to maintain the warm ionized medium (WIM) of the Milky Way and thus may also ionize the downward facing surfaces of neutral halo material (\\citealp{WVW02,B-HM99}; \\citealp*{DSF00,MC93}). However, high ionization species are also detected toward these complexes, suggesting that alternative forms of ionization and sources of Lyman continuum radiation may be important. Conductive interfaces or turbulent mixing layers arising from the interaction of the neutral gas with a very low density halo component have often been cited as mechanisms able to complete the picture of these anomalous structures \\citep{Fox+04}.  We have embarked on an extended multi-wavelength study of the large-scale complexes with WHAM to explore the physical conditions in the ionized gas of IVCs and HVCs and to study the relationship of this component to the neutral gas. Here we present the first map of the ionized component of Complex L, first designated by \\citet{WvW91} to be 21 cm emission present between $-190$ km~s\\ensuremath{^{-1}}\\ $< \\ensuremath{v_{\\mathrm{LSR}}} < -85$ km~s\\ensuremath{^{-1}}\\ in the approximate range $341\\deg < \\ell < 348\\deg$ and $31\\deg < b < 41\\deg$. In this work, we examine a slightly larger region of the sky that shows intermediate-velocity emission in both 21 cm and H\\ensuremath{\\alpha}\\ at slightly higher latitudes and longitudes that may be physically related to the gas previously defined to be Complex L.  There have been a few previous detections of H\\ensuremath{\\alpha}\\ emission from Complex L. \\citet*{WVW01} and \\citet{Putman+03} find emission ($\\ensuremath{I_{\\mathrm{H}\\alpha}} \\sim$ 0.2 -- 1 R) toward several directions near higher column-density clumps of the complex. These emitting regions are also reported to have very high [\\ion{N}{II}]/H\\ensuremath{\\alpha}\\ ratios (1 -- 2.7). Here, we present the first map of the entire ionized component of Complex L. ",
        "conclusions": "We have presented the first map of the ionized gas in the intermediate-velocity cloud Complex L. The ionized gas is very clearly spatially and kinematically linked to the neutral gas. However, the intensity of the H\\ensuremath{\\alpha}\\ emission arising from the ionized component of the complex varies independently from the 21 cm emission of the neutral component. Interesting trends including a small systematic offset in the mean velocity of each component and a hints of a relationship in a large sample of \\ensuremath{I_{\\mathrm{H}\\alpha}}\\ and \\ensuremath{N_{\\mathrm{H\\;I}}}\\ measurements may start to provide clues about the details of what sustains the ionization in IVCs and HVCs.  In addition to mapping out other complexes in H\\ensuremath{\\alpha}, we will soon begin exploring the physical conditions in the ionized component of these halo clouds by leveraging the multi-wavelength capabilities of WHAM."
    },
    "0410/astro-ph0410396_arXiv.txt": {
        "abstract": "From newly-obtained high-resolution, high signal-to-noise ratio spectra the abundances of the elements La and Eu have been determined over the stellar metallicity range $-3<$[Fe/H]$<+0.3$ in 159 giant and dwarf stars. Lanthanum is predominantly made by the $s$-process in the solar system, while Eu owes most of its solar system abundance to the $r$-process. The changing ratio of these elements in stars over a wide metallicity range traces the changing contributions of these two processes to the Galactic abundance mix.  Large $s$-process abundances can be the result of mass transfer from very evolved stars, so to identify these cases, we also report carbon abundances in our metal-poor stars.  Results indicate that the $s$-process may be active as early as [Fe/H]$=-2.6$, alalthough we also find that some stars as metal-rich as [Fe/H]$=-1$ show no strong indication of $s$-process enrichment.  There is a significant spread in the level of $s$-process enrichment even at solar metallicity. ",
        "introduction": "Abundances of elements determined in stars over a very wide metallicity range contain vital clues to the chemical evolution of our Galaxy.  This is especially true for the neutron capture ($n$-capture) elements, those heavier than the Fe peak (Z$>$30). Syntheses of $n$-capture elements occur in a variety of fusion episodes in the late stages of stellar evolution, with neutron densities that range over factors of 10$^{15}$.  During quiescent He-burning, neutrons generated via the  \\iso{13}{C}($\\alpha$,n)\\iso{16}{O} and \\iso{22}{Ne}($\\alpha$,n)\\iso{25}{Mg} reactions are captured by heavy element seed nuclei.  The neutron densities are relatively low (N$_n\\sim$10$^{8}$~cm$^{-3}$; \\citealt{Pagel1997}), so nearly all possible $\\beta$-decays will have time to occur between successive neutron captures.  Synthesis of successively heavier isotopes progresses along the ``valley of $\\beta$-stability''.  This synthesis route is called the $s$-process, and it is responsible for about half of the isotopes of the $n$-capture elements.  Much larger neutron densities ($\\geq$10$^{23}$~cm$^{-3}$; \\citealt{Pagel1997}) are generated via p~+~e~$\\rightarrow$ n~+~$\\nu$ during the high-mass-star core collapse that results in Type~II supernovae (SNe).  Extremely neutron rich nuclei are formed in a matter of seconds; in this so-called $r$-process, it is not necessary to have pre-existing heavy-element seed nuclei.  The $n$-capture rates exceed $\\beta$-decay rates, creating nuclei out to the neutron drip line.  These extremely neutron rich nuclei experience multiple $\\beta$-decays back to the valley of $\\beta$-stability after the very quick extinction of $n$-capture events as the Type~II SN envelopes are ejected.  The $r$-process is also responsible for about half of the solar system $n$-capture isotopes but not always the same ones created in the $s$-process. Thus, $n$-capture elements can be composed of some pure $r$-process, pure $s$-process, and some mixed-parentage isotopes.  Abundance studies of very low metallicity stars cannot measure isotopic abundances from typical absorption lines; the isotopic wavelength offsets are usually very small compared with other line-broadening effects. The only exceptions for the Z$>$30 elements to date are Eu \\citep{Sneden2002, Aoki2003} and Ba \\citep{Magain1995, Lambert2002}.  Therefore, to help understand the relative importance of the two $n$-capture synthesis mechanisms throughout the history of the Galaxy, stellar elemental abundance comparisons are made between elements whose solar system isotopic abundances are dominated by the $s$-process and those due mainly to the $r$-process.  Beginning with \\citet{Spite1978}, such comparisons have nearly always focused on the relative abundances of Ba (Z=56, $\\approx$85\\% $s$-process in the solar system) and Eu (Z=63, $\\approx$97\\% $r$-process; see \\citealt{Burris2000} and references therein).  The massive stars that are althought to provide the site of the $r$-process evolve on much shorter time scales than the lower mass stars that manufacture $s$-process elements.  On the average, then, older stars ought to show a pure or almost pure $r$-process signature \\citep{Truran1981}.  While simple in principle, observational evidence for the onset of the $s$-process in a general way is still lacking.  A recent and extensive study of $n$-capture elements in metal-poor giant stars by \\citet{Burris2000} offers a somewhat inconclusive answer to the question of when low-mass stars begin to contribute to the Galactic chemical mix.    They found that  the $s$-process may begin to contribute at lower metallicities than expected, an effect somewhat blunted by the considerable range in Ba/Eu abundance ratios at any given metallicity.   In contrast, \\citet{Mashonkina2003} found that the onset of the $s$-process occurred at higher metallicities and that evolution of the $s$-/$r$-process ratio could be tied to Galactic stellar populations.  There is much less variation in the Ba/Eu ratio in the latter study.  Unfortunately, the observational and analytical uncertainties in [Ba/Eu] abundance ratios are substantial and may contribute considerably to the ambiguity of previous results.  Often, Ba abundances are derived from just a handful of \\ion{Ba}{2} transitions (mainly the 4554~\\AA\\ resonance line, now sometimes accompanied by the 5853, 6141, and 6496~\\AA\\ lines), and Eu abundances from the \\ion{Eu}{2} transitions at 4129 and 4205~\\AA ~\\citep{Francois1996, McWilliam1998, Burris2000}. Detailed hyperfine and isotopic substructure analyses have been published for \\ion{Eu}{2} features along with reliable transition probabilities, so accurate Eu abundances may be routinely determined \\citep{LawlerEu}.  Barium abundance uncertainties now dominate the problem.  The near-ground-state \\ion{Ba}{2} lines most commonly measured are usually very strong, while higher excitation lines (at 3891, 4130, and 4166~\\AA) are almost all undetectable. The low-excitation transitions are often saturated, yielding Ba abundances that are very sensitive to adopted values of stellar microturbulent velocities. In In addition, Ba has five abundant isotopes (\\iso{134}{Ba}-\\iso{138}{Ba}) which are synthesized in different proportions by $r$- and $s$-processes.  Therefore, elemental abundances depend on the assumed $r$/$s$ fraction, which is what one is ultimately attempting to determine; these difficulties are especially acute for the 4554~\\AA ~line. From these effects, it is difficult to assess the Ba abundances in most stars to better than $\\pm$0.2~dex. It is not at all obvious whether observed star-to-star scatter in [Ba/Eu]\\footnote{We adopt the standard spectroscopic notation where   [A/B]=[log$_{10}\\epsilon(A)-$log$_{10}\\epsilon(B)]_*-[$log$_{10}\\epsilon(A)-$log$_{10}\\epsilon(B)]_{\\odot}$, where log$_{10}\\epsilon(X)=$log$_{10}(N_X/N_H)+12$ . However, for the purposes of this paper, it is usually easier to work in the log$~\\epsilon$ notation, where  log$_{10}\\epsilon(A/B)=$log$_{10}(N_A/N_H)-$log$_{10}(N_B/N_H)$. } ratios at a given [Fe/H] metallicity is primarily indicative of astrophysical variations or Ba measurement uncertainties.  Fortunately, neighboring rare earth elements whose solar system abundances are mainly due to the $s$-process may be easily observed in stars over a wide metallicity range. In particular, La (Z=57, $\\approx$75\\% $s$-process in the solar system; \\citealt{Burris2000}) has many detectable features of \\ion{La}{2}. This species has enjoyed a recent laboratory study that yielded accurate transition probabilities and hyperfine structure constants \\citep{LawlerLa}. There are \\ion{La}{2} lines of different strength throughout the near-UV to red spectral regions, and all of them yield consistent abundances from the solar spectrum and from $r$-process-rich, low-metallicity stars such as \\bd17, \\cs22, and BPS BPS CS 31082-001 \\citep{LawlerLa, CS22892, Hill2002}.  Finally, \\iso{139}{La} is the only abundant naturally-occurring isotope of La, so $r$- and $s$-process contribution estimates are irrelevant to the determination of total elemental La abundances \\citep{Lodders2003}.  In this paper we report new La and Eu abundances for 159 giant and dwarf stars.  Our aim is to refine measurements of the $s$-/$r$-process ratio to the point where measurement uncertainties no longer obscure the onset of nucleosynthetic contributions from asymptotic giant branch (AGB)  stars.  Our data also allow us to address $s$-process evolution in the Galactic stellar populations that we have sampled.  In \\S \\ref{obs} we present the spectroscopic data, in \\S \\ref{ana} we discuss the abundance analyses, in \\S \\ref{compare} we discuss our results in relation to previous work, and in \\S \\ref{disc} we discuss the Galactic trend in La/Eu ratios and the relationship between La/Eu and stellar kinematics.  Section \\ref{con} contains some concluding remarks. ",
        "conclusions": "\\label{con}}   In this study we observed 159 giant and dwarf stars across a wide range of metallicity in order to measure the evolution of the abundance ratio La/Eu, a proxy for the $s$-/$r$-process ratio.  We   have found that the $s$-/$r$-process ratio does not increase monotonically with [Fe/H].  However, there is evidence for evolution in this ratio.  At low metallicity, the abundances of La and Eu are approximately equal; near solar metallicity, La is consistently more abundant.   The $s$-process contribution to individual stars varies widely even at near-solar metallicities.  However, we find that on the average the dispersion in La and the dispersion in Eu individually are equal, indicating that the $r$-process contribution, while smaller overall, also varies.  Some of the variation in La/Eu can be attributed to the intrusion of different Galactic stellar populations.  We  find that when stars are separated by velocity, very little scatter in La/Eu remains.  Rather, stars separate into  high-velocity and low-velocity groups.  The former group has essentially a single value of La/Eu (although there is some spread, and a few stars stand out in this respect; the overall spread in abundances is about 0.2 dex here).  The latter group shows considerable dispersion still (about 0.4 dex overall), although at a higher overall La/Eu value. This variation is further reduced when stellar age is considered.  The age correlation and the velocity data argue  most strongly for evolution in the $s$-/$r$-ratio in the  disk, with very little change throughout the halo.  This sample also includes thick disk stars identified by other studies, although these are not readily distinguishable from the halo based solely on La/Eu.  Neither stellar population considerations nor measurement  uncertainties can account for the persistently high La/Eu ratio at low metallicities, where the $r$-process is believed to be dominant.  None  of the stars studied here has a ``pure $r$-process\" value of La/Eu,  in contrast to the $r$-process rich star \\cs22, studied with the same La and Eu atomic data.  It is still unclear whether this indicates  that the $s$-process was indeed active in the halo at very low metallicities and established a ``floor\" in La/Eu for old stars or simply that the ``pure $r$-process\" value of La/Eu, despite extensive recent work, is in error."
    },
    "0410/astro-ph0410505_arXiv.txt": {
        "abstract": "We have studied the aperiodic X-ray timing and color behavior of the accreting millisecond pulsars SAX J1808.4--3658, XTE J1751--305, XTE J0929--314, and XTE J1814-338 using large data sets obtained with the {\\it Rossi X-ray Timing Explorer}. We find that the accreting millisecond pulsars have very similar timing properties to the atoll sources and the low luminosity bursters. Based on the correlation of timing and color behavior SAX J1808.4--3658 can be classified as an atoll source, and XTE J0929--314 and XTE J1814--338 are consistent with being atoll sources, but the color behavior of XTE J1751--305 is different. Unlike in other atoll sources the hard color in the extreme island state of XTE J1751--305 is strongly correlated with soft color and intensity, and is not anti-correlated with any of the characteristic frequencies. We found previously, that the frequencies of the variability components of atoll sources follow a universal scheme of correlations. The frequency correlations of the accreting millisecond pulsars are similar, but in the case of SAX J1808.4--3658 and XTE J0929--314 shifted compared to those of the other atoll sources in a way that is most easily described as a shift in upper and lower kilohertz QPO frequency by a factor close to 1.5. Although, we note that the shift in lower kilohertz QPO frequency is based on only one observation for SAX J1808.4--3658. XTE J1751--305 and XTE J1814--338, as well as the low luminosity bursters show no or maybe small shifts. ",
        "introduction": "Many low mass X-ray binaries containing a neutron star show kilohertz QPOs, quasi periodic oscillations with frequencies ranging from a few hundred to more than a 1200 Hz. The low frequency ($<200$ Hz) part of the power spectrum taken from their X-ray lightcurves is usually dominated by a broad band-limited noise component. In addition to the band-limited noise several Lorentzian components below 200 Hz are present (see e.g. \\cite{vstr02:apj568}). These Lorentzians become broader as their characteristic frequency decreases (\\citealp*{psaltis99:apj520}; \\citealp{vstr02:apj568}). Therefore, several features can, only at high frequencies, be classified as a QPO \\citep[traditionally defined by the condition FWHM $<$ centroid frequency/2, see][]{vdk95:xrb252}. In addition to these timing features in the persistent emission, 13 sources show millisecond oscillations (270--619 Hz) during thermonuclear X-ray bursts \\citep[see][for a review]{muno04:rossi}. Although these burst oscillations show a slight drift, they are thought to be connected with the spin of the neutron star and were the first direct indication that these sources harbour neutron stars with millisecond spins.  Based on the correlated behavior of their timing properties and the behavior of the sources in color-color diagrams, that are used to study their X-ray spectral properties, many of the neutron-star low-mass X-ray binaries are traditionally classified as either ``atoll'' or ``Z'' sources \\citep[][]{hasinger89:aa225}. Recently, \\citet{muno02:apj568} and independently \\citet{gierlinski02:mnras331} used the large {\\it Rossi X-ray Timing Explorer} data sets now available for several of the atoll sources, to show that some of the atoll sources look much more like Z sources when plotted in color-color diagrams than was previously thought. Although the appearances in the color-color diagram of the two classes now is more similar, the classification still holds, as clear differences between the two classes in both timing and color behavior remain \\citep{muno02:apj568,barret02:apj576,olive03:apj583,vstr03:apj596,reig04:apj}. Note, that next to the differences in timing and color behavior there is also a difference in luminosity between the two classes; where the Z sources are all very luminous \\citep[$\\sim 3 \\times 10^{38} erg s^{-1}$][]{ford00:apj537}, the atoll sources occur over a wide range of luminosities \\citep[$10^{36-38} erg s^{-1}$][]{ford00:apj537}.  The sources 1E 1724--3045, GS 1826--24, and SLX 1735--269, are all low-luminosity neutron stars that show thermonuclear X-ray bursts. The power spectra of these ``low-luminosity bursters'' resemble those of the atoll sources in their hard states (\\citealp{olive98:aa333}; \\citealp{barret00:apj533}; \\citealp*{belloni02:apj572}); they show a broad band-limited noise component and several broad Lorentzian components. However, where the characteristic frequencies of the low-frequency Lorentzians were found to be similar to those in the atoll sources, the characteristic frequency of the highest frequency Lorentzian seemed to be higher \\citep[][]{belloni02:apj572}.  In 1998 the first accreting millisecond pulsar was found with the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) during an outburst of SAX J1808.4--3658 \\citep[][]{wijnands98:nat394}. During this outburst no thermonuclear X-ray bursts were observed and the source did not show kilohertz QPOs above 200 Hz \\citep[][]{wijnands98:apj507}. The low frequency power spectrum was very similar to that of the atoll sources \\citep[][]{wijnands98:apj507} and both the atoll sources and SAX J1808.4--3658 followed one relation when the characteristic frequency of a low-frequency Lorentzian was plotted versus that of the band-limited noise \\citep[][]{wijnands99:apj514}. During a new outburst of SAX J1808.4--3658 in October 2002, kilohertz QPOs (up to 700 Hz) and burst oscillations were observed in an accreting millisecond pulsar for the first time \\citep[][]{wijnands03:nat424,chakrabarty03:nat424}. Recently, four additional accreting millisecond pulsars were discovered; XTE J1751--305, XTE J0929--314, XTE J1807--294, and XTE J1814--338 \\citep[see][for a observational review]{wijnands04:bepposax}. Of these XTE J1807--294 showed kiloherz QPOs (Markwardt, private communication) and XTE J1814--338 showed burst oscillations \\citep{strohmayer03:apj596}.  Here we present an extensive analysis of the aperiodic-timing and color behavior of the accreting millisecond pulsars SAX J1808.4--3658, XTE J1751--305, XTE J0929--314, and XTE J1814--338. We study the correlations between the characteristic frequencies of their various timing features, and compare these with those of four well-studied atoll sources and three low-luminosity bursters. We find, that the accreting millisecond pulsars have very similar timing properties as the atoll sources and the low luminosity bursters, although discrepancies occur for SAX J1808.4--3658 and XTE J0929--314. We also compare the behavior of the accreting millisecond pulsars in the color-color diagrams to that of the atoll sources and find that one pulsar, XTE J1751--305, exhibits unusual behavior while the others act like ordinary atoll sources. In the next section we describe the observations and data analysis. We suggest that readers not interested in the technical details of observations and analysis skip immediately to the results section. ",
        "conclusions": "\\subsection{Correlations of Timing Features} \\label{sec.compare_timing}  The frequencies of the variability components of the atoll sources, 4U 0614+09, 4U 1608-52, and 4U 1728-34, follow a universal scheme of correlations when plotted versus $\\nu_{u}$. In Figure \\ref{fig.freq_freq} we present such a plot for SAX J1808.4--3658, XTE J1751--305, XTE J0929--314, and XTE J1814--338 (colored points). We also include (black points) results for the atoll sources 4U 0614+09, 4U 1728-34 \\citep{vstr02:apj568}, 4U 1608-52 \\citep{vstr03:apj596}, and Aql X--1 \\citep{reig04:apj}. For Aql X--1 we only include the results for the extreme island state for which $\\nu_{u}$ could unambiguously be identified. For 4U 1728-34 we also include a recent result by \\citet*{migliari03:mnras345}, where for the first time twin kilohertz QPOs were observed in the island state. As a third group of sources we have included the low luminosity bursters 1E 1724--3045, GS 1826--24, and SLX 1735--269 (grey points) described in \\citet{belloni02:apj572}. \\citet{belloni02:apj572} constructed power spectra only up to 512 Hz, which led to an ill-constrained L$_u$. Here we re-analyze the timing properties of groups C--R of \\citet{belloni02:apj572} using the same method as described in \\S \\ref{sec.anal_timing} (Tables \\ref{tbl.numax}, \\ref{tbl.qvalues}, and \\ref{tbl.rms}). With 0.0039--4096 Hz power spectra this leads to a much better constrained L$_u$.  SAX J1808.4--3658 shows similar relations between its characteristic frequencies as the atoll sources do. However, with respect to these sources, the relations of SAX J1808.4--3658 are shifted. XTE J0929--314 shows a similar shift, although this is only based on one power spectrum. XTE J1751--305 and XTE J1814--338 as well as the low luminosity bursters show relations consistent with those of the atoll sources. The shift occurs only between the characteristic frequencies of the low frequency components on the one hand and $\\nu_{u}$ (and probably $\\nu_\\ell$) on the other (see below): the mutual relations between the characteristic frequencies of the {\\it low frequency components} are identical for all sources (see Figure \\ref{fig.nu_vs_nub}).  To determine the actual shift factors between the frequency relations of SAX J1808.4--3658, XTE J0929--314, and those of the atoll sources we consider the $\\nu_i$ vs. $\\nu_u$ relations, where $\\nu_i$ is either $\\nu_b$, $\\nu_h$, or $\\nu_{\\ell ow}$. We only use that part of the relations for which $\\nu_u < 600$ Hz, as the behavior of the low-frequency components above 600 Hz is complex (see also \\S \\ref{sec:physical}). We take a $\\nu_i$ vs. $\\nu_u$ relation of one source (for example SAX J1808.4--3658) together with that of the atoll sources and perform a power law fit. As our points have errors in both coordinates, we use the FITEXY routine of Press \\citet{press92:comp274} that performs a straight-line fit to data with errors in both coordinates. We take the logarithm of both $\\nu_i$ and $\\nu_u$ so that fitting a straight-line becomes equivalent to fitting a power law. Before fitting we multiply $\\nu_u$ (or $\\nu_i$) with a shift factor. We let this  shift factor run between 0.1 and 3 with steps of 0.001. The fit with the minimal $\\chi^2$ then gives the shift factor. The errors in the shift factor use $\\Delta\\chi^{2}$ = 1.0 (corresponding to a 68\\% confidence level). In Table \\ref{tbl.shftfactors} we list the shift factors both in $\\nu_u$ and $\\nu_i$ for the different relations and sources. Next to the sources that show a clear shift in Figure \\ref{fig.freq_freq} (SAX J1808.4--3658 and XTE J0929--314) we also include the other millisecond pulsars as well as the low luminosity bursters. We also determine the shift factor for the single $\\nu_\\ell$ observation of SAX J1808.4--3658. We determine the shift factor if only $\\nu_\\ell$ is shifted, and the shift factor if both $\\nu_\\ell$ and $\\nu_u$ are shifted. For the latter we first multiply $\\nu_u$ with 1.454, the weighted average of the $\\nu_u$ shift factors for the $\\nu_b$, $\\nu_h$, and $\\nu_{\\ell ow}$ vs. $\\nu_u$ relations of SAX J1808.4--3658 (see below).  The shift factors in $\\nu_u$ for SAX J1808.4--3658 are between 1.4 and 1.5 when determined with the $\\nu_b$ and $\\nu_h$ vs. $\\nu_u$ relation. For the $\\nu_{\\ell ow}$ vs. $\\nu_u$ relation the factor is higher, but this is based on fewer points which results in a larger error on the shift factor. From this point on we take the shift in $\\nu_u$ to be the weighted average of the three which results in 1.454$\\pm$0.009. After shifting $\\nu_u$ with this factor, the additional shift needed for the single $\\nu_\\ell$ point is 1.359, similar to the shift factors found for $\\nu_u$. For XTE J0929--314 the $\\nu_u$ shift factors determined with the $\\nu_b$ and $\\nu_h$ vs. $\\nu_u$ relation are similar to those in SAX J1808.4--3658, but note that these results are based on only one point from XTE J0929--314. The $\\nu_{\\ell ow}$ vs. $\\nu_u$ relation, just like for SAX J1808.4--3658, gives a higher value but with large error. For some of the other sources small shifts might exist, but unlike for SAX J1808.4--3658 the frequency vs. frequency relations for these sources have up to know only been observed over a small frequency range. The frequency relations should be measured over larger frequency ranges to confirm or reject the idea that small shifts occur.  The simplest explanation for this shift is that some physical difference between these sources affects $\\nu_{u}$ and $\\nu_\\ell$. To make the relations of SAX J1808.4--3658 and XTE J0929--314 coincide with those of the atoll sources, $\\nu_{u}$ and $\\nu_\\ell$ must be multiplied with {$\\sim1.5$}. In Figure \\ref{fig.freq_freq_shifted} we replot Figure \\ref{fig.freq_freq}, but now with the $\\nu_{u}$ and $\\nu_\\ell$ of SAX J1808.4--3658 and XTE J0929--314 multiplied by 1.454. The alternative, that $\\nu_\\ell$ and the characteristic frequencies of all the low frequency components, L$_{b}$, L$_h$, L$_{\\ell ow}$, and the narrow QPOs (see \\S \\ref{sec.identify} for a discussion about the narrow QPOs) are affected is more complex. In this case to make the relations coincide the characteristic frequencies of all the low frequency components have to be multiplied with about 0.35, $\\nu_\\ell$ would have to be multiplied with about 0.8, and $\\nu_{hHz}$ would have to remain the same. It is simpler if $\\nu_{u}$ and $\\nu_\\ell$ are the shifted frequencies. Note, that the resulting $\\chi^2$/dof of the power law fits (see Table \\ref{tbl.shftfactors}) are independent of whether $\\nu_{u}$ or the frequencies of the low frequency components are shifted.  We tried to further test our idea that the shift occurs in $\\nu_{u}$ and not in the frequencies of all the low frequency components, by plotting $Q$ versus $\\nu_{\\rm max}$ for both L$_u$ and L$_h$ (see Figure \\ref{fig.nu_vs_q}). The $Q_u$ vs. $\\nu_{u}$ plot showed a shift in $\\nu_{u}$ similar to that in Figure \\ref{fig.freq_freq}, while the $Q_h$ vs. $\\nu_{h}$ plot did not show a clear shift, consistent with the simpler interpretation above. However, the scatter in the $Q$ versus $\\nu_{\\rm max}$ correlations is much larger than of the frequency-frequency correlations in Figure \\ref{fig.freq_freq}. Neither plots of fractional rms versus $\\nu_{\\rm max}$ nor linking Figure \\ref{fig.freq_freq} to the atoll source states led to extra clues.  Might we have wrongly identified the highest frequency Lorentzian component as L$_u$ while it really is L$_\\ell$? There are a number of arguments against this, all based on similarities with the atoll source class of which SAX J1808.4--3658 seems part (see \\S \\ref{sec.1808_colors}). First, L$_\\ell$ was detected together with the component identified as L$_u$ in SAX J1808.4--3658 by \\citet{wijnands03:nat424} (our group 3, see \\S \\ref{sec.1808_2002}) when the source was almost in the lower left banana (see \\S \\ref{sec.1808_colors}). In group 3 L$_\\ell$ is weak and L$_u$ is strong (see Table \\ref{tbl.rms}); this is common behavior for atoll sources when they are almost in the lower left banana \\citep{mendez01:apj561,vstr02:apj568,vstr03:apj596}. Second, in the island and extreme island states of SAX J1808.4--3658 the component identified as L$_u$ is a strong broad component that is present over a broad frequency range while L$_\\ell$ is not. This is again common behavior for the atoll sources, where L$_\\ell$ was observed in an island state only once, and then as a narrow and weak QPO \\citep{migliari03:mnras345}.  \\subsection{Identifying Timing Features} \\label{sec.identify}  The frequency relations of the accreting millisecond pulsars and the low luminosity bursters of Figure \\ref{fig.freq_freq}, and especially the shift between the relations of SAX J1808.4--3658 and the atoll sources, provide new clues about the identification of some power spectral components. We assume here that the shift in Figure \\ref{fig.freq_freq} is a shift in $\\nu_{u}$; the conclusions drawn below do not change if we use the alternative, a shift in $\\nu_{\\ell ow}$, $\\nu_{h}$, and $\\nu_{b}$ (see \\S \\ref{sec.compare_timing}).  The strong, broad highest-frequency Lorentzian in the extreme island state of the atoll sources could previously be identified as either L$_u$ or L$_{hHz}$ based on its frequency of about 200 Hz. The rms fractional amplitude vs. frequency relations of 4U 0614+09, 4U 1608-52, and 4U 1728-34 gave a hint that this Lorentzian was L$_u$ \\citep{vstr03:apj596}. Now we observe the same shift in frequency for this Lorentzian as for the component which in other states is unambiguously L$_u$ (see Fig. \\ref{fig.freq_freq}), supporting the conclusion that it is indeed L$_u$. This conclusion is further strengthened by the results for the low luminosity bursters 1E 1724--3045, GS 1826--24, and SLX 1735--269. These sources, which always have a power spectrum similar to that in the extreme island state of the atoll sources, show a clear correlation between the frequency identified as $\\nu_{u}$, and $\\nu_{\\ell ow}$, $\\nu_{h}$, and $\\nu_{b}$ (Fig. \\ref{fig.freq_freq}). This would not be expected if this frequency is really $\\nu_{hHz}$, as in the atoll sources $\\nu_{hHz}$ remains approximately constant while the other frequencies change. Note also that with the addition of the extreme island state results for Aql X--1, the atoll sources now also show these same correlations in the extreme island state (when $\\nu_{u}$ is around 200 Hz). So, we identify the broad low frequency features previously reported by \\citet{wijnands98:apj507} from the 1998 outburst as manifestations of L$_u$.  The broad L$_{\\ell ow}$ component found in the extreme island states of the atoll sources and in the low luminosity bursters has now also been identified in the accreting millisecond pulsars. It was suggested that this component could be identified with the lower kilohertz QPO based on extrapolations of frequency-frequency relations \\citep[see][]{psaltis99:apj520,belloni02:apj572,vstr02:apj568,vstr03:apj596}, however, it was also noted that it could be a different component \\citep{vstr03:apj596}. Here we observe that if we make the relations of SAX J1808.4--3658 and XTE J0929--314 coincide with those of the atoll sources by multiplying $\\nu_{u}$ with 1.454, the $\\nu_{\\ell ow}$ vs. $\\nu_{u}$ relation of SAX J1808.4--3658 does indeed coincide with the atoll sources one, but the one $\\nu_\\ell$ point of SAX J1808.4--3658 is still offset from the atoll relation. We also have to multiply $\\nu_\\ell$ with a similar factor of about 1.5 to make the L$_\\ell$ point of SAX J1808.4--3658 fall on the atoll relation. If we apply the alternative shift (see \\S \\ref{sec.compare_timing}) we have to multiply $\\nu_{\\ell ow}$ with about 0.35, but $\\nu_\\ell$ with a different factor of about 0.8 to make the relations coincide. Therefore it seems probable that L$_{\\ell ow}$ and L$_\\ell$ are different components.  For SAX J1808.4--3658 the characteristic frequency of L$_{hHz}$ seems to be slightly correlated with the characteristic frequencies of the other components (see Fig. \\ref{fig.freq_freq}). The L$_{hHz}$ points also seem to lie on an extension of the L$_{\\ell ow}$ relation in Figure \\ref{fig.freq_freq}. In the other atoll sources L$_{\\ell ow}$ disappears and L$_{hHz}$ appears when the source goes from the extreme island to the island state. This happens at $\\nu_u \\sim350$ Hz in Figure \\ref{fig.freq_freq_shifted}. It might be that in SAX J1808.4--3658 L$_{\\ell ow}$ is present in the island state, but is not detected. It is not detected as a separate component when its characteristic frequency is in the range of $\\nu_{hHz}$, causing the correlation of L$_{hHz}$ with $\\nu_u$ in Figure \\ref{fig.freq_freq}. However note that the scatter in the L$_{hHz}$ relation of the atoll sources is large and that a similar scatter in SAX J1808.4--3658 might accidentally cause the correlation of L$_{hHz}$ with $\\nu_u$.   In the atoll sources 4U 0614+09, 4U 1608--52, and 4U 1728--34 the behavior of the band-limited noise in the lower left banana state is complex. First, the characteristic frequency of the band-limited noise increases from $\\sim0.3$ to $\\sim15$ Hz. At this stage the band-limited noise is broad. Then at $\\sim15$ Hz this noise component appears to ``transform'' into a narrow QPO whose frequency increases up to $\\sim50$ Hz, while what appears to be another band-limited noise component appears at lower characteristic frequencies. In our previous papers \\citep{vstr02:apj568,vstr03:apj596} we call the band-limited noise component that becomes a QPO, as well as the QPO it becomes, L$_{b}$ and the ``new'' broad band-limited noise appearing at lower frequency L$_{b2}$. In SAX J1808.4--3658 the same phenomenon seems to occur in groups 2 and 3. Based on this similar behavior the $\\sim15$ Hz Lorentzian can be identified as L$_{b}$ and the broad component at $\\sim10$ Hz as L$_{b2}$. However, in Figure \\ref{fig.freq_freq_shifted} for $\\nu_{u} \\gtrsim 600$ Hz the L$_{b}$ relation of SAX J1808.4--3658 follows the relation of the L$_{b2}$ component of the atoll sources. The L$_{b2}$ points of SAX J1808.4--3658 fall below that relation. This difference between L$_{b}$ in SAX J1808.4--3658 and in the other atoll sources can also be seen for $\\nu_b \\gtrsim 10$ Hz in Figure \\ref{fig.nu_vs_nub}. It seems that in this regime the frequency behavior of the band-limited noise in SAX J1808.4--3658 deviates from that in the other atoll sources.   In addition to the broad components L$_{b}$, L$_h$, L$_{\\ell ow}$, the low luminosity bursters \\citep{belloni02:apj572} and the atoll source 4U 1608--52 \\citep{vstr03:apj596,yoshida93:pasj45} show a narrow QPO that has $\\nu_{\\rm max}$ between $\\nu_{b}$ and $\\nu_{h}$. The QPO in 4U 1608--52 is not the same as that in the low luminosity bursters (called L$_{LF}$) but might be a sub-harmonic \\citep{vstr03:apj596}. Therefore we will call it L$_{LF/2}$. L$_{LF}$ and L$_{LF/2}$ are not only present in neutron star X-ray binaries but could also be identified in the black hole candidate GX 339--4 \\citep{belloni02:apj572,vstr03:apj596}. Also, a narrow QPO found in Cyg X--1 by \\citet{pottschmidt03:aa407} might be L$_{LF/2}$ \\citep{vstr03:apj596}. Until now these QPOs were only found at relatively low frequencies ($< 3$ Hz). In this work we find a QPO with a $\\nu_{\\rm max}$ between $\\nu_{b}$ and $\\nu_{h}$ at high frequencies ($\\approx 47$ Hz) in SAX J1808.4--3658 (groups 2 and 3), and at low frequencies (0.7--1.0 Hz) in XTE J1814--338 (group 29) and XTE J0929--314 \\citep[groups 23--25, previously reported by][]{markwardt02:apj575}. In Figure \\ref{fig.narrowqpo} we plot our results of the SAX J1808.4--3658, XTE J1814--338, and XTE J0929--314 on Figure 13 of \\citet{vstr03:apj596}. Based on the power law relations of Figure \\ref{fig.narrowqpo} we can identify the QPOs in XTE J1814--338 and XTE J0929--314 as L$_{LF/2}$ and the QPOs in SAX J1808.4--3658 as L$_{LF}$.    \\subsection{Timing-Color Correlations and States}  The timing properties of the accreting millisecond pulsars SAX J1808.4--3658, XTE J1751--305, XTE J0929--314, and XTE J1814--338 are very similar to those of the atoll sources. Also based on its correlated timing and color behavior SAX J1808.4--3658 can be classified as a classical atoll source (see \\S \\ref{sec.1808_colors}). In the time span for which timing results could be obtained, XTE J0929--314 and XTE J1814--338 showed little change in power spectra and colors. Therefore it is hard to classify these sources as atoll sources although their correlated timing and color behavior is consistent with that of a typical atoll source (see \\S \\ref{sec.0929_colors} and \\S \\ref{sec.1814_colors}). XTE J1751--305 is different from a classical atoll source with respect to the behavior of its hard color (see \\S \\ref{sec.1751}). In atoll sources in the island and extreme island states the characteristic frequencies and hard color are anti-correlated with each other, and decoupled from soft color and intensity, which are correlated. The hard color in XTE J1751--305 is instead strongly correlated with soft color and intensity, and is not anti-correlated with any of the characteristic frequencies.  To explain the one-to-one relation between soft color and intensity in an extreme island state of the atoll source 4U 1608--52 \\citet{vstr03:apj596} used the following scenario: In the extreme island state of 4U 1608--52 the disk is far away from the neutron star and cannot be seen in the PCA energy spectrum; the soft color therefore only depends on the spectral properties of the neutron star \\citep{gierlinski02:mnras337}. As the accretion rate (intensity) increases, the neutron star surface becomes hotter and the soft color increases \\citep{vstr03:apj596}. This scenario can also be applied to the horizontally extended extreme island state of SAX J1808.4--3658. In an energy spectral study of the 1998 outburst of SAX J1808.4--3658, \\citet{gierdonebarret02:mnras331} found that the soft component likely originates from the neutron star surface.  In the horizontally extended island state branches  that are observed in the atoll sources (including SAX J1808.4--3658), contrary to the horizontal branch of Z sources, the timing properties remain similar when the intensity changes. As noted by \\citet{olive03:apj583}, this aspect of the behavior can be explained by the scenario of \\citet{vdk01:apj561} for the parallel track phenomenon in the intensity versus lower kilohertz QPO frequency diagram that occurs in the lower left banana of the atoll sources \\citep[see e.g.][]{mendez00:texas}. If we look at these island state branches in SAX J1808.4--3658 during the 2002 outburst (groups 5--15) we observe two parallel tracks (see Figure \\ref{fig.parallel}), similar to those in the lower left banana of the atoll sources. This is a further indication that the same physical phenomenon is causing the source to drift between parallel tracks in the lower left banana of the atoll sources and along the island state branches as suggested by \\citet{vstr03:apj596}. We observe one parallel track from group 5 to 8 and the next from group 9 to 13. Both parallel tracks last 4 days; this is longer than the parallel tracks in the lower left bananas which usually last about a day.   \\subsection{Ingredients for a Physical Model} \\label{sec:physical}  Currently, to our knowledge, there is no physical model \\citep[see][for a discussion of models]{vstr03:apj596} that explains the mutual relations between the characteristic frequencies of the power spectral features (e.g. Figures \\ref{fig.freq_freq} and \\ref{fig.nu_vs_nub}). Here we point out some clues that might help to construct such a model. As described above in \\S \\ref{sec.identify}, in the lower left banana of the atoll sources (among which we now also classify SAX J1808.4--3658) the behavior of the low frequency components is complex; several timing features, additional to the ones observed in the island states, appear (see e.g. groups 2 and 3 of SAX J1808.4--3658). In Figure \\ref{fig.freq_freq_shifted} this happens at $\\nu_u \\approx 600$ Hz. The Z sources also show a complex low frequency behavior \\citep[see Figure 14 in][]{vstr03:apj596}, also at high $\\nu_u$. If $\\nu_u$ indeed represents the Keplerian frequency at the inner disk radius, a high $\\nu_u$ means that then the inner disk is close to the neutron star surface, so the more complex behavior might be caused by relativistic or magnetic effects that are stronger close to the neutron star. Note that the power spectra in the extreme island state of some sources also show extra timing features, but these are in the form of weak narrow QPOs that do not dominate the power spectra.  Because of the complex behavior of the low-frequency noise above $\\nu_u \\approx 600$ Hz in Figure \\ref{fig.freq_freq_shifted} we concentrate now on the more simple frequency relations below $\\nu_u \\approx 600$ Hz in Figure \\ref{fig.freq_freq_shifted}.  The low-frequency relations below $\\nu_u \\approx 600$ Hz can be fitted with power laws (see Table \\ref{tbl.plfits}). For these fits we use the same method as in \\S \\ref{sec.compare_timing}, which uses the FITEXY routine of Press \\citet{press92:comp274}. Note that although the $\\chi^2$/dof are relatively high, this is not due to systematic deviations from a power law but to scatter around it. The power law indexes of the $\\nu_{b}$, $\\nu_{h}$, and $\\nu_{\\ell ow}$ vs. $\\nu_u$ relations fall in the range 1.8--2.7, progressively steeper towards lower frequency components. The errors on the power law indexes of the L$_{LF}$ and L$_{LF/2}$ relations are large as these relations are only based on a few points, but also these power law indexes are consistent with this picture.  Here we use $\\nu_{\\rm max}$ as the characteristic frequency. Although $\\nu_{\\rm max}$ has proven to be of great practical value and may represent a real underlying physical frequency \\citep[e.g., the inner frequency of a disk annulus, see][]{belloni02:apj572}, this is not assured. Two other frequencies, that have physical interpretations in terms of a damped harmonic oscillator, are $\\nu_0$, the centroid frequency of the Lorentzian and $\\nu_{\\rm damped}$ \\citep[for a definition, see][]{titarchuk02:apj578}. These frequencies can be calculated from our fit results $\\nu_{\\rm max}$ and $Q$ \\citep[see][]{vstr03:apj596}. Note that using $\\nu_0$ instead of $\\nu_{\\rm max}$ does not give further insights in the shift between SAX J1808.4--3658, and the atoll sources (see \\S \\ref{sec.compare_timing}). For the narrow Lorentzians the shift factor will be the same as $\\nu_{\\rm max}$ is approximately equal to $\\nu_0$, for the broad features $\\nu_0$ is badly constrained and a shift cannot be determined. For the $\\nu_{h,0}$ vs. $\\nu_{u,0}$ relation of the atoll sources 4U 0614+09, 4U 1608-52, and 4U 1728-34 \\citet{vstr03:apj596} found a quadratic relation. A quadratic relation was also found for the $\\nu_{h,0}$ vs. $\\nu_{u,0}$ relations of several Z sources \\citep{stella98:apj492,psaltisetal99:apj520,jonker98:apj499,homan02:apj568}. The relativistic precession model \\citep{stella98:apj492} predicts a quadratic relation between the Lense-Thirring precession frequency and the Keplerian frequency of the inner disk; $\\nu_{LT} = 4.4\\times10^{-8} I_{45} M^{-1} \\nu_s \\nu_K^2$ Hz, where $I_{45}$ is the moment of inertia in units of $10^{45}$ g cm$^2$, $M$ is the mass of the neutron star in $M_{\\odot}$, and $\\nu_s$ is the neutron star spin frequency. However, the required $I_{45}/M$ values were too large for L$_{h}$ to be the fundamental Lense-Thirring precession frequency \\citep{psaltisetal99:apj520}. Also the $\\nu_{h,0}$ vs. $\\nu_{u,0}$ relations coincided for 4U 1728--34 and 4U 1608--52, which is not expected in this interpretation in view of their different spin frequencies (363 Hz and 620 Hz, respectively, assuming burst oscillations happen at the spin frequency). The relation of the frequency difference between the 410 Hz QPO and the $\\sim401$ Hz pulse spike of SAX J1808.4--3658 was found to be consistent with a quadratic relation by \\citet{wijnands03:nat424}. We note, that this frequency difference is always about half of $\\nu_{h,0}$ (see Table \\ref{tbl.410hz}).  Although the relations between the characteristic frequencies of the low frequency components and $\\nu_u$ are different between SAX J1808.4--3658, XTE J0929--314 and the other sources (see \\S \\ref{sec.compare_timing}), the  mutual relations between the low frequency components are the same for all atoll sources (Fig. \\ref{fig.nu_vs_nub}). The $\\nu_{h}$ vs. $\\nu_{b}$ relation is even the same for black hole candidates and atoll sources \\citep[][]{wijnands98:apj507,vstr02:apj568}. So, the $\\nu_{h}$ vs. $\\nu_{b}$ relation does not depend on spin, mass, magnetic field, or the presence of a surface. Note, that the frequency range of $\\nu_{h}$ and $\\nu_{b}$ is about a factor 10 lower for black hole candidates than for neutron star low mass X-ray binaries. Therefore L$_{b}$ and L$_{h}$ might be properties of the accretion disk that scale inversely with mass \\citep{wijnands99:apj514}.  The frequency correlations of SAX J1808.4--3658 and XTE J0929--314 are shifted compared to those of the other sources in a way that is most easily described as a shift in upper and lower kilohertz QPO frequency by about a factor 1.5 (see \\S \\ref{sec.compare_timing}). In most models $\\nu_{u}$ reflects the Keplerian frequency at the inner edge of the accretion disk (e.g. \\cite{miller98:apj508}). If this is indeed the correct interpretation, then for similar characteristic frequencies of the low frequency components, this inner radius of the disk can somehow get closer to the neutron star surface in the atoll sources, low luminosity bursters, and the accreting millisecond pulsars XTE J1751--305 and XTE J1814--338 than in the accreting millisecond pulsars SAX J1808.4--3658 and XTE J0929--314. It is unclear which physical parameter is the cause of this. In Table \\ref{tbl.msp} we list several characteristic parameters for the accreting millisecond pulsars. There is no obvious relation between the shift and any of the parameters. In particular, there is no relation to spin frequency: SAX J1808.4--3658, and XTE J0929--314 which share one set of frequency relations have spin frequencies of respectively 401 and 185 Hz, and the other sources which share another set of relations have spin frequencies that range from 314 to 620 Hz. Also the magnetic field strength $B$, which may be stronger for the accreting millisecond pulsars than for the atoll sources \\citep[see e.g.][]{wijnands03:nat424,strohmayer03:apj596}, is not a good candidate as not all accreting millisecond pulsars show the shift. The fact that the shift occurs by a factor close to $\\sim3/2$ suggests a possible relation with recently proposed parametric resonance models for kilohertz QPOs \\citep[e.g.][]{abramowicz03:pasj55}. In these models 2:3 frequency resonances between general relativistic orbital/epicyclic frequencies play a central role. However, similarly to the remarks about inner disk radius above, it is not clear what would set SAX J1808.4--3658 and XTE J0929--314 apart from the other sources.  This work was supported by NWO SPINOZA grant 08--0 to E.P.J. van den Heuvel, by the Netherlands Organization for Scientific Research (NWO), and by the Netherlands Research School for Astronomy (NOVA). This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center. We would also like to thank the referee, whose comments helped us to improve the paper."
    },
    "0410/astro-ph0410733_arXiv.txt": {
        "abstract": "We explore the possibility for the formation of Population III binaries. The collapse of a rotating cylinder is simulated with a three-dimensional, high-resolution nested grid, assuming the thermal history of primordial gas. The simulations are done with dimensionless units, and the results are applicable to low-mass as well as massive systems by scaling with the initial density. We find that if the initial angular momentum is as small as $\\beta \\approx 0.1$, where $\\beta$ is the ratio of centrifugal force to pressure force, then the runaway collapse of the cloud stops to form a rotationally-supported disk. After the accretion of the envelope, the disk undergoes a ring instability, eventually fragmenting into a binary. If the initial angular momentum is relatively large, a bar-type instability arises, resulting in the collapse into a single star through rapid angular momentum transfer. The present results show that a significant fraction of Pop III stars are expected to form in binary systems, even if they are quite massive or less massive. The cosmological implications of Population III binaries are briefly discussed. ",
        "introduction": "\\label{intro}  In present-day galaxies, it is known that more than half of main-sequence stars are found in binaries or low-order multiple systems not only for low-mass stars but also for massive stars (e.g., Tohline 2002; Kroupa 2004). Regarding the upper mass limit, a very massive binary of the order of $100M_\\odot$ is found in the Large Magellanic Cloud \\citep{Ostrov02}. Furthermore, it has been revealed that the binary fraction of pre-main-sequence stars is comparable to that of main-sequence stars \\citep{ Mathieu1994, Prosser94}. This implies that binaries possibly form prior to the pre-main-sequence phase.  Recent studies on the formation of Population III stars have shown that the first stars in the universe are likely to be as massive as $100-1000M_\\odot$ \\citep{BCL99, Abel00, Abel02} or to have formed in a bimodal initial mass function with peaks of $\\sim 1M_\\odot$ and several $100M_\\odot$ \\citep{NU99, NU01}. In light of present-day star formation, a portion of Population III stars appear able to form in binary systems. Population III binaries could play an important role for the thermal history of the universe through the formation of massive binary black holes (BHs), accreting binaries, and Population III Type Ia supernovae (SNe). For instance, the accreting binaries may emit ultraviolet radiation in a longer timescale than the lifetime of massive stars themselves and therefore can be a reionization source in the early universe \\citep{Umemura04}.  The binary formation for present-day stars has been studied through elaborate numerical simulations \\citep{Bonnell94,Boss86,Tsuribe99,MH03,NL03,MTM04}. However, no study has been made on the formation of Population III binaries. So far, studies of Population III protostars have focused on the core collapse of single stars which proceeds through the cooling by primordial hydrogen molecules \\citep{PSS83,ON98,Omukai00}. The previous studies imply that (1) the initial stellar core is on the order of $M_{\\rm core}=10^{-2} M_\\odot $, which is not so different from the present-day protostars, (2) the mass accretion rate is as high as a few tens of $c_s^3/G$ with the sound speed $c_s$, which is numerically $\\dot{M}\\approx 10^{-3} -10^{-1} M_\\odot {\\rm yr}^{-1}$ for $\\approx 100$K gas because no coolant works to cool down to less than $100$K in primordial gas, and (3) the thermal history follows the evolution with an adiabatic exponent of $\\gamma \\approx 1.1$ until reaching the central density of $n \\approx 10^{20}{\\rm cm}^{-3}$, above which hydrogen molecules are dissociated and therefore the temperature increases with $\\gamma=5/3$.  On the basis of these findings, we study the formation of Population III binary stars. We perform a three-dimensional hydrodynamic simulation with nested grids in order to resolve the core collapse and the subsequent fragmentation. In the next section, the numerical model is described. In \\S \\ref{Results}, we present the numerical results for the collapse and fragmentation of Population III stellar core. \\S \\ref{Discussion} is devoted to the discussion of the cosmological implications of Population III binaries. ",
        "conclusions": "\\label{Discussion}  We have simulated the collapse of a rotating cylinder while assuming the primordial gas thermal history. As a result, we have found that to form a binary system, the formation of a rotationally supported disk is indispensable, which is realized only by a lower rotation parameter. The rotation-supported disk undergoes ring-mode instability, eventually fragmenting in to binary cores. In contrast, the case with a relatively high rotation parameter leads to a bar-mode instability and the angular momentum efficiently transferred by the bar, resulting in the collapse into a single stellar core. These results show that the initial conditions with less angular momentum are more prone to binary formation and therefore imply that a significant fraction of Population III stars are expected to form in binary systems.  In the present simulation, we have started with high initial density. If we use the scaling for mass and linear size and renormalize the physical quantities by a lower initial density such as $n_{\\rm c0} \\la 1 \\times 10^{5}{\\rm cm}^{-3}$, then the simulation predicts the formation of a massive binary with several hundred $M_\\odot$ and a separation of $\\approx 100$AU. \\citet{NU01} showed that the mass of Population III stars could bifurcate according to the initial filament density. If the initial density is above the critical density of H$_2$ molecules, $n_{\\rm c0} > 10^{5}{\\rm cm}^{-3}$, then the fragment mass is peaked around $1M_\\odot$, whereas the peak is shifted to $\\approx 100 M_\\odot$ for $n_{\\rm c0} \\la 1 \\times 10^{5}{\\rm cm}^{-3}$.  The present simulations show that the binary formation can occur in the wide mass range of 1 to several $\\times 100 M_\\odot$, if we start with the initial density of 10$^{3}$ -- 10$^{12}$ cm$^{-3}$.   The recent theoretical analyses of the evolution of metal-free stars predict that the fate of Population III stars can be classified as follows (e.g., Heger \\& Woosley 2002): \\begin{enumerate} \\renewcommand{\\labelenumi}{\\arabic{enumi}.} \\item A star of $1M_\\odot \\la M \\la 8 M_\\odot$ evolves into a white dwarf after asymptotic giant branch mass loss.  \\item A star of $8M_\\odot \\la M \\la 25 M_\\odot$ results in a Type II SN.  \\ item A star of $25M_\\odot \\la M \\la 140 M_\\odot$ probably collapses into a BH.  \\item A star of $140M_\\odot \\la M \\la 260 M_\\odot$ is partly or completely disrupted by the electron-positron pair instability.  \\item A star with mass of $M \\ga 260 M_\\odot$ collapses completely to a BH without ejecting any heavy elements. \\end{enumerate} Combining such Population III stellar evolution with the present results, we may expect a wide variety of Population III binaries. If a binary is as massive as $\\ga 300 M_\\odot$, it can result in a binary BH. Recently, \\citet{Belczynski04} argued that a Population III binary BH could be the intensive gravitational wave burst sources, which may be detected by the Laser Interferometer Gravitational-Wave Observatory. If there is a slight mass difference in a massive binary, a system composed of a $\\sim 100 M_\\odot$ Population III star and several $100 M_\\odot$ BHs may form. Then the mass ejected from a massive star can accrete onto a companion BH, resulting in a possible source of ionizing radiation in the early universe \\citep{Umemura04}. The Population III BH accretion may be an important clue for the early reionization that is suggested by the {\\it Wilkinson Microwave Anisotropy Probe} \\citep{RO04a,RO04b}. If a low-mass Population III binary forms, Population III Type Ia SN might occur. In this case, heavy elements synthesized in a more rapidly evolving star can pollute a companion star by stellar wind. The evolution is likely to be sensitive to the metallicity-dependent opacity of accretion flow. Population III Type Ia SNe seem worth studying in the future, because they may play an important role in the thermal history in the universe."
    },
    "0410/astro-ph0410219_arXiv.txt": {
        "abstract": "A toy model for gamma-ray burst-supernovae (GRB-SN) is discussed by considering the coexistence of baryon poor outflows from black holes (BHs) and a powerful spin-connection to surrounding disk, giving rise to consistent calorimetry as described by van Putten (2001a) in a variant of the Blandford-Znajek process (BZ, 1977). In this model the half-opening angle of the magnetic flux tube on the horizon is determined by the mapping relation between the angular coordinate on the BH horizon and the radial coordinate on the surrounding accretion disk. The GRB is powered by the baryon poor outflows in the BZ process, and the associated SN is powered by very small fraction of the spin energy transferred from the BH to the disk in the magnetic coupling (MC) process. The timescale of the GRB is fitted by the duration of the open magnetic flux on the horizon. It turns out that the data of several GRB-SNe are well fitted with our model. ",
        "introduction": "As is well known, the Blandford-Znajek (BZ) process is an effective mechanism for powering jets from quasars and active galactic nuclei (Blandford {\\&} Znajek 1977; Rees 1984). In the BZ process energy and angular momentum are extracted from a rotating balck hole (BH), and transferred to the remote astrophysics load by the open magnetic field lines. Recently, much attention has been paid on the issue of the long duration gamma-ray bursts (GRBs) powered by the BZ process (Paczynski 1993, 1998; Meszaros {\\&} Rees 1997). A detailed model of GRBs invoking the BZ process is given by Lee, Wijers, {\\&} Brown 2000, hereafter L00), and they concluded that a fast-spinning BH with strong magnetic field $\\sim 10^{15}G$ can provide energy $\\sim 10^{53}ergs$ to power a GRB within 1000 seconds.   Recently the observations and theoretical considerations have linked long duration GRB with ultra-bright type Ib/c supernovae (Galama et al. 1998; Bloom et al. 1999; Galama et al. 2000). The first candidate was provided by SN1998bw and GRB980425, and the recent HETE-II burst GRB030329 has greatly enhanced the confidence in this association (Stanek et al. 2003; Hjorth et al. 2003).  Not long ago Brown et al. (2000, hereafter B00) worked out a specific scenario for GRB-SN connection. They argued that the GRB is powered by the BZ process, and the SN is powered by the energy dissipated into the disk through closed magnetic field lines coupling the disk with the BH. The latter energy mechanism is referred to as the magnetic coupling (MC) process, which is regarded as one of the variants of the BZ process (Blandford 1999; Putten 1999; Li 2000, 2002; Wang, Xiao {\\&} Lei 2002, hereafter W02). It is shown in B00 that about $10^{53}ergs$ are available to power both GRBs and SNe. However, they failed to distinguish the fractions of the energy for these two objects.  More recently, van Putten et al. (van Putten 2001a; van Putten \\& Levinson 2003; van Putten et al. 2004) worked out a poloidal topology for the open and closed magnetic field lines, where the separatrix on the horizon is defined by a finite half-opening angle. The duration of GRB is set by the lifetime of rapid spin of the BH. It is found that GRB and SN are powered by a small fraction of the BH spin energy. This result is consistent with the observations, i.e., the duration of GRB with tens of seconds, the true GRB energies distributed around $5\\times 10^{50}ergs$ (Frail et al. 2002), and the aspherical SNe kinetic energies with $2\\times 10^{51}ergs$ ( Hoflich et al. 1999).  In this paper we propose a toy model for GRB-SN by considering coexistence of the BZ and MC processes (CEBZMC), where the configuration of the magnetic field is based the works of Wang et al. and van Putten et al. (W02; Wang et al. 2003, hereafter W03; van Putten 2001a; van Putten et al. 2004). The transferred energy and the duration for GRBs are calculated in the evolving process of the half-opening angle of the magnetic flux on the BH horizon, while the transferred energy for SNe is calculated in the MC process. Thus the association of GRBs and SNe is explained reasonably and consistently.  This paper is organized as follows. In \\S \\ 2 the configuration of the magnetic field is described with the half-opening angle $\\theta _{BZ} $ of the magnetic flux tube on the horizon, and we find that $\\theta _{BZ} $ can evolve to zero in some range of the power-law index $n$ of the magnetic field on the disk. In \\S \\ 3 the toy model for GRB-SN based on BH evolution is described in the parameter space consisting of the BH spin $a_ * $ and the power-law index $n$. It is shown that the half-opening angle $\\theta _{BZ} $ plays an important role of determining (i) the powers for GRBs and SNe, (ii) the duration of GRB. In \\S \\ 4 we calculate the duration of GRB and the energies of GRB-SN, and compare with the other models of GRBs invoking the BZ process. It is found the duration of GRB obtained in our model is shorter than that of the other models. It turns out that this model is in excellent agreement with the observations, owing to considering the MC effects, while too much energy than the needed energy of GRB is produced in the other models without the MC effects. Finally, in \\S \\ 5, we summarize our main results and discuss some problems concerning our model. Throughout this paper, the geometric units $G = c = 1$ are used. ",
        "conclusions": "In this paper we discuss a toy model for GRB-SN, where the energy for GRBs and SNe are powered by the BZ and MC processes, respectively, and the duration for GRB is estimated by the lifetime of the half-opening angle $\\theta _{BZ} $. The energy extracted in the BZ process for GRBs is much less than that extracted in the MC process for SNe. This result is consistent with the observations: the true energy of GRB $E_\\gamma = 5\\times 10^{50}ergs$ and the aspherical SNe of kinetic energy $E_{SN} = 2\\times 10^{51}ergs$. For a set of GRBs the observed true energy $E_\\gamma $ and duration $T_{90} $ can be well fitted by taking adequate values of the power-law index $n$ and magnetic field $B_H $. However, there are still several issues related to this model.  (i) In this paper we take the evolving path of the \\textbf{RP} in the region \\textbf{IA} to describe the association of GRB-SN as shown in Figure 5. In this case the GRB will terminate before the SN event stops. In fact the evolving path in region\\textbf{ IB} can be also used to describe the association of GRB-SN, and it implies that the \\textbf{RP} will arrive at the characteristic curve of $\\tilde {P}_{MC} = 0$ ahead of reaching the characteristic curve of $\\tilde {P}_{BZ} = 0$. However, we think that the GRB will still terminate before the SN event stops, because the disk will be destroyed during the explosion of the SN, and the BZ process can not work without the magnetic field supported by its surrounding disk.  (ii) In order to highlight the effects of the BZ and MC processes in powering GRB-SN, we neglect the effects of the disk accretion in this model. In fact, disk accretion plays a very important role in BH evolution. The slowing down of the BH spin will be delayed significantly since the accreting matter delivers remarkable amount of energy and angular momentum to the BH. Strictly speaking, disk accretion cannot be halted steadily by the MC process. The MC effects on disk accretion is a very complicated, and it should be treated dynamically.  (iii) According to Kruskal-Shafranov criteria (Kadomtsev 1966), the screw instability will occur if the toroidal magnetic field becomes so strong that the magnetic field lines turns around itself about once. In our latest work (Wang et al. 2004), we argued that the state of CEBZMC always accompanies the screw instability. In this model the half-opening angle $\\theta _{BZ} $ is related to the infinite radial parameter $\\xi $ by equation (\\ref{eq4}). We expect to derive a greater $\\theta _{BZ} $ by considering the restriction of the screw instability to the MC region. Thus the ratios $\\eta _{GRB} $ and $\\eta _{SN} $ will be changed. We shall improve this model in future."
    },
    "0410/astro-ph0410518.txt": {
        "abstract": "{ Studies of distant galaxies have shown that ellipticals and large spirals (Schade et al. 1999; Lilly et al. 1998) were already in place 8 Gyr ago, leading to a very modest recent star formation (Brinchmann and Ellis, 2000) in intermediate mass galaxies (3-30 $10^{10}$ $M_{\\odot}$). This is challenged by a recent analysis (Heavens et al. 2004) of the fossil record of the stellar populations of $\\sim$ $10^{5}$ nearby galaxies, which shows that intermediate mass galaxies formed or assembled the bulk of their stars 4 to 8 Gyr ago. Here we present direct observational evidence supporting this findings from a long term, multi-wavelength study of 195 z$>$ 0.4 intermediate mass galaxies, mostly selected from the Canada France Redshift Survey (CFRS). We show that recent and efficient star formation is revealed at IR wavelengths since $\\sim$ 15\\% of intermediate mass galaxies at z$>$ 0.4 are indeed luminous IR galaxies (LIRGs), a phenomenon far more common than in the local Universe. The star formation in LIRGs is sufficient in itself to produce 38\\% of the total stellar mass of intermediate mass galaxies and then to account for most of the reported stellar mass formation since z=1. \\\\ Observations of distant galaxies have also the potential to resolve their star formation and mass assembly histories. The high occurrence of LIRGs is easily understood only if they correspond to episodic peaks of star formation, during which galaxies are reddened through short IREs (infrared episodes).  We estimate that each galaxy should experience 4 to 5 $\\times$ $(\\tau_{\\rm IRE}/0.1Gyr)^{-1}$ IREs from z=1 to z=0.4, $\\tau_{\\rm IRE}$ being the characteristic timescale. An efficient and episodic star formation is further supported by the luminosity-metallicity relation of z$\\sim$ 0.7 emission line galaxies, which we find to be on average metal deficient by a factor of $\\sim$ 2 when compared to those of local spirals. We then examine how galaxy IREs can be related to the emergence at high redshift of the abundant population of galaxies with small size (but not with small mass), blue core and many irregularities. We show that recent merging and gas infall naturally explain both morphological changes and episodic star formation history in a hierarchical galaxy formation frame.\\\\ We propose a simple scenario in which 75$\\pm$25\\% of intermediate mass spirals have recently experienced their last major merger event, leading to a drastic reshaping of their bulges and disks during the last 8 Gyrs. It links in a simple manner distant and local galaxies, and gives account of the simultaneous decreases during that period, of the cosmic star formation density, of the merger rate, and of the number densities of LIRGs, compact and irregular galaxies, while the densities of ellipticals and large spirals are essentially unaffected. It predicts that 42, 22 and 36\\% of the IR (episodic) star formation density is related to major mergers, minor mergers and gas infall, respectively. ",
        "introduction": "Understanding how and when galaxies and the stars in them were formed is still a considerable challenge for astrophysicists. A global picture of galaxy evolution during the last 8 Gyrs has emerged from the results of all-inclusive galaxy surveys.  The apparently unchanged number density of ellipticals and of large spirals at z$\\sim$ 0.7 (Schade et al. 1999; Lilly et al. 1998; Conselice et al, 2004a) has lead to a picture for which most of the recent star formation  occurred in small mass galaxies (Cowie et al. 1996; Brinchman and Ellis, 2000, hereafter BE2000). This picture has been called \"galactic downsizing\" to illustrate its apparent contradiction to expectations from hierarchical models of galaxy formation. It is only consistent with a modest fraction of stars recently formed, since present-day stars are mostly locked in spheroids (bulge of spirals and ellipticals, Fukugita, Hogan \\& Peebles, 1998).  This is challenged by the evolution of the cosmic stellar mass density which is based on a detailed modelling of galaxy spectral energy distributions (Cole et al. 2001; Dickinson et al. 2003,; Fontana et al. 2003; Drory et al. 2004; Rudnick et al. 2003). Even if they suffer from various uncertainties (cosmic variance, assumptions on stellar populations, metal abundances, extinction and initial mass function), all these studies predict that 30 to 50\\% of the mass locked in stars in present day galaxies actually condensed into stars at $z<1$. These estimates agree with integrations of the universal star formation history, if it properly accounts for the light re-emitted by dust at IR wavelengths (Flores et al. 1999; Chary and Elbaz, 2001). Heavens et al. (2004, hereafter, H2004) have derived a peak of star formation at modest redshifts (z=0.5-0.6), mostly due to intermediate mass galaxies (3-30 $10^{10}M_{\\odot}$) which have assembled/formed 60\\% of their stellar masses during the last 8 Gyrs (from integrations of their Figure 2). There is now a reasonable agreement that a significant part of the present-day stars (30\\% to 50\\%) was formed at recent epochs. %A possible explanation of the apparent discrepancy with BE2000 is provided by an examination of the Dickinson et al's Figure 7. While the BE2000 points in their highest redshift bins ($0.4-0.8$; $0.8-1$) fit well with other findings, their point in the lowest redshift bin (0.2-0.4) might be affected by small statistics at low redshift (the so-called cone effect) and to some lack of z$\\sim$ 0.3 galaxies in several CFRS fields (see Crampton et al. 1995).  Where are these relatively recent stars born and how they have been formed ? We examine the major evolutionary trends since z=1 from all inclusive surveys. Five to eight Gyrs ago, the galaxy population included $ 70$\\% starbursts ( defined as W$_O$([OII])$> 15\\AA$ ), a fraction which has sharply decreased to 17\\% in the local Universe (e.g. Hammer et al. 1997). Both UV (Lilly et al. 1996) and IR (Flores et al. 1999) comoving luminosity densities have been declining rapidly since $z=1$.  \\subsection{ Most of the star formation since z=1 is related to the IR luminosity density}  At IR wavelengths the comoving luminosity density was 70 times larger at $z=1$ than it is today (Elbaz et al. 1999; Elbaz and Cesarsky, 2003). At $z=1$ it is mostly due to luminous IR galaxies (LIRGs, defined as L$_{IR} > 10^{11} L_\\odot$ following Sanders and Mirabel, 1996), which are mainly powered by star formation  (Fadda et al 2002). LIRGs form stars at very high rates (from  20 to several hundred of $M_\\odot/yr$) and are responsible for the bulk of the cosmic infrared background (Elbaz et al. 2002). An early conclusion from studies of mid infrared emission and SEDs of galaxies is that from 50 to 66\\% of the CSFD at $z\\sim 0.7$ is related to far infrared emission (Flores et al. 1999), mostly due to bright mid-IR sources (LIRGs). The large impact of star formation revealed by the IR has been recently confirmed by Spitzer counts of 24 $\\mu$m sources (Papovich et al. 2004).  Distant LIRG morphologies encompass those of disks, irregulars and LCGs, with a noticeable fraction of them showing signs of merging or interactions (Zheng et al. 2004a). If sustained at the observed high rates, their star formation can double their mass in less than 1 Gyr (Franceschini et al. 2003; Zheng et al. 2004a). Liang et al. (2004b) have derived their extinction, star formation rate and metal abundances from a detailed analysis of their high quality spectra from VLT. Distant LIRGs present a metal content less than half of that of local bright disks. Liang et al. (2004b) and Zheng et al. (2004a) have suggested that many distant disks are still forming a significant fraction of their stellar mass content.  \\subsection{ Can we relate the UV luminosity density decline to the evolution of the merging rate ?}  Peculiar galaxies are found to be the most rapidly evolving galaxy population and to be responsible for most of the decline of the UV luminosity density (Brinchman et al. 1998; Lilly et al. 1998, BE2000, Conselice et al, 2004b). Besides showing unusual or irregular morphologies, peculiar galaxies generally display small half light radii (Lilly et al. 1998), while their stellar masses are comparable to those of large, more regular galaxies (BE2000). It has been often suggested that the morphologies of peculiar galaxies are related to interactions and that merging could play a major role in their number density decline. This is broadly consistent with the large expected amount of merging events during the last 8 Gyrs. From number counts of pairs up to $z=1$, Le F\\`evre et al. (2000) estimated that each luminous (L*) galaxy experiences 1.1 to 1.8 $\\times$ (7 10$^8$ yrs/$\\tau_{merger}$) major merging event (with a companion with half and one fourth of its luminosity, respectively). In this study the characteristic time over which the merger is completed, $\\tau_{merger}$ was assumed to range from 0.4 to 1 Gyr following expectations from simulations (Hernquist \\& Mihos, 1995). A recent study made in the near IR (Bundy et al. 2004) leads us to suspect a slower merger history, and accounting for the various selection effects (colour \\& relative spatial resolution) would decrease the Le F\\`evre et al. estimates to half their original values. Merging appears still to be an important ingredient in galaxy evolution at z$<$ 1, and this matches rather well with a recent stellar mass production.  Several clues also favor a merger origin for many peculiar galaxies. In the CFRS, Hammer et al. (2001) have studied in detail the most compact ones ($r_{50}<$ 3.5 kpc), hereafter called luminous compact galaxies (LCGs) which show blue U-V colours but not necessarily blue B-K colours. LCGs are found to be dusty enshrouded starbursts (SFR averaged to 40 M$_{\\odot}/yr$) superimposed on a mix of metal-rich old stellar population and intermediate age stars. Hammer et al (2001) suggested that LCGs could be remnants of mergers or interactions, and could be an important stage in the bulge formation in many spirals. Detailed studies of few local counterparts by Barton \\& Zee (2001) and by Ostlin et al. (2001) have revealed tidally triggered central starbursts or complex velocity fields typical of completed mergers, which also explain their narrow emission lines. There are also theoretical grounds which support that merger remnants could be observed in large numbers in the distant Universe. Although after a major merger, the dynamical relaxation time is short (few $10^{7}$ yrs, Barnes \\& Hernquist, 1992), hydro-dynamical models (Tissera et al. 2002; see also Baugh et al. 1996) show that enhanced star formation can be sustained from several $10^{8}$ yrs to $\\sim$ 2 $10^{9}$ yrs after the collision, a period during which formation of stars in the surrounding disk can be accommodated.  %Peculiar (small irregular, mergers and compact) galaxies are found to be the most rapidly evolving galaxy populations and responsible of most of the decline of the UV luminosity density (Guzman et al. 1997; Brinchman et al. 1998; Lilly et al. 1998). The important population of faint \\& blue compact galaxies at intermediate redshifts has been revealed from HDF studies (Guzman et al. 1997; Phillips et al. 1997) as a mixture of HII like galaxies and central starbursts in disk like galaxies. Due to its shallower selection, CFRS has detected their luminous and massive counterparts (see e.g. Guzman et al. 2003). These have been classified by Hammer et al (2001) as luminous compact galaxies (LCGs) because while they display blue rest frame U-V colours, their rest-frame B-K colours take the whole range from Irr to early type values. Hammer et al (2001) found that they are dusty enshrouded starbursts (SFR averaging to 40 M$_{\\odot}/yr$) superimposed to a mix of metal rich old stellar population \\& intermediate age stars. They have suggested that LCGs could be remnants from mergers or interactions, and could be an intermediate stage in the bulge formation, before the completion of star formation in the surrounding disk. An active formation of the bulge through tidally-triggered star formation is supported by a study of few local counterparts by Barton \\& van Zee (2001), and this also reports for their narrow emission lines. Fabry Perot observations of two LCGs at low redshift have revealed very irregular velocity fields which probably originate from mergers (Ostlin et al. 2001). Although after a major merger, the dynamical relaxation time is short (few $10^{7}$ yrs, Barnes \\& Hernquist, 1992), hydro-dynamical models (Tissera et al. 2002; see also Baugh et al. 1996) show that enhanced star formation can be sustained from several $10^{8}$ yrs to $\\sim$ 2 $10^{9}$ yrs after the collision, a period during which the formation of stars in the surrounding disk can be accommodated.\\newline  \\subsection{The present study}  Thus there are good reasons to suspect that merging probably has an important impact on the observed recent formation of stars. Because a significant part of the recent star formation is related to dust-enshrouded galaxies, it calls for a panchromatic study of galaxy properties. This study is based on long-term follow-up observations of the CFRS by our team, using HST, ISO, VLA and VLT. It has the goal to link the morphological changes seen in optical wavelengths to the star formation revealed at IR wavelengths, and, more ambitiously, to be able to see an overall picture of galaxy evolution over the last 8 Gyr. It targets 195 $z> 0.4$ galaxies, mostly from the sample of the CFRS which is probably the most widely studied survey of galaxies for its statistical properties. CFRS is essentially complete up to z$\\sim$ 1, encompassing all luminous ($M_B<-20$) galaxies (Hammer et al. 1997) with stellar masses from 3 to 30 10$^{10} M_{\\odot}$, those being responsible for the bulk of the star formation history reported by H2004. There is a general consensus that intermediate mass galaxies dominate the stellar mass budget at $z < 1$ (80\\% of the total mass according to BE2000, 65\\% from integrations of H2004).  In Section 2 we present the data and a comparison of the various estimates of star formation rate is given in Section 3. A solid estimate of the stellar mass fraction formed in the past 8 Gyrs is provided in Section 4, as well as new insights on the star formation histories of individual galaxies. Section 5 describes the major morphological changes during this elapsed time, which are linked to the galaxy star formation history in Section 6, leading to a simple scenario which accounts for all changes in the intermediate mass galaxy population, in the framework of the hierarchical galaxy formation theory. $\\Lambda$CDM cosmology ($H_0$=70 km$s^{-1}Mpc^{-3}$, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$) is adopted throughout the paper. ",
        "conclusions": "We have re-examined here the properties of distant galaxies mostly on the basis of a panchromatic follow-up study of the CFRS. We find that: \\begin{itemize} \\item SFRs based on the [OII]$\\lambda$3727 emission line lead to severely underestimated values, which can lead to misleading results on the star formation history of distant galaxies \\item 15\\% of intermediate mass galaxies at 0.4$<$ z $<$ 1 are indeed LIRGs and the associated star formation density is sufficient in itself to form 38\\% of the total stellar mass in present-day intermediate mass galaxies \\item all emission line galaxies (starbursts and LIRGs) at z$\\sim$ 0.7 have gas phases which are metal deficient by a factor 2 when compared to those of local spirals \\item most distant field galaxies (but E/S0) show significant star formation in their 1kpc radius central cores \\item  optical/spectral properties of LIRGs are so similar to those of other galaxies that only IR measurements are able to describe how the star formation density is distributed between the different galaxy types \\item the only way to interpret the high occurrence of LIRGs is to associate them to episodic star formation histories for most galaxies, in agreement with the hierarchical galaxy formation scenario \\end{itemize}  These episodic and violent star formation events can be simply associated to major changes in galaxy morphologies. We have then tested three galaxy formation scenarios against eight rather robust features of galaxy evolution. The best scenario accounts for a recent formation of bulges in spirals through merging or tidally triggered central bursts, followed by an inside out formation of the stars in disks. In this scenario, at least half of the bulge stellar content was made earlier in their progenitors, before the last major phase of accretion. It then provides a simple origin for the reported recent stellar mass production (30 to 50\\% of the stellar mass formed since z=1) and for a recent peak in the cosmic star formation history (H2004). It is consistent with the simultaneous decreases, during the last 8 Gyrs, of the UV and IR luminosity density, of the merging rate, and of the number densities of LIRGs and of LCGs. It also naturally accounts for all the major changes of the galaxy morphologies in the last 8 Gyrs, as well as for an approximately unchanged density of large disks at $z\\sim 1$. Because of the striking similarities between properties of galaxies at different luminosities (average $O/H$ abundances, fraction of merging events and of compact galaxies), possibly most present-day spirals, independently of their mass range, have recently experienced their last major phase of the building of both their bulges and their disks. More than a third of the present-day stellar mass formed at $z < 1$ is associated with more than half of the stellar mass formed in spirals during violent events which have shaped them as the Hubble sequence describes them today.\\newline  All the important numbers used here to establish our scenario (e.g. fraction of LIRGs and of various morphological types, SFR, stellar masses etc...) have been systematically cross-checked and found in good agreement with the litterature. We also find that biases related to the I band selection criterion used in the CFRS provide minor effects. This is likely due to the fact that dust-enshrouded starbursts (such as LIRGs) show optical properties which mimic those of normal starbursts, in such a way that optical observations are generally unable to distinguish them. This is well reported in our scenario: the addition of a highly reddened burst to a normal high-z galaxy has no severe impact on its optical continuum. The way IR light impacts our optical view of the galaxy evolution is then rather subtle and it explains well the often associated wording \"hidden star formation\". Without observing in the IR and accounting only from the UV light, all distant galaxies apparently show modest star formation rate and stellar mass increase. %However, by using moderate spectral resolution and very large S/N spectroscopy available from large telescopes, estimates of the star formation from Balmer lines, after proper correction of underlying absorption (Flores et al. 2004; Liang et al. 2004a), provide values which are broadly consistent with IR results. \\newline  A significant number of recent major mergers is qualitatively consistent with the dramatic increase of type 1 and 2 AGN out to z=1 reported by Miyaji, Hasinger \\& Schmidt (2000).  Major mergers are particularily efficient in forming large bulges which can account for the numerous population of present-day early-type spirals (see also Zheng et al, 2004b). The link between distant and nearby galaxies could be further investigated: what is the relative impact of both star formation and mass aggregation at z$<$ 1 for present-day spirals and ellipticals ? Could we match our scenario with the Milky Way and M31 past histories ? While it is unlikely that a major merger has recently occurred in the Milky Way, the M31 system appears to be much more complex and the significant intermediate age population in its halo supports the idea of a recent galaxy merger (Brown et al, 2003). It is also necessary to predict more accurately the relative fraction of star formation occurring during major mergers, minor mergers, interactions and gas infall, respectively. This requires unambiguous estimates of the actual merging rate and of the fraction of LCGs which are end products of merging events. We believe that estimates of velocity fields of distant galaxies with the VLT (with FLAMES/GIRAFFE) will address these issues, as well as the evolution of the Tully Fisher relationship. Large follow up of the deepest fields studied by HST and Spitzer should be done, using moderate spectral resolution spectroscopy to derive SFR, extinction and chemical gas properties, as well as velocity fields provided by 3D spectroscopy."
    },
    "0410/hep-ph0410176_arXiv.txt": {
        "abstract": "\\noindent In the {\\it Top-down} scenarios, the decay of super-heavy particles ($m \\sim 10^{12-16} $GeV), situated in dark-matter halos not very far from our Galaxy, can explain the ultra-high-energy (UHE) cosmic-ray spectrum beyond the Griesen-Zatasepin-Kuzmin  cut-off. In case the dynamics of this decay is governed by the minimal supersymmetric standard model, a major component of the UHE cosmic-ray flux at PeV-EeV energies could be given by the lightest {\\it neutralino} $\\chi^0_1$, that is the lightest stable supersymmetric particle. Then, the signal of UHE $\\chi^0_1$'s on earth might emerge over the interactions of a comparable neutrino component. We compute the event rates for the resonant production of {\\it right selectrons} ($\\tilde{e}_R$) and {\\it right squarks} ($\\tilde{q}_R$) in mSUGRA, when UHE neutralinos of energy $E_{\\chi}\\gsim 10^5$ GeV scatter off electrons and quarks in an earth-based detector like IceCube. When the resonant channel dominates in the total $\\chi^0_1e,\\chi^0_1q$ scattering cross section, the only model parameters affecting the corresponding visible signal rates turn out to be the  physical masses of the resonant right-scalar and of the lightest neutralino. We compare the expected number of supersymmetric events with the rates corresponding to the expected Glashow $W$ resonance and to the continuum UHE $\\nu N$ scattering for realistic power-law spectra. We find that the event rate in the leptonic selectron channel is particularly promising, and can reach a few tens for a one-year exposure in IceCube.  Finally, we note that UHE neutralinos at much higher energies (up to hundreds ZeV) may produce {\\it sneutrino} ($\\tilde{\\nu}$) resonances by scattering off relic neutrinos in the Local Group hot dark halo. The consequent $\\tilde{\\nu}$ {\\it burst} into hadronic final states could mimic  {\\it Z-burst} events, although with quite smaller conversion efficiency. ",
        "introduction": "The presence of a ultra-high-energy (UHE) component in the cosmic-ray spectra beyond the Griesen-Zatasepin-Kuzmin cutoff (GZK) \\cite{gzk}, as revealed initially by the experiment Fly's Eye \\cite{flys_eye}, subsequently by AGASA \\cite{agasa}, and marginally by Hires \\cite{hires}, is presently an open problem in high-energy astrophysics. The highest-energy particles in the cosmic-ray spectrum are apparently protons rather than photons, because of their characteristic hadronic-shower fluorescence light shapes. Beyond the GZK energies, protons, slowed down by the cosmic Black Body Radiation, have energy-loss length well below 50 Mpc. This implies that possible sources of these extremely-high-energy particles should be within our Galaxy, or inside our Local Group of Galaxies or nearby Clusters. However, there are no such powerful sources in Milky Way or nearby Virgo Group, that are correlated with the arrival map of the UHE cosmic rays (UHECR). There is a timid, and yet unconfirmed, clustering of AGASA data along the Super-Galactic Plane, not sufficient to probe the UHECR confinement in observable GZK volumes. On the contrary, UHECR events exhibit isotropy, a signature reminiscent of most distant cosmic edges. It is not clear whether the UHECR spectrum is really showing a GZK cut-off \\cite{agasa,hires}. In order to make our UHECR understanding even more confused, there are now first evidences of clustered UHECR events by AGASA correlated with known BL Lacs sources \\cite{Gorbunov02}. This UHECR-BL Lacs connection calls for the UHE Blazing Jet of active galactic nuclei (AGN) as a source of particle acceleration. Those Bl Lacs sources are mostly at distances above GZK lengths.  Although the GZK puzzle is still experimentally not completely settled, in the past and also recently, there have been quite a few suggestions about possible sources and mechanisms for accelerating particles to such extreme energies \\cite{suggestn,flux_rev}. \\\\ A class of proposals, generically called {\\it Top-down} scenarios, tries to explain the high-energy part of the spectrum through the decay of super-massive particles ($m \\sim 10^{12-16} \\rm{GeV}$) \\cite{massive,ber,susy_oth,toldra,drees}, or by their eventual binary collapse and annihilation in place of time-tuned decays \\cite{Farg-Khlop04}. These particles are generally predicted in Grand Unified Theories (GUT). Their decay lifetime should be comparable with the present age of the Universe, so that their decay products would give rise today to the highest energy part of the cosmic-ray spectrum. Also topological defect decays can give rise to the UHECR primaries \\cite{suggestn}. In any case, the decay dynamics of the super-heavy particles is controlled by the model assumed for particle interactions. In the standard model (SM), the decay after fragmentation finally results into photons, neutrinos and hadrons. In supersymmetry (SUSY), UHE secondary stable neutralinos ($\\chi^0_1$) may be born, too \\cite{ber,susy_oth,toldra,drees}.   The minimal supersymmetric standard model (MSSM) is a well motivated extension of the standard model of elementary particles \\cite{Haber:1984rc}. SUSY stabilizes the Higgs-boson mass under radiative corrections by introducing an entire spectrum of supersymmetric partners for the SM particles. The MSSM  predicts a stable weakly-interacting heavy neutral particle, called neutralino ($\\chi^0_1$), which can be also a good candidate for cold dark matter\\footnote{In general, in the MSSM, there are 4 neutralinos. Here, we are talking about the lightest one $\\chi^0_1$.}.  If the masses of supersymmetric particles are at the TeV scale (as needed to stabilize the Higgs-boson mass), we expect to discover them in a few years at the CERN LHC \\cite{lhc}. In the case  SUSY is the correct theory of particle interactions, it will govern also the decays of super-heavy particles. In \\cite{susy_oth,toldra,drees}, the super-heavy particle decays in the MSSM extension of the SM have been investigated in detail. These decays would involve the entire spectrum of SUSY partners. At the end of the decay chain, one would be left with a spectrum (calculable in the MSSM) of all the  stable particles, that is SM particles plus  stable neutralinos. Hence, a crucial prediction of these models is the presence of neutralinos in the UHECR spectrum produced by super-heavy objects.   \\vskip 0.2cm Another  possible explanation for the GZK paradox considers relic light neutrinos as a calorimeter in resonance with incoming UHE (ZeVs) neutrinos. The latter may come freely from distances above GZK from far BL Lacs (hence explaining the UHECR-BL Lacs correlations) or gamma ray bursts (GRB). They may also arise from topological defects spread all over dark halos (Galactic or extra-galactic). This model, the so-called {\\it Z-burst} model \\cite{zburst}  will be  discussed in the following in view of a new and  analogous ${\\it sneutrino}$-burst model.    \\vskip 0.2cm According to the discussion above, we shall assume in this study a rather hard neutralino spectrum (i.e. $\\sim E_{\\chi}^{-1.5}$) that should be typical for  the GUT-relics fragmentation. In order to estimate the effects of model dependence in the spectral index, we will also compare our results with the ones corresponding to a softer {\\it Fermi-like} flat spectrum, going as $\\sim E_{\\chi}^{-2}$. We shall not consider a much harder {\\it Z-burst}$-$like spectrum  ($\\sim E_{\\chi}^{-1}$) that is usually used in different context, although such a spectrum could be of relevance for a ${\\it sneutrino}$-burst model.   \\vskip 0.2cm One of the major goals for the upcoming UHECR experiments (like Amanda, IceCube, Antares, Nestor) is to verify these predictions. In \\cite{pnath}, the possible signatures of up-going UHE neutralinos in the future satellite-based detector EUSO \\cite{euso} have been investigated. The resonant production of quite heavy squarks ($m \\sim 0.8 - 1.2 ~\\rm TeV$) in the $\\chi_1^0$ scattering off nucleons in the earth's atmosphere, after crossing all the earth volume, has been considered. Heavy squarks are selected so that the neutralino interaction cross sections are moderate, and the neutralino flux is not completely stopped by the earth screen (while the corresponding neutrino background does not survive the passage through the earth). An analogous process leading to a smaller neutrino earth opacity has been also considered earlier,  while discussing the up-going $\\tau$ neutrinos \\cite{Fargion2002a}. Note that, in EUSO, the UHE neutralino interaction in the atmosphere is polluted by more abundant up-going-horizontal $\\tau$ air-showers induced by UHE $\\nu_{\\tau}$ earth skimming in the terrestrial crust \\cite{skim_tau}.  \\vskip 0.2cm In the present article, we will focus on the interaction of the highest energy neutralinos with both electrons and nucleons  in a earth-based detector like IceCube \\cite{icecube}. In particular, we will study the possibility to detect  the resonant production of the SUSY partners of electrons and quarks, when their masses are larger than  the present experimental limits derived from direct searches at high-energy colliders ($m_{\\tilde e}\\gsim 100$ GeV for selectrons \\cite{emass}, $m_{\\tilde q}\\gsim 300$ GeV for squarks\\footnote{Even lighter squarks are allowed for $m_{\\tilde g}\\gg m_{\\tilde q}$.} \\cite{tevatron}). We will also take into account the constraints on the supersymmetry parameter space coming from bounds on the relic density of the lightest neutralino, in case the latter is responsible for most of the cold dark matter in the universe.  Neutralinos with appropriate energy can produce a right/left selectron ($\\tilde e_R$/$\\tilde e_L$) resonance in the detector via the scattering off the detector electrons $$ \\chi_1^0 + e^- \\rightarrow \\tilde e_{R,L}  \\rightarrow X $$ (see  Fig.~\\ref{feyn_diag}). Here, $X$ stands for some visible final state  arising from the $\\tilde e_{R,L}$ decay, that will be eventually revealed by the detector. As far as the scalar mass is not too heavy and $\\chi_1^0$ has a relevant gaugino component, for resonant $\\chi_1^0$ energies this channel gives the dominant contribution to the scattering cross section, and one can neglect $t$-channel effects. \\\\ The kinematics and signature  is somehow similar to the $W^-$ resonant production  in UHE antineutrino-electron scattering \\cite{glashow} $$ \\bar \\nu_e + e^- \\rightarrow W^-  \\rightarrow X \\; , $$ that is expected to occur at the $Glashow$ resonant  energy $E_{\\nu} = M_W^2/(2  m_e) \\simeq 6.3 $ PeV. Since we can predict the exact characteristics of this channel (like its hadronic and electromagnetic showering and its final visible energies), the corresponding  events could be used to  calibrate the  SUSY resonance observation.  Right/left-squark ($\\tilde q_R/\\tilde q_L$) resonances can be produced through the neutralino scattering off the quarks in the detector nucleons $$ \\chi_1^0 + q \\rightarrow \\tilde q_{R,L}  \\rightarrow X \\; . $$ Although, at the partonic level, this process has the same resonant structure as the selectron channel, due to the quark momentum spread inside the nucleon, the cross-section resonance peak is smeared, and the $\\chi_1^0$ incident-energy resonant characteristics are lost.  The background coming from the UHE $\\nu$ interaction with nucleons ($\\nu + q  \\rightarrow X $) by $t$-channel processes could create a major continuous background  overlapping or overshadowing  the squark signal. In the following, after computing the relevant $\\chi_1^0$ event rates, we will address the issue of disentangling the SUSY signal from  the SM resonant and continuous background given by the competitive UHE neutrinos.   \\vskip 0.2cm Higher-energy neutralinos [with $E_{\\chi}\\sim m_{\\tilde \\nu }^2/(2 m_{\\nu})$] could interact with light relic neutrinos in the Local Group or Super-Galactic Plane, and produce sneutrino  $\\tilde \\nu$ resonances in the process \\cite{Fargion:2001pu} $$ \\chi_1^0 + \\nu \\rightarrow \\tilde \\nu  \\rightarrow X \\; . $$ The following sneutrino decay  into final states containing hadrons can in principle mimic a  {\\it Z-burst} event \\cite{zburst}. In this paper, we will also go through  a $\\tilde \\nu$-{\\it burst} scenario.  \\vskip 0.4 cm While calculating the neutralino scattering cross-sections, we will stick to a particular model of SUSY breaking, namely the minimal supergravity mediated model (mSUGRA) \\cite{msugra}. The mSUGRA and its variants are theoretically well motivated. The parameter space of mSUGRA is characterized by five quantities: $m_0$ and $m_{1/2}$, the universal scalar and gaugino masses, respectively; $A_0$, the universal trilinear scalar coupling (all taken at the scale of grand unification); $\\tan \\beta$, the ratio of the vacuum expectation values of the two Higgs doublets; the sign of $\\mu$, the Higgs mixing parameter  (${\\cal W} \\ni \\mu H_1 H_2$, where $H_1, H_2$ are the two higgs superfields in the superpotential $\\cal W$).  While the sign of $\\mu$ is arbitrary, its magnitude is determined by requiring the radiative electroweak symmetry breaking.    In most of the mSUGRA parameter space, the lightest neutralino is almost a pure bino $\\tilde B$, the superpartner of the $U(1)$ gauge field \\footnote{In general, a neutralino is a linear combination of the two Higgs superpartners $\\tilde H_1$, $\\tilde H_2$ (higgsinos), a  $\\tilde W_3$ (wino) and a $\\tilde B$ (bino).}. This is particularly important for $\\chi_1^0$ being a good candidate for cold dark matter \\cite{dark_mat}. Then, the {\\it right} slepton  and  {\\it right} squark  are  maximally coupled to the lightest neutralino, and resonant $\\chi_1^0 f\\to \\tilde f_R$ cross sections are in general enhanced with respect to $\\chi_1^0 f\\to \\tilde f_L\\;$ cross sections (where $f=e,q$).  We will then restrict our study to the right-scalar resonant production and decay into the same initial particles $\\chi_1^0 f$. The corresponding production rates will turn out to be quite straightforward to compute, the only critical parameters being the physical masses of the particle involved in the process (that is the resonance and neutralino masses, $m_{\\tilde f_R}$ and $m_\\chi$). \\\\ In any case, one should keep in mind that the model predicts further new channels associated to the left scalar resonances (sort of \\textit{twin shadows} of the right scalars) whose rate is in general suppressed with respect to the right scalar by both the lower (and in general more model-dependent) relevant decay branching fractions (see Section 3), and the heavier resonance masses (since in general $m_{\\tilde f_L}\\gsim m_{\\tilde f_R}$).    The neutralino energies that are relevant for the resonant reactions $\\chi_1^0 q\\to \\tilde q$, $\\chi_1^0 e\\to \\tilde e$, $\\chi_1^0 \\nu\\to \\tilde \\nu$ are strictly determined by the resonance and target masses [$E_{\\chi}\\sim m_{\\tilde f}^2/(2 m_{target})$], that is \\bea E_{\\chi}&\\gsim& \\frac{m_{\\tilde q}^2}{2 m_p}\\sim 10^5 \\;\\; {\\rm GeV}  , \\label{rang}\\\\ E_{\\chi}&\\sim& \\frac{m_{\\tilde e}^2}{2 m_e}\\simeq 10^7 \\div 10^9  \\;\\; {\\rm GeV} , \\label{range} \\\\ E_{\\chi}&\\sim& \\frac{m_{\\tilde \\nu}^2}{2 m_{\\nu}}\\simeq 10^{14}  \\;\\; {\\rm GeV}, \\eea respectively ($m_p,m_e,m_\\nu$ being the proton, electron, relic-neutrino masses), where we assumed $m_{\\tilde q}> 300$ GeV, $m_{\\tilde e}\\simeq 100-1000$ GeV, $m_{\\tilde \\nu}\\simeq 200$ GeV, and $ m_{\\nu}\\simeq 0.1$ eV.   The plan of this paper is the following. In Section 2, we discuss our assumptions on the  UHE-neutralino spectrum. In Section~3, we present  the cross-sections for the UHE-neutralino interaction with electrons and quarks leading to selectron and squark resonances. In Section~4,  we  present the event rates for the selectron and squark resonant production by UHE $\\chi_1^0$ for a 1-km$^3$ ice detector, in benchmark mSUGRA scenarios that takes into account cosmological bounds on the neutralino relic density. They will then  be compared with the number of events expected to arise from the {\\it Glashow} resonance $\\bar\\nu_e e \\rightarrow W^-$, and from the $\\nu (\\bar \\nu )-$nucleon interaction via charged- and neutral-current processes, assuming comparable UHE-neutrino and neutralino spectra. In  Section~5, we  address the issue of disentangling the SUSY events in the detector by looking at their peculiar energy spectrum and fluence for visible decay products. In  Section~6, we will discuss a possible {\\it $\\tilde \\nu$-burst} scenario, corresponding to the resonant interaction of  UHE $\\chi_1^0$'s with light relic neutrinos in the Local Group or SuperCluster hot dark halos. Finally,  in Section~7, we present our conclusions. ",
        "conclusions": "If elementary-particle interactions  are governed by supersymmetry, then an important component of the UHECR may be given by the lightest stable supersymmetric particle, namely the lightest neutralino.  For instance, some of the models trying to explain the UHECR spectrum beyond the GZK cut-off via super-heavy particle decays, when including  supersymmetry, predict a flux of UHE neutralinos arriving to the earth. One of the major challenges of the future cosmic-ray experiments is to verify such predictions.  In this paper, we computed the event rates for  right selectrons and right squarks produced on resonance, when UHE neutralinos scatter off the electrons and quarks in a detector like IceCube. We assumed a well-motivated model of supersymmetry, that is minimal supergravity, and analyzed the signal for moderate scalar masses, when the resonant scattering is expected to be dominant on the $t$-channel scattering.  The right scalar sector turns out to be particularly straightforward to analyze.  One has in general $B(\\tilde {e}_R \\rightarrow e\\;\\chi^0_1)\\simeq 1$, and $B(\\tilde {q}_R \\rightarrow q\\;\\chi^0_1)\\simeq 1$. As a consequence, after constraining the mSUGRA parameter space by negative collider searches and neutralino relic-density bounds, the relevant phenomenology depends only on the physical masses of the resonant right scalar and lightest neutralino. Any other parameter dependence, such as the one arising from the neutralino physical compositions and couplings,  drops off. We performed a detailed study in scenarios inspired to the post-WMAP mSUGRA benchmarks of \\cite{Battaglia:2003ab}.   For comparison, event rates for the expected resonant $W$ signal  in UHE anti-neutrino scattering off electrons in matter were also presented.   Two different  phenomenological forms for the neutralino fluxes were assumed to calculate the event rates: power laws with exponent either $\\beta_0=1.5$  or $\\beta_0=2$. We set the flux normalization on the basis of present bounds on the UHE neutrino flux from the AGASA and Fly's Eye experiments (assuming, quite conservatively, that interaction rates of neutrino and neutralino with matter are of comparable strength).   For similar neutrino and neutralino fluxes, in the leptonic channel the event rates for SUSY particle resonances are very promising (up to a few tens of events per year in IceCube). They are remarkably greater than the leptonic Glashow $W$ signal for most of the post-WMAP benchmarks characterized by moderate scalar masses ($m_{\\tilde e_R}\\lsim 500$ GeV), that fall  either into the `bulk' or into the co-annihilation region. The corresponding hadronic signal is less structured in the energy spectrum, and penalized by the larger resonance masses ($m_{\\tilde q_R}\\gsim 600$ GeV). A few events per year in Icecube are expected in the most promising benchmarks, in the hadronic case. Signal in the `funnel' and `focus point' benchmarks is depleted by the large $m_0$ down to less than 0.01 event/year.   We showed here for the first time the detailed spectrum of the visible energy released in the $\\bar \\nu_e e^-$ reaction in underground detectors, and discussed how to calibrate the SUSY  signal through the expected $W$ spectrum.   We also considered the possibility of resonant production of sneutrinos in the interaction of UHE neutralinos with light relic neutrinos in the extended Galactic or Local Group  hot neutrino halo. The visible decay products of the sneutrinos might trigger air-showers beyond GZK energies, thus mimicking the so called {\\it Z-burst} process. However, even in the most promising scenarios, we found that the {\\it sneutrino-burst} event rate would be more than one order of magnitude smaller than the {\\it Z-burst} rate (assuming similar fluxes for UHE neutrinos and neutralinos). Furthermore, the UHE neutralino energies tuned for the {\\it sneutrino-burst}  are usually higher than for neutrinos in {\\it Z-burst}, making it less attractive than the standard {\\it Z-burst} model.  In conclusion, we found a well-defined and interesting  window of mSUGRA parameters that can be tested by the next-generation underground detectors. The corresponding signal may exceed in some cases the standard neutrino ones. Its imprint may be disentangled in a realistic way. These high-energy astrophysical traces could offer a new opportunity to  high-energy astrophysics to anticipate  SUSY  discovery at colliders.  \\vskip 1.2cm {\\bf Acknowledgments} \\\\ This work was partially supported by the RTN European Programme MRTN-CT-2004-503369 (Quest for Unification).  \\newcommand{\\plb}[3]{{Phys. Lett.} {\\bf B#1} #2 (#3)}                  % \\newcommand{\\prl}[3]{Phys. Rev. Lett. {\\bf #1} #2 (#3) }        % \\newcommand{\\rmp}[3]{Rev. Mod.  Phys. {\\bf #1} #2 (#3)}             % \\newcommand{\\prep}[3]{Phys. Rep. {\\bf #1} #2 (#3)}                   % \\newcommand{\\rpp}[3]{Rep. Prog. Phys. {\\bf #1} #2 (#3)}             % \\newcommand{\\prd}[3]{Phys. Rev. {\\bf D#1} #2 (#3)}                    % \\newcommand{\\np}[3]{Nucl. Phys. {\\bf B#1} #2 (#3)}                     % \\newcommand{\\npbps}[3]{Nucl. Phys. B (Proc. Suppl.) {\\bf #1} #2 (#3)} % \\newcommand{\\sci}[3]{Science {\\bf #1} #2 (#3)}                 % \\newcommand{\\zp}[3]{Z.~Phys. C{\\bf#1} #2 (#3)} \\newcommand{\\epj}[3]{Eur. Phys. J. {\\bf C#1} #2 (#3)} \\newcommand{\\mpla}[3]{Mod. Phys. Lett. {\\bf A#1} #2 (#3)}             % \\newcommand{\\apj}[3]{ Astrophys. J.\\/ {\\bf #1} #2 (#3)}       % \\newcommand{\\jhep}[2]{{Jour. High Energy Phys.\\/} {\\bf #1} (#2) }% \\newcommand{\\astropp}[3]{Astropart. Phys. {\\bf #1} #2 (#3)}            % \\newcommand{\\ib}[3]{{ ibid.\\/} {\\bf #1} #2 (#3)}                    % \\newcommand{\\nat}[3]{Nature (London) {\\bf #1} #2 (#3)}         % \\newcommand{\\app}[3]{{ Acta Phys. Polon.   B\\/}{\\bf #1} #2 (#3)}% \\newcommand{\\nuovocim}[3]{Nuovo Cim. {\\bf C#1} #2 (#3)}         % \\newcommand{\\yadfiz}[4]{Yad. Fiz. {\\bf #1} #2 (#3);             % Sov. J. Nucl.  Phys. {\\bf #1} #3 (#4)]}               % \\newcommand{\\jetp}[6]{{Zh. Eksp. Teor. Fiz.\\/} {\\bf #1} (#3) #2; {JETP } {\\bf #4} (#6) #5}% \\newcommand{\\jetpl}[6]{{\\em Pisma Zh. Eksp. Teor. Fiz.\\/ } {\\bf #1} #2 (#3) [{\\it JETP Lett.\\/} {\\bf #4} #5 (#6)]}     % \\newcommand{\\philt}[3]{Phil. Trans. Roy. Soc. London A {\\bf #1} #2 (#3)}                                                          % \\newcommand{\\hepph}[1]{hep--ph/#1}           % \\newcommand{\\hepex}[1]{hep--ex/#1}           % \\newcommand{\\astro}[1]{astro--ph/#1}         %"
    },
    "0410/astro-ph0410055_arXiv.txt": {
        "abstract": "{   Observations of the nearby starburst galaxy \\mbox{NGC 253} in the 21-cm line reveal the presence of neutral hydrogen in the halo, up to 12 kpc from the galactic plane. This extra-planar \\HI\\ is found only in one half of the galaxy and is concentrated in a half-ring structure and plumes which are lagging in rotation with respect to the disk. The \\HI\\ plumes are seen bordering the bright H$\\alpha$ and X-ray halo emission.  It is likely that, as proposed earlier for the H$\\alpha$ and the X-rays, also the origin of the extra-planar \\HI\\ is related to the central starburst and to the active star formation in the disk. A minor merger and gas accretion are also discussed as possible explanations.\\\\ The \\HI\\ disk is less extended than the stellar disk.  This may be the result of ionization of its outer parts or, alternatively, of tidal or ram pressure stripping. ",
        "introduction": "\\mbox{NGC 253} is a barred Sc galaxy in the Sculptor group of galaxies, at a distance of 3.9 Mpc \\citep{kar03}. It is one of best nearby examples of a nuclear starburst galaxy. Outside of its very active nuclear region, \\mbox{NGC 253} shows a global morphology in gas and stars expected for a normal spiral galaxy.\\\\ \\begin{figure*}[!ht] \\includegraphics[width=\\textwidth]{channelpanel_compressed.eps} \\caption{NGC 253: \\HI\\ channel maps. The ellipse outlines the bright optical disk (R$_{25}$). Contours are at -4.6, -2.3, 2.3 ($1.75\\sigma$), 4.6, 9.2, 18.4, 36.8, 73.6, 147.2, and 294.4 mJy/beam, with negative contours dashed. The heliocentric radial velocities (km/s) are given in the upper left corners. The 70'' beam is shown by the filled circle in the bottom left panel.} \\label{channels} \\end{figure*} The nuclear starburst is thought to be forming stars at a fairly high rate and to produce a superwind \\citep{hec90} which brings material into the halo. Starburst, disk and halo of \\mbox{NGC 253} have been studied at various wavelengths. The X-ray emission has been observed with ROSAT \\citep{pie00}, with Chandra \\citep{str02} and with XMM-Newton \\citep[][ data on the halo unpublished]{pie01}.  The ROSAT observations have revealed diffuse soft X-ray emission from the nucleus, the disk and a halo extending up to a distance of 9 kpc from the disk. Narrow band H$\\alpha$ imaging \\citep{hoo96} has shown the presence of a diffuse ionized gas (DIG) component. The radio continuum study by \\citet{car92} has revealed a bright disk and a radio halo extending to a height of at least 9 kpc above the plane of \\mbox{NGC 253}. Infrared studies with IRAS \\citep{ric93,alt98} and ISO \\citep{rad01} show that there is also dust in the halo region.\\\\ Studies of a number of nearby spirals with active star formation show considerable amounts of \\HI\\ in their halos \\citep{fra01,fra03}. Also the well-known IVCs and HVCs of our Galaxy \\citep{wak97,woe04} may represent a similar halo component. The origin of such cold halo gas is not known with certainty, but it is thought that galactic fountains set up by the active star forming regions of the disk may be at least partly responsible.  One would therefore expect a significant amount of extra-planar gas also in \\mbox{NGC 253}, because of the starburst and of the star formation which is very active over a large fraction of its disk. But so far no extra-planar \\HI\\ had been found in this galaxy \\citep[cf.][]{puc91,kor95}.\\\\ Here we present the results of a new and deeper \\HI\\ study of \\mbox{NGC 253} which shows that there is indeed \\HI\\ in the halo of this starburst galaxy. Although the central starburst and the active star formation in the disk are likely to cause gas to flow into the halo, there can be other explanations for the origin of this extra-planar \\HI. For example, in many galaxies, including the Milky Way, observations show ongoing accretion. We will discuss the different possibilities for the origin of this extra-planar \\HI.\\\\ \\begin{figure*}[!ht] \\includegraphics[width=\\textwidth]{figure_1_compressed.eps} \\caption{Multi-wavelength image of the nearby spiral galaxy \\mbox{NGC 253}. The deep optical image is from \\citet{mal97}, the blue shows the DSS optical disk, the red the X-ray emission (0.1 - 0.4 keV) from ROSAT \\citep{pie00}, and the green contours our Compact Array \\HI\\ observations. Contours are at 0.18, 0.36, 0.72, 1.4, 2.9, 5.8, 12 and 23$\\times 10^{20}$ cm$^{-2}$. The circular beam has a full half-width of 70$^{\\prime\\prime}$.\\label{fig1}} \\end{figure*} ",
        "conclusions": "The present \\HI\\ observations of \\mbox{NGC 253} have revealed the presence of extra-planar gas in the northeast half of the galaxy. Furthermore, they have shown that the \\HI\\ disk is unperturbed and apparently truncated well inside the stellar disk.\\\\ These results will be discussed here in connection with the nuclear starburst and galactic fountains in the central regions of the disk. Also the possibility of a minor merger and gas infall will be given some consideration.\\\\ \\begin{figure}[!t] \\includegraphics[width=\\columnwidth]{anom_on_opt_compressed.eps} \\caption{Contour map of the anomalous \\HI\\ superimposed on the optical image of \\mbox{NGC 253} from the Digitized Sky Survey. Contours are at 1.8, 3.6, 7.2 and 14$\\times 10^{19}$ cm$^{-2}$.} \\label{fig7} \\end{figure} \\subsection{Structure and kinematics of the extra-planar \\HI}  We have constructed models of the 3-D structure and kinematics of \\mbox{NGC 253} consisting of a thin \\HI\\ disk and a thick disk with different kinematics. The picture we derive for the latter is that of a thick disk with a large, central cavity. As seen in Figs. \\ref{fig3} and \\ref{fig4}, this structure, resembling that of a half-ring, seems only present in the northeast side of the galaxy. The neutral gas forming it is located above and below the outer border of the northeast half of the disk. The ring rotates more slowly than the thin disk. Radial motions, both inward and outward, may also be present. The \\HI\\ plumes seen above and below the disk may be filaments or wall-like extensions of the ring to high z-distances. The cavity in the halo outlined by ring and plumes has a radius of about 6 kpc.\\\\  Two questions arise in connection with the puzzling northeast/southwest halo asymmetry.  One is whether there is no extra-planar gas at all on the southwest part of \\mbox{NGC 253}, either in neutral or ionized form.  The other is what might have caused such asymmetry.\\\\ As to the first question, it should be noted, as already pointed out earlier, that in the northeast part of the galaxy the \\HI\\ emission is barely detected and a factor 2 weaker emission in the southwest would be below detection. Therefore deeper \\HI\\ observations are necessary to confirm the asymmetry and place better upper limits on the \\HI\\ emission in the southwest halo region. At any rate, the present limits already imply an unexpectedly low column density of \\HI.  For comparison, \\HI\\ densities as found in the halo of \\mbox{NGC 891} and of \\mbox{NGC 2403} would have been detected with the present observations of \\mbox{NGC 253}.    There is the possibility that the asymmetry in \\HI\\ is simply caused by complete ionization over most of the halo including all of the southern part. This will be briefly considered below.\\\\ As to the cause of the asymmetry, it is interesting to note that an asymmetry is also seen in the X-rays plumes and in the H$\\alpha$, although not as pronounced as in \\HI, whereas no striking asymmetry is present in the inner galaxy disk either in the optical, in the infrared (see the Two Micron All Sky Survey, Jarrett et al. 2003, and the 60 $\\mu$m dust, Rice 1993), or in the \\HI. In H$\\alpha$ and the radio continuum \\citep{car92} the northern spiral arm is brighter than the southern arm. It is plausible that these asymmetries are related to the asymmetry found for the extra-planar \\HI\\ emission as will be discussed in section 4.2.   \\subsection{Galactic fountains and starburst}  The main source of energy input into the halo is probably the central starburst. A significant amount of the flux comes, however, also from the region of active star formation in the central regions of the disk, within a 4 kpc radius \\citep{ulv00}, which is particularly bright in the radio continuum \\citep[cf.][]{car92}, the H$\\alpha$ and the soft X-rays.  The superbubble expanding into the halo is therefore the combined effect of the nuclear starburst and to a large extent also of the {\\bf star formation} in the inner disk or ring.\\\\ The relative locations of \\HI, radio continuum, X-rays and H$\\alpha$ provide important clues for understanding the 3-dimensional picture. There is a correspondence between \\HI\\ and radio continuum features. The shapes of the halo extensions in the 0.33 GHz image \\citep{car92} look similar to the shape of the \\HI\\ halo distribution. The \\HI\\ plume (to the SE) appears at the same position as the radio continuum spur observed by \\citet{car92}. This spur might mark the eastern edge of the outflow cone of matter.  As already noted, the H$\\alpha$ and the X-ray emission are displaced with respect to the \\HI\\ plumes. Similarly, there also seems to be a displacement between the X-rays and the H$\\alpha$ emission \\citep[see Fig. 10b of][]{str02}.  \\citet{str02} have made several models to explain the X-ray and the H$\\alpha$ observations. The shape of the \\HI\\ distribution as derived here, supports their model 5 (see their Fig. 11d and Fig. \\ref{fig8} in this paper). In this model, the X-ray emission originates from a shock-heated shell around a tenuous hot superbubble, which is surrounded by a cooler shell of \\HII\\ and \\HI\\ that was dragged up from a thick disk. The cartoon in Fig. \\ref{fig8}, partly adapted from \\citet{str02}, illustrates schematically the gas structure of \\mbox{NGC 253}. The thick layer in this picture is formed by the \\HI\\ clouds in the lower halo. The cavity in the halo cloud layer is filled by a hot (about $10^6$ K) halo medium that is observed in the soft X-rays \\citep[see sketch (Fig. 10) in ][]{pie00}.\\\\  \\begin{figure}[!ht] \\includegraphics[width=\\columnwidth]{model_compressed.eps} \\caption{Same as figure 11d of \\citet{str02}, but with an \\HI\\ shell added. The various components are described in the text. The dashed lines on the right outline parts that have not been detected.} \\label{fig8} \\end{figure}  In the picture suggested here, the absence of extra-planar \\HI\\ over the largest part of the halo region of \\mbox{NGC 253} (the \\HI\\ cavity described above), especially when compared with other galaxies with \\HI\\ halos, raises interesting questions. Why is there less halo gas in \\mbox{NGC 253} than in \\mbox{NGC 891} and 2403? and even more striking, if the halos of \\mbox{NGC 891} and 2403 are the result of galactic fountains, why should there be less gas in the halo of \\mbox{NGC 253}, where there is a starburst and a high degree of star formation activity?  Is it possible that there was a neutral gas halo in \\mbox{NGC 253} and that it has been swept away by the expanding superbubble? Or, alternatively, could there be extra-planar gas everywhere, filling the halo of \\mbox{NGC 253} also in the southwest part, and be all ionized?  This should be addressed with deeper H$\\alpha$ studies, although it seems to be contradicted by the presence of fainter H$\\alpha$ emission in the SW part where the \\HI\\ is missing. The extraplanar \\HI\\ in the northeast side might be part of the gas swept away, or, in the case of a fully ionized halo, it could be cold (or cooling down) gas not ionized because of its higher density.\\\\  \\subsection{Minor merger, gas infall}  The observed extraplanar \\HI\\ could be accreted gas from a merger with a satellite. Perhaps, such a minor merger could also explain the triggering of the starburst in an otherwise normal, non-interacting galaxy like NGC 253. The age of the starburst and the superbubbles is a few times 10$^{7}$ yr. It would take approximately 10$^{9}$ yr to spread accreted gas symmetrically over the whole disk. These timescales for the infall and for the starburst would all be consistent with the observed asymmetric \\HI\\ halo distribution. There are also indications from the optical (see section 4.4) that a minor merger may have occurred.\\\\ A minor merger would naturally account for the anomalous \\HI\\ and the asymmetric picture. The question, however, is whether it could produce the remarkable velocity-space continuity and also the symmetry with respect to the major axis of the extra-planar \\HI\\ with the disk of \\mbox{NGC 253} pointed out above and shown in Fig. \\ref{fig5}. The X-rays and H$\\alpha$ would still have to be explained by the superwind from the starburst, while their apparent connection with the \\HI\\ plumes could be due to the superwind hitting the tidal wreckage. All considered, a merger event seems unlikely in this case.   \\subsection{The truncated \\HI\\ disk}  What causes the cutoff at the edge of the \\HI\\ disk? According to \\citet{bla97} there is H$\\alpha$ emission along the major axis of \\mbox{NGC 253} beyond the SW edge of the \\HI\\ disk and its velocity is in agreement with that of the \\HI. They conclude that most likely the hot stars in the disk of \\mbox{NGC 253} itself are responsible for the ionization. For the disk to be ionized in its outer parts by the stars and the starburst, its outer layer has to be warped.  As noted in the previous section, there is evidence in the present data for the presence of such a warp as also found by \\citet{puc91} and Koribalski et al.(1995). The position and inclination angles of the disk change somewhat outside a 9' radius.  Ionization of the outer hydrogen disk by the central starburst seems, therefore, to be possible. Perhaps one should also consider the possibility of strong UV radiation from the halo region. In this connection, the reason why the extra-planar \\HI\\ on the northeast side is not ionized may lie in its relatively high densities.\\\\ A different mechanism for reducing the size of the \\HI\\ disk in \\mbox{NGC 253} is ram pressure stripping. However, the Sculptor group, of which \\mbox{NGC 253} is a member, is not a compact group but a loose filament of galaxies extended along the line of sight \\citep{jer98, kar03}.  It is therefore unlikely to contain a significant amount of intra-group medium capable of causing the stripping.\\\\ A third possibility is that the gaseous disk of \\mbox{NGC 253} has been truncated as the result of a minor merger. A number of observational elements suggest that \\mbox{NGC 253} might have undergone an encounter. The galaxy shows an unusual stellar halo \\citep{mal97} with an optical loop extending to the south at the southwest edge of the galaxy \\citep{bec82} (see also Fig. \\ref{fig1}). A minor merger has already been mentioned above as a possible, but unlikely cause for the anomalous gas."
    },
    "0410/astro-ph0410749_arXiv.txt": {
        "abstract": "We present preliminary iron abundances and $\\alpha$ element (Ca, Mg) abundance ratios for a sample of 22 Red Giant Branch (RGB) Stars  in the Sagittarius galaxy (Sgr), selected near the RGB-Tip. The sample is representative of the Sgr dominant population. The mean iron abundance is [Fe/H]=-0.49. The $\\alpha$ element abundance ratios are slightly subsolar, in agreement with the results recently prensented by \\cite{boni}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410263_arXiv.txt": {
        "abstract": "We present the first results of a campaign to obtain orbital solutions of subdwarf B (sdB) stars from the Edinburgh-Cape survey. We have obtained blue spectra of 35 sdBs, 20 of which have been observed in more than two epochs. 15 out of the 35 are certain binaries with a few other objects showing radial velocity variations with small amplitude, possibly long period sdB binaries. We have secured the orbital parameters for 2 of the 15 systems and narrowed down the orbits of another one to a small range of periods. These preliminary results only use data taken up to December 2003. ",
        "introduction": "In March 2002 we started a campaign to obtain orbital solutions of 100 subdwarf B (sdB) binaries from the Edinburgh-Cape (EC) survey (Stobie et al. 1997). This work was intended as the Southern Hemisphere counterpart of that undertaken by Maxted et al. (2001) and Morales-Rueda et al. (2003). In their work Morales-Rueda et al. (2003) select their targets from two main sources, the Palomar-Green (PG) and the Kitt Peak-Downes (KPD) catalogues. The PG catalogue is biased against targets that show a Ca II H-line (as these were taken out of the catalogue), and thus, against sdB binaries with main sequence companions. On the other hand the EC survey does not show these biases, as the object selection is based on colours only, and is ideal to obtain an ``unbiased'' sample of sdB binaries that can then be compared with binary evolution theories (Han et al. 2003). We should mention that some bias might still remain as in the cases where the main sequence star was somewhat brighter than the sdB star (i.e. combined colour redder than $U-B$ = $-$0.4 ), the binary would have been classified probably as a late A star and would have been taken out of the catalogue. ",
        "conclusions": "Their orbital periods place EC00404$-$4429, EC02200$-$2338 and EC12327$-$1338 in the group of sdB binaries formed via the common envelope (CE) ejection channel (see fig. 10 of Han et al. 2003). The minimum mass of the companion, calculated assuming the canonical mass of 0.5\\,M$_{\\odot}$ for the sdB, indicates that the companions are probably white dwarfs, thus these binaries have probably formed through the second CE path (Han et al. 2003).  We have added the orbital periods of these three new systems to the orbital period distribution of sdBs known, Fig. 4. The most interesting features of this distribution are the excess of binaries at orbital periods around 1 d and that the very long period systems (tens and hundreds of days) are starting to appear in the sample. The comparison of the observed distribution with theoretical predictions (Han et al. 2003) is invaluable to determine, among other important evolutionary unknowns, plausible values for the CE ejection efficiency.  \\begin{figure}[!ht] \\plotfiddle{hist2004.ps}{2.5cm}{-90}{40}{40}{-150}{120} \\vspace{35mm} \\caption{Orbital period distribution for the sdB binaries known. The figure is an update of Morales-Rueda et al. (2003) and Morales-Rueda et al. (2004).} \\end{figure}"
    },
    "0410/astro-ph0410439_arXiv.txt": {
        "abstract": "SCUBA-2, which replaces SCUBA (the Submillimeter Common User Bolometer Array) on the James Clerk Maxwell Telescope (JCMT) in 2006, is a large-format bolometer array for submillimeter astronomy. Unlike previous detectors which have used discrete bolometers, SCUBA-2 has two dc-coupled, monolithic, filled arrays with a total of $\\sim$10,000 bolometers. It will offer simultaneous imaging of a 50 sq-arcmin field of view at wavelengths of 850 and $450\\rm\\ \\mu m$. SCUBA-2 is expected to have a huge impact on the study of galaxy formation and evolution in the early Universe as well as star and planet formation in our own Galaxy. Mapping the sky to the same S/N up to 1000 times faster than SCUBA, it will also act as a pathfinder for the new submillimeter interferometers such as ALMA. SCUBA-2's absorber-coupled pixels use superconducting transition edge sensors operating at 120~mK for performance limited by the sky background photon noise. The monolithic silicon detector arrays are deep-etched by the Bosch process to isolate the pixels on silicon nitride membranes. Electrical connections are made through indium bump bonds to a SQUID time-domain multiplexer (MUX). We give an overview of the SCUBA-2 system and an update on its status, and describe some of the technological innovations that make this unique instrument possible. ",
        "introduction": "SCUBA (the Submillimetre Common User Bolometer Array) on the James Clerk Maxwell Telescope has been one of the most successful instruments ever built for a ground-based telescope\\cite{Holl99}.  Examples of scientific highlights from SCUBA are the discovery of an unknown population of dusty star-forming galaxies\\cite{Smai97} and detailed observations of dusty debris disks around nearby stars where planets are forming\\cite{Holl03}.  SCUBA-2 will leverage new technologies for a revolutionary improvement in sensitivity and mapping speed.  The science case for SCUBA-2 is discussed in detail elsewhere\\cite{SPIE02WSH}.  Less than 1~square degree of sky has been mapped in the sub-millimeter to any great depth (e.g. close the the extragalactic confusion limit). SCUBA-2 will be primarily a survey instrument that will map large areas of sky.  The new sub-millimeter interferometers like ALMA and the SMA are limited as survey instruments by their narrow fields of view.  SCUBA-2 will act as a pathfinder for these instruments, filling the JCMT's $8\\times8$-arcmin field of view with dc-coupled pixels. SCUBA-2's field of view is 12 times that of SCUBA.  Also, SCUBA is undersampled by a factor of four, taking 64 pointings (a total of 128 s) for SCUBA to produce a fully-sampled image at 850 and $450\\rm \\mu m$ simultaneously.  SCUBA-2 will produce a fully-sampled image in just one pointing at $850\\rm\\ \\mu m$ or four pointings at $450\\rm\\ \\mu m$.  This means that SCUBA-2 will be able to map large areas of sky about 1000 times faster than SCUBA (see Table~\\ref{table:1}).  With fully-sampled image planes and no sky chopping, SCUBA-2 will have better image fidelity and map dynamic range than SCUBA.  Imaging at two colours simultaneously, SCUBA-2 will allow us to study the physical properties of dust emission and to exploit the better angular resolution at $450\\rm\\ \\mu m$ when weather permits. SCUBA-2 will also carry out deep imaging to the confusion limit in only a few hours.  \\begin{table*}[bht] \\caption{Estimated performance of SCUBA-2.} \\label{table:1} \\renewcommand{\\tabcolsep}{1pc} % \\renewcommand{\\arraystretch}{1.2} % \\newcommand{\\m}{\\hphantom{$-$}} \\begin{tabular}{@{}lcccc} \\hline Parameter           &\\multispan2{\\hfill$450\\rm\\ \\mu m$\\hfill} &\\multispan2{\\hfill$850\\rm\\ \\mu m$\\hfill}\\\\ &SCUBA&{\\bf SCUBA-2}&SCUBA&{\\bf SCUBA-2}\\\\ \\hline Per-pixel NEFD ($\\rm mJy/\\sqrt{Hz}$)   & 400& {\\bf 80} & 90& {\\bf 25} \\\\ Point source NEFD ($\\rm mJy/\\sqrt{Hz}$)  & 400& {\\bf 113} & 90& {\\bf 21} \\\\ Point source extraction speed relative to SCUBA   & 1& {\\bf 13} & 1& {\\bf 18} \\\\ Large area mapping speed relative to SCUBA   & 1& {\\bf 850} & 1& {\\bf 1250} \\\\ \\hline \\end{tabular}\\\\[2pt] \\end{table*} ",
        "conclusions": ""
    },
    "0410/astro-ph0410113_arXiv.txt": {
        "abstract": "We consider the Abelian model of a complex scalar field coupled to a gauge field within the framework of General Relativity and search for cosmological solutions. For this purpose we assume a homogeneous, isotropic and uncharged Universe and a homogeneous scalar field. This model may be inserted in several contexts in which the scalar field might act as inflaton or quintessence, whereas the gauge field might play the role of radiation or dark matter, for instance. Particularly, we propose two such models: (i) in the first, the inflaton field decays to massive vector bosons that we regard as dark-matter; (ii) in the second, due to its coupling to radiation the scalar field is displaced from its ground state and drives an accelerated expansion of the Universe, playing the role of quintessence. We observe that the equations are quite simplified and easier to be solved if we assume a roughly monochromatic radiation spectrum. ",
        "introduction": "The study of scalar fields in cosmological models has been proven to be very fruitful both in inflationary theories and in quintessence models.~\\cite{inflation,Kremer,Linde,KT,revq,quint}  It turns out that the success of these models also raises the need to investigate the interaction of these scalar fields with other forms of matter.  The inflationary theories are based upon the dynamical evolution of a scalar field whose coupling to other fields is normally ignored until reheating process takes place~\\cite{inflation,Linde,KT,reheating}. In the first reheating models a phenomenological friction term $\\Gamma\\dot\\phi$ was inserted into the equation of motion of the scalar field in order to simulate the energy transfer to other forms of matter~\\cite{gammaphiponto}. The theories of reheating have been improved since then, and today they are very more sophisticated, taking under consideration quantum effects and many kinds of possible decays of the inflaton to other particles~\\cite{ressonance,KLS}. Anyway, how reheating really occurred is still an object of intense investigation~\\cite{Kof}.  The very recent measurements of type-IA supernovae~\\cite{SNIa} and the analysis of the power spectrum of the CMBR~\\cite{CMBR} provided strong evidence for a present accelerated expansion of the Universe~\\cite{Carroll,accel}. That offered a new setting for the scalar field cosmological models. In these contexts they are named quintessence fields~\\cite{revq}, and recently many interesting phenomenological models of interacting quintessence have been proposed~\\cite{interacq}.  In the present work a very simple cosmological model is investigated in which a complex scalar field interacts with a vector field. We believe that the kind of interaction treated in this model is very special because it raises naturally by requiring a local $U(1)$ gauge symmetry. This is the principle beyond the electromagnetic interaction and it inspired the development of the other fundamental interactions. According to the current view all the forms of interactions are completely determined by certain internal symmetry groups~\\cite{qft}.  Our starting point will be the following Lagrangian: \\be {\\cal L}_{\\rm mat}=D_\\mu\\phi(D^\\mu\\phi)^*-V(\\phi\\phi^*) -{1\\over 4}F_{\\mu\\nu}F^{\\mu\\nu}, \\ee{1} where \\be D_\\mu\\equiv\\nabla_\\mu+iqA_\\mu,\\qquad F_{\\mu\\nu}\\equiv\\partial_\\mu A_\\nu- \\partial_\\nu A_\\mu. \\ee{1a} The Lagrangian (\\ref{1}) is clearly invariant under local $U(1)$ gauge transformations: $\\phi\\rightarrow \\exp[{iq\\alpha(x)}]\\phi$ and $A_\\mu\\rightarrow A_\\mu -\\partial_\\mu\\alpha(x)$, and associated with this symmetry is the conservation of the four-current \\be {\\cal J}^\\mu=iq[\\phi (D^\\mu\\phi)^*-\\phi^* (D^\\mu\\phi)]. \\ee{1b}  Our basic procedure will be to introduce the Lagrangian (\\ref{1}) in the action with the Einstein-Hilbert Lagrangian and derive the Einstein's equations for that system (Sec.2). Then we shall investigate solutions to those equations that may be applicable in cosmological models. In Sec.3 we ascribe an interpretation to the contribution of each field for the energy density. We also discuss about the flow of energy among the fields. The Fourier expansion of the vector field and spatial average procedures are performed in Sec.4 as a first step towards the application to concrete models. Considerations about the possibilities of such applications comes in Sec.5, and the equations are solved for the simplified case of a roughly monochromatic radiation spectrum. Finally, we conclude in Sec.6 with brief remarks and future prospects. All the fields are treated classically. We choose units such that $c=\\hbar=k=1$ and adopt the convention that Greek indices run from $0$ to $3$ whereas Latin indices run from $1$ to $3$. ",
        "conclusions": "To sum up, we remark that the qualitative treatment performed here point out to interesting mechanisms of accounting for the dark-matter in the Universe (or at least a part of it), and also provides a possible candidate for quintessence. However, the quintessence behavior displayed here differs from the usual quintessence models in the sense that the scalar field is not tending asymptotically to the minimum of its potential. Rather, its minimum is not a stable position and the scalar field is removed from its ground state due to its coupling to a background field of vector bosons.\\footnote{A similar mechanism was employed to produce a short secondary stage of inflation. See~\\cite{felder} and references therein.}  All the same, our study is not complete and a quantitative confrontation of the predictions yield by these models to observations is needful. Besides taking into account the full spectrum, it is also necessary to address the issue of whether the models treated here satisfy the constraints from the anisotropy of the CMBR, the large-scale galaxy distribution and the SN-Ia data. That will be performed in a future work.  \\ack This work was supported by CNPq (Brazil)."
    },
    "0410/astro-ph0410325_arXiv.txt": {
        "abstract": "Synthetic integrated spectral properties of the old Galactic open clusters are studies in this work, where twenty-seven Galactic open clusters of ages $\\geq$ 1Gyr are selected as the working sample. Based on the photometric observations of these open clusters, synthetic integrated spectrum has been made for the stellar population of each cluster. The effects of blue straggler stars (BSSs) on the conventional simple stellar population (SSP) model are analyzed on an individual cluster base. It is shown that the BSSs, whose holding positions in the color-magnitude diagrams (CMDs) cannot be predicted by the current single-star evolution theory, present significant modifications to the integrated properties of theoretical SSP model. The synthesized integrated spectral energy distributions (ISEDs) of our sample clusters are dramatically different from the SSPs based on isochrone only. The BSSs corrected ISEDs of stellar populations show systematic enhancements towards shorter wavelength in the spectra. When measured with wide-band colors in unresolvable conditions, the age of a stellar population can be seriously under-estimated by the conventional SSP model. Therefore, considering the common existence of BSS component in real stellar populations, a considerable amount of alternations on the conventional ISEDs should be expected when applying the technique of evolutionary population synthesis (EPS) to more complicated stellar systems. ",
        "introduction": "By many years of substantial practice, evolutionary population synthesis (EPS) has been proved to be a powerful tool for the study of the evolution of stellar systems. EPS method is based on the great success of the understanding of stars, including stellar structure and evolution (without taking into account rotation, magnetic field and binarity), and the properties of stellar radiation from their atmospheres. As the basic building blocks of synthetic spectra of galaxies, simple stellar populations (SSPs) are built with the integration of the library of isochrones based on the theoretical evolutionary tracks, the library of stellar spectrum of individual stars, and the stellar initial mass function (IMF). The first two ingredients are given at a certain chemical composition, while the IMF is usually considered as universal (see Bressan et al. 1994, for a review).  As the solid base for studying complex stellar systems, EPS method is very successful in many aspects. In analyzing the stellar populations of composite systems such as galaxies, EPS reproduces the overall spectral properties of the stellar components, and also derives star formation and chemical evolution history (Schulz et al. 2002). In the case of simple stellar systems such as star clusters, relying on almost perfect isochrone-fitting to the photometric observations of the systems, the conventional SSP model can give fairly good interpretations upon most of the evolutionary features of star clusters, i.e., the main sequence (MS), red giant branch (RGB), horizontal branch (HB) or red clump giants on the integrated light of these components. However, exceptions do exist. It is often observed that a certain number of definite member stars have holding positions in the color-magnitude diagrams (CMDs) that cannot be predicted by the single-star evolution theory, and consequently the integrated contributions of these stragglers to the total radiation are ignored by the conventional SSP model. Most of the unresolved objects are related to binary scenario and stellar collisional events, such as blue straggler stars (BSSs) in stellar systems. In the case of open clusters, being the most luminous blue objects, BSSs can raise a big challenge to the SSP model (Deng et al. 1999; Schiavon et al. 2004).  Since the first identification by Sandage (1953) in the cluster M3, BSSs have been observed in almost all star clusters, dwarf galaxies and presumably in all other stellar systems (Stryker 1993). These enigmatic stars locate bluer and brighter of the cluster turnoff region in a CMD, and appear $\\it{straggling}$ away from the regular evolutionary path, hence are named 'blue stragglers'. As for the nature of these objects, observations suggest that BSSs are more massive than the regular stars in the turnoff region, thus they are most likely the remnants of stellar merger events, or still binary systems with luminous blue components.  The primary ways of the formation of BSSs are dynamical evolution of close binaries, and collisions in the high density area in the clusters. (Ferraro et al. 1997; Piotto et al. 1999; Ferraro et al. 2003). As suggested by mounting evidence, if BSSs are partially the result of mass transfer or coalescence during resonant interactions involving binary systems (Leonard 1989; Bacon et al. 1996), they would be the visible tracers of binary populations, and additionally they provide an opportunity for learning how interactions in binary systems affect stellar evolution (Pols \\& Marinus 1994). Meanwhile, appreciable frequency of single-single star collisions generally occur only in the cores of the densest clusters (Hills \\& Day 1976), where BSSs can provide observational constraints on stellar collision cross sections for tidal captures, direct collisions, and the corresponding productions in dense stellar environment (Bailyn 1995; Lombardi et al. 2002). Regardless of the origins of BSSs, when mergers occur, many of the cluster properties are affected. Remarkable central concentration of BSSs provides an effective examination on cluster dynamics, such as two-body relaxation in binary systems, core collapse rate, mass segregation and depletion of low mass stars in stellar systems, and energy equipartition among the cluster member stars. Moreover, BSSs do exist in significant numbers in both young and old populations.  They provide a remarkable hotter component with respect to the rest of the cluster (Manteiga et al. 1989) and bring on possible bluer integrated colors and integrated spectral energy distributions (ISEDs) (Deng et al. 1999). Therefore, the presence of numerous BSSs has to be accounted in population synthesis definitely for star clusters, and possibly for dwarf ellipsoidals, ellipticals and other galaxies. It is important to gather empirical information on BSSs in order to characterize correctly the overall properties of populations in stellar systems according to integrated light studies.  Based on the assumption that all members are born simultaneously therefore having the same age and metallicity, open and globular clusters are considered as the best SSP templates in the real world (Bica \\& Alloin 1986; Battinelli et al. 1994). In this paper, our attention is put on the BSSs in Galactic open clusters. The reason is that, for the chemical composition range covered by Galactic population I clusters, BSSs are the most luminous blue objects whose contributions to the integrated light are more prominent, and this is especially outstanding for old open clusters when the red clump instead of blue HB is populated. The BSS effects on the ISEDs become far less important for the Galactic globular clusters, where the most luminous blue objects are the HB stars.  Since the direct observation of cluster integrated light (Bica \\& Alloin 1986) can not attack the problems of field contamination and gravitational evaporation of low-mass stars, it is not possible to get the proper ISED of a cluster in that way. Pure theoretical efforts based on the current theory of stellar evolution also fail at this point simply due to the stragglers. Keeping these in mind, we propose a semi-empirical approach of building cluster ISED in our previous work (Deng et al 1999): after careful membership analysis, the general stellar ingredients well fitted by a theoretical isochrone is represented by the conventional SSP model, in which the 'missing' low mass stars can be recovered with proper initial mass function, so that the true meaning of SSP is conserved; and all the stragglers are included into the ISED by derivation of the spectra from observations.  The importance of BSSs on the studies of stellar formation and evolution, and population analysis of galaxies have been stressed by numerous work (see Stryker 1993 for a review). What we are going to show in this paper, by using old Galactic open clusters, is that the effects of BSSs on the integrated light of a stellar system and on the theoretical model of SSP. With a sample of twenty-seven Galactic open clusters of ages $\\geq$ 1 Gyr, we present in this work a set of integrated spectra and modified integrated B-V color of the clusters. Comparison between the conventional approach and our results is made aimed at investigating BSSs contribution to the conventional SSP model. A short report of our working sample is given in the next section. In Section 3, population synthesis is performed for each cluster. BSSs are treated as hydrogen-burning MS stars, whose spectra are selected from Lejeune (et al. 1997,1998) stellar spectral library, given the physical quantities by fitting evolutionary tracks to their observed positions in the CMDs. The other stellar ingredients in a cluster are approximated with Padova isochrone (Bertelli et al. 1994). The final ISED of a cluster is the combination of these two components. BSSs contribution to the conventional ISEDs varies for different clusters depending on the basic quantities of the cluster and the BSS component. Discussions addressing these issues are given in Section 4. In Section 5, BSSs modification to the integrated properties of SSP model in terms of wide-band B-V color is presented, and the corresponding consequences on age determination of a SSP is also considered. The concluding remarks on our work are presented in Section 6. ",
        "conclusions": "Aimed at detecting BSSs effects on the integrated properties of conventional SSP model, twenty-seven Galactic open clusters older than 1Gyr and their BSS contents are studied in this work. In general, the ISEDs of our sample clusters have been dramatically modified, and such modification is emphasized in UV and blue bands; In terms of photometry, the integrated B-V color of the sample clusters becomes bluer, this leads to an under-estimation either in age or in metallicity when fitting the observed colors of a target unresolved population with that of a conventional SSP model. The present results are summarized and discussed in the following,  \\begin{enumerate}  \\item In this work, twenty-seven Galactic open clusters older than 1Gyr are selected as our working sample. The photometric data of BSSs including B-V colors and V-band magnitudes, N$_{BS}$ and N$_2$ number are directly quoted from AL95. The basic parameters of sample clusters from recent literatures are used when they are available.  \\item We scheme a semi-empirical approach of building cluster ISED. The general stellar ingredients including all member stars well fitted by a theoretical isochrone are represented by the conventional SSP model. BSSs are treated as hydrogen-burning MS stars. Their positions in the CMD are fitted with evolutionary tracks of masses higher than the turnoff. Then, a theoretical stellar spectrum is assigned to each BSS, according to parameters ($\\it{T_{eff}}$ and $\\it{log}$ $\\it{g}$) derived from the best fitting. Finally, the combination of these two components constructs the modified ISED of the cluster.  \\item We have divided the BSSs modifications to conventional ISEDs into three cases, namely, normal star dominate (Fig. 4a), BSSs important (Fig. 4b) and BSSs dominate (Fig. 4c). The modifications are mainly related to the BSSs richness in the clusters and their positions in the observed CMDs. For the first case, BSSs generally locate near the cluster turnoff in the CMD, and the number of BSSs is small when compared with turnoff region stars, therefore they can not produce obvious contribution to the ISED of the cluster. For the second case, BSSs effects get enhanced generally due to: (a). there is a large number of bright BSSs in the cluster; (b). The N$_{BS}$ is small, but there is a few (even just 1!) BSSs which are much brighter than the cluster turnoff point; and (c). the cluster is very member-poor. The third case is an excessive situation. There is at least one extremely luminous BSS in the cluster and/or the cluster is extremely poorly populated. The bright star(s) is so bright that controls almost all the integrated light of the cluster.  \\item We have depicted the alternation of B-V color of our sample clusters caused by BSSs. As shown in Figure 16, the B-V color becomes significantly bluer.  \\item We have discussed the effects of metallicity on BSSs contribution. As shown in Figure 18, they are well correlated in the metal-poor case (Z $<$ 0.02). In Figure 19, the ratio of the alternation of B-V color increases with metallicity increasing. But the correlation does not follow the case of Z $>$ 0.02.  \\item We have scaled the age underestimation of a cluster by fitting the observed B-V color with conventional SSP model. Based on the assumption that all the sample clusters have solar metallicity (Z=0.02), we have shown in Figure 17 that the age of an observed stellar population is seriously under-estimated by conventional SSP model. As clearly shown in Figure 17, 30\\% uncertainties in ages can be given for the least affected clusters, and for many cases, such corrections amounts to twice of the conventional age.  \\item We have also detected the age and metallicity uncertainties when fitting the observed cluster ISED with conventional SSP model, since decreasing either age or metallicity will also make the ISED bluer as inclusion of BSSs does. Taking old open cluster NGC 188 and NGC 2682 as examples, the results of the best fitting show that conventional SSP model can fit perfectly most of the features of the composite ISEDs including BSSs contribution if we let the age or metallicity free, and both age and metallicity can change by a factor of 2.  \\item The modification due to BSSs to the theoretical SSP model corresponding to each cluster show that BSS population is very important, and can be crucial for stellar population analysis of complex stellar system. This is especially true when applying the current EPS technique to unresolved stellar systems. Considering open clusters as typical descendants of regular star forming activities, and regulating our results to the age and metallicity ranges covered by present data set, SSP model derived from single star evolution theory is seriously altered. Therefore, the contribution of such a component should be considered when applying EPS method to more complicated stellar systems.  \\end{enumerate}  The present results rely on the photometry from different sources, and the formation mechanisms and the specific properties of BSSs in each cluster vary greatly among the sample, therefore we are limited to draw our conclusion on a single cluster base. With careful cluster physical parameter measurements and membership analysis of individual BSSs, the current results are still meaningful and can be regarded as reference for further investigation on this problem. To reveal the effects of BSSs on stellar populations in all different cases of age, metallicity and environment, systematic observations together with simulations of stellar dynamical processes at all conditions are needed."
    },
    "0410/astro-ph0410108.txt": {
        "abstract": "{\\large} We present high quality long-slit spectra for three nearby powerful radio galaxies \u00ad-- 3C293, 3C305, PKS1345+12. These were taken with the aim of characterising the young stellar populations (YSP), and thereby investigating the evolution of the host galaxies, as well as the events that triggered the activity. Isochrone spectral synthesis modelling of the wide wavelength coverage spectra of nuclear and off-nuclear continuum-emitting regions have been used to estimate the ages, masses and luminosities of the YSP component, taking full account of reddening effects and potential contamination by activity-related components. We find that the YSP make a substantial contribution to the continuum flux in the off-nuclear regions on a radial scale of 1 \u00ad-- 20~kpc in all three objects. Moreover, in two objects we find evidence for reddened post-starburst stellar populations in the near-nuclear regions of the host galaxies. The YSP are relatively old (0.1 \u00ad-- 2~Gyr), massive ($10^9 < M_{YSP} < 2\\times10^{10}$ M$_\\odot$) and make up a large proportion ($\\sim$1 \u00ad-- 50\\%) of the total stellar mass in the regions of galaxies sampled by the observations. Overall, these results are consistent with the idea that AGN activity in some radio galaxies is triggered by major gas-rich mergers. Therefore, these radio galaxies form part of the subset of early-type galaxies that is evolving most rapidly in the local universe. Intriguingly, the results also suggest that the radio jets are triggered relatively late in the merger sequence, and that there is an evolutionary link between radio galaxies and luminous/ultra-luminous infrared galaxies. ",
        "introduction": "Powerful radio sources are invariably associated with early-type host galaxies and their luminous emission lines and powerful radio emission make them stand out at high redshifts. Therefore, they are potentially unique probes of the evolution of massive early-type galaxies in the early universe. However, if we are to use them in this way, it is crucial to understand the link between the quasar/jet activity and evolutionary processes in the host galaxies. In this context it is notable that the evolution of the co-moving number density of powerful radio sources (Dunlop \\& Peacock 1990), shows a marked similarity to that of the evolution of the global star formation rate in the universe (Madau et al. 1996). This suggests that the triggering of the quasar/jet activity is intimately linked to the evolution of the general galaxy population.  Observations of powerful radio galaxies in the local universe provide important clues to the triggering mechanism for the jet and quasar activity. In particular, the imaging studies of Heckman et al. (1986) and Smith, Heckman \\& Illingworth (1989) provide evidence for morphological features such as double nuclei, arcs, tails and bridges which suggest that, in a substantial subset of radio galaxies, the AGN activity  is triggered by the accretion of gas during galaxy mergers and interactions. This is backed up by spectroscopic studies (Tadhunter et al. 1989, Baum et al. 1990) which show emission line kinematics consistent with an accretion origin for the warm gas in many radio galaxies, although a cooling flow origin is a viable alternative to mergers for some objects at the centres of galaxy clusters.  An attraction of the merger hypothesis for the triggering of the activity is that, potentially, it can help to explain the redshift evolution of the radio source population and its similarity to that of the star formation history of the universe. This is because, in a hierarchical galaxy evolution  scenario, galaxy mergers and interaction are not only a major driver for star formation, but they can also provide the fuel for AGN and jet activity (e.g. Kauffmann \\& Haehnelt 2000). Thus, radio galaxies may represent a particular, post-merger phase, in the evolution of giant elliptical galaxies.  Despite the promise of the merger hypothesis for the triggering of the activity, the following uncertainties remain: \\begin{itemize} \\item What types of mergers or galaxy interactions trigger radio jet activity? \\item At what stage during the merger are the jets triggered? \\item Do all massive early-type galaxies go through one or more radio galaxy phase(s) as they build up via galaxy mergers? \\item What is the relationship between powerful radio galaxies and other types of merging systems in which a massive early-type galaxy is the end point, for example, the ultraluminous infrared galaxies (ULIGS)? \\end{itemize} Unfortunately the quantitative information that can be extracted from the existing morphological and kinematical data is limited. Therefore it is important to develop alternative techniques for investigating the triggering events. One promising alternative is to use the properties of the starbursts that are expected to be triggered as part of the merger process.  Although a substantial proportion of radio galaxies show blue or UV excesses compared with normal early-type galaxies (Lilly \\& Longair 1984, Smith \\& Heckman 1989), various activity-related  components can contribute to the optical/UV continuum and thereby complicate the interpretation of the continuum spectral energy distributions (SEDs) in terms of the underlying stellar populations. Activity-related continuum components include: scattered AGN light (e.g. Tadhunter et al. 1992, Cimatti et al. 1993, Cohen et al. 1999), nebular continuum from the emission line nebulae (Dickson et al. 1995) and direct light from moderately extinguished quasar nuclei (Shaw et al. 1995). However, with careful modelling of spectroscopic and polarimetric data it is possible to accurately quantify the contributions of these components to the optical/UV continua and thereby study the underlying stellar populations of the host galaxies. In this way it has been possible to show that young stellar populations (YSP) make a significant contribution to the optical/UV continua  in 25 \u00ad-- 40\\% of powerful radio galaxies at low and intermediate redshifts ($z < 0.7$: Aretxaga et al. 2002, Tadhunter et al. 2002, Wills et al. 2002, 2004).  A further complicating factor is that, even if young stellar populations can be identified, they are not necessarily related to the normal evolution of the host galaxies. It has been proposed that the shocks induced by the radio jets as they sweep through the ISM of the host galaxies may trigger star formation, and be a major cause of the alignment effect in high redshift radio galaxies (e.g. Rees 1989). However, apart from a few well-publicised examples in which there is good evidence that the excess UV emission along the radio axis is due YSP rather than scattered light or nebular continuum (e.g. Minkowski's object: van Breugel et al. 1985), direct observational evidence for jet-induced star formation is lacking in the overwhelming majority of radio galaxies.    In order to realise the full potential of radio galaxies as evolutionary probes, it is now important to establish the detailed properties of their YSP. In this, the first in a series of papers in which we will investigate the nature of the YSP in radio galaxies, we present the results of the detailed analysis of the continuum spectra of three nearby powerful radio galaxies --- 3C293, 3C305, PKS1345+12.  We assume the cosmological parameters $H_0 = 75$ km s$^{-1}$ Mpc$^{-1}$ and $q_0 = 0.0$ throughout this paper. For these parameters 1.00 arcsecond corresponds to 0.82, 0.75 and 1.98~kpc at the redshifts of 3C293, 3C305 and PKS1345+12 respectively. ",
        "conclusions": "The results presented in this paper demonstrate the potential of studies of the young stellar populations  for investigating the nature of the events that trigger the activity in radio galaxies, the order of events, the timescales, and the relationship between radio galaxies and other types of active objects such as ULIGs. The main results can be summarised as follows. \\begin{itemize} \\item{\\bf Nature of the YSP and the mergers.} The YSP are identified with post-starburst populations that are relatively  old and massive, and represent a large fraction of the total stellar mass in the regions sampled by our spectra. These results imply that we are observing the aftermath of {\\it major mergers} between massive galaxies, at least one of which must have been a gas-rich disk galaxy.  \\item{\\bf Spatial distribution.} Although the YSP are concentrated close to the nuclei of the host galaxies in at least two of our sample objects, in all cases the YSP are spatially extended on a scale of 1 \u00ad-- 20~kpc.  \\item{\\bf Triggering and timescales.} The fact that the post-starburst populations detected in our sources are relatively old compared with the typical ages of extragalactic radio sources, suggests that the activity associated with the observed high surface brightness radio sources was triggered late in the merger sequence --- up to 0.5 -- 2~Gyr after the starburst responsible for the YSP.  \\item{\\bf Luminous and ultraluminous infrared galaxies.} Our results point to a strong link between the radio galaxies and at least a subset of the LIG and ULIG populations; they provide evidence that some LIGs and ULIGs may {\\it evolve into} radio galaxies. \\end{itemize}  It is now important to extend such studies to larger samples of radio galaxies, in order to put these results on a firmer footing, and investigate whether they hold for the general population of radio galaxies with detected YSP.   \\subsection*"
    },
    "0410/astro-ph0410107_arXiv.txt": {
        "abstract": "The direct observation of extrasolar planets and their spectra is coming into reach with the new generation of ground-based near-IR interferometers, like the Very Large Telescope Interferometer (VLTI). The high contrast between star and planet requires an excellent calibration of atmospheric distortions. Proposed techniques are the observation of color-differential or closure phases. The differential phase, however, is only in a first order approximation independent of atmospheric influences because of dispersion effects. This might prevent differential phase observations of extrasolar planets. The closure phase, on the other hand, is immune to atmospheric phase errors and is therefore a promising alternative. We have modeled the response of the closure phase instrument AMBER at the VLTI to realistic models of known extrasolar planetary systems taking into account their theoretical spectra as well as the geometry of the VLTI. We present a strategy to determine the geometry of the planetary system and the spectrum of the extrasolar planet from closure phase observations in a deterministic way without any a priori assumptions. We show that the nulls in the closure phase do only depend on the system geometry but not on the planetary or stellar spectra. Therefore, the geometry of the system can be determined by measuring the nulls in the closure phase and braking the remaining ambiguity due to the unknown system orientation by means of observations at different hour angles. Based on the known geometry, the planet spectrum can be directly synthesized from the closure phases. ",
        "introduction": "\\label{sec:intro}  Since the first extrasolar planet candidate was discovered in 1995 orbiting the solar-like star 51\\,Peg (Mayor \\& Queloz 1995), more than 130 extrasolar planets have been detected by radial velocity surveys (e.g. Mayor et al. 2003 for recent discoveries). So far, the only observational information on the spectrum of an extrasolar planet has been obtained for the transiting system HD209458 (e.g. Charbonnneau et al. 2002, Vidal-Madjar et al. 2004).  \\begin{figure}[h] \\begin{center} \\includegraphics[width=9cm]{vjoergens_f1.eps} \\end{center} \\caption[example] { \\label{fig:clphase} {\\bf Spectral signal of planet in the closure phase.} \\emph{Top panel:} theoretical near-IR spectrum of the extrasolar planet orbiting 51\\,Peg displaying CO and H$_2$O absorption bands (Sudarsky, Burrows \\& Hubeny 2003). \\emph{Bottom panel:} closure phases in milliradian calculated by us based on the theoretical spectrum of 51\\,Peg\\,b shown above as well as on the stellar spectrum for a simulated observation with the near-IR instrument AMBER/VLTI using the three telescopes UT1, UT3 and UT4 (baseline length 102m, 62m, 130m). It is evident that the spectrally resolved closure phases contain a wealth of spectral information of the planet.} \\end{figure}  In order to gain deeper insights into the physics of extrasolar planets and directly detect e.g. their atmospheres, observations with milliarcsec spatial resolution and at an unprecedented dynamical range are required (e.g. 10$^{-5}$ in the near-IR for 51\\,Peg). Such observations are coming into reach with the new generation of ground-based near-IR interferometers, like the VLTI. Observations of high contrast systems, like extrasolar planets, require an excellent calibration of atmospheric distortions. It has been proposed to observe extrasolar planets and their spectra through the differential phase or closure phase method (e.g. Quirrenbach \\& Mariotti 1997, Akeson \\& Swain 1999, Quirrenbach 2000, Lopez \\& Petrov 2000, Segransan 2001, Vannier et al. 2004). An overview of the mathematical basis for the observations of interferometric phases of extrasolar planets has also been provided by Meisner (2001) and in particular on the web page referenced therein.  Differential phases are measured between different wavelength bands, which, in a first order approximation, are subject to the same optical path length fluctuations introduced by atmospheric turbulence. The resulting differential phase errors can therefore be eliminated in the data reduction to a certain degree. However, higher order effects due to random dispersion in the atmosphere as well as in the delay lines (caused mainly by variations of the humidity of the air) can be a significant source of systematic errors and might prevent one from reaching the necessary precision to observe differential phases of extrasolar planets.  A more promising approach is the observation of closure phases, which are measured for a closed chain of three or more telescopes. The closure phase $\\Phi_\\mathrm{cl}$ is the sum of the individual single-baseline phases $\\Phi_{ij}$ measured on the baseline between telescope $i$ and $j$ of the array: \\begin{eqnarray*} \\Phi_\\mathrm{cl} = \\Phi_{12} + \\Phi_{23} + \\Phi_{31}  . \\end{eqnarray*}  The closure phase is independent of atmospheric phase errors and allows therefore a self-calibration, which was first recognized by Jennison (1958).  Closure phase observations from the ground will be possible with the AMBER (Astronomical Multi-BEam combineR) instrument at the VLTI operated by ESO (Petrov et al. 2003, Malbet et al. 2004) . AMBER is operating in the near-IR J, H, and K bands from 1.1 to 2.4$\\mu$m. It is currently undergoing commissioning at Cerro Paranal and is scheduled to be available for regular observations in 2005. With a multi-beam combiner element and a fringe dispersing mode, it has the capability of combining the light of three telescopes (UTs or ATs), and performing spectrally resolved differential and closure phase observations. There are three different spectral resolutions available: low (R=$\\lambda$/$\\Delta \\lambda$=35), medium (R=500\\dots1000) and high (R=10000\\dots15000).  We have modeled the response of the AMBER instrument at the VLTI for known extrasolar planetary systems taking into account their theoretical spectra as well as the geometry of the VLTI. The closure phase depends not only on the contrast ratio between planet and star but also on the (generally unknown) system geometry as well as on the interferometer geometry and the hour angle of observation. In the following, we present a method to determine the planetary system geometry and the planet spectrum from closure phase observations in a deterministic way without any a priori assumptions.   \\begin{figure}[t] \\begin{center} \\includegraphics[width=7.2cm,angle=90]{vjoergens_f2.eps} \\end{center} \\caption[example] { \\label{fig:nulls} {\\bf Relation between the closure phase and the corresponding single-baseline phases.} The figure shows simulated phase measurements for a system comprised of two point sources with a 3 milliarcsec separation and a constant flux ratio of the components of 10$^{-3}$ observed at zenith with the VLTI telescopes UT1, UT3 and UT4. The three top panels display the phases calculated for the single-baselines B$_{12}$, B$_{23}$ and B$_{31}$. (We note that they cannot be measured without a phase reference (e.g. PRIMA) and are displayed here only for theoretical considerations.) The bottom panel shows the simulated closure phase for these three telescopes. All phases are in millirad. It is evident that the nulls in the closure phase are always also nulls in the single-baselines phases. In particular, every second null of a single-baseline phase is also a null in the closure phase. } \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410277_arXiv.txt": {
        "abstract": "The recently suggested swing interaction between fast magnetosonic and Alfv\\'en waves  \\citep{paper1} is generalized to inhomogeneous media. We show that the fast magnetosonic waves propagating across an applied non-uniform magnetic field can parametrically amplify the Alfv\\'en waves propagating along the field through the periodical variation of the Alfv\\'en speed. The resonant Alfv\\'en waves have half the frequency and the perpendicular velocity polarization of the fast waves. The wavelengths of the resonant waves have different values across the magnetic field, due to the inhomogeneity in the Alfv\\'en speed. Therefore, if the medium is bounded along the magnetic field, then the harmonics of the Alfv\\'en waves, which satisfy the condition for onset of a standing pattern, have stronger growth rates. In these regions the fast magnetosonic waves can be strongly {\\lq absorbed\\rq}, their energy going in transversal Alfv\\'en waves. We refer to this phenomenon as {\\lq {\\it Swing Absorption}\\rq}. This mechanism can be of importance in various astrophysical situations. ",
        "introduction": "Wave motions play an important role in many astrophysical phenomena. Magnetohydrodynamic (MHD) waves may transport momentum and energy, resulting in heating and acceleration of an ambient plasma. A variety of waves have recently been detected in the solar atmosphere using the SOHO and TRACE spacecraft. Hence, an understanding of the basic physical mechanisms of excitation, damping and the interaction between the different kinds of MHD wave modes is of increasing interest \\citep{rob1}. Formally speaking, there is a group of direct mechanisms of wave excitation by external forces (e.g.\\ turbulent convection, explosive events in stellar atmospheres, etc.) and wave dissipation due to non-adiabatic processes in a medium (such as viscosity, thermal conduction, magnetic resistivity, etc.). There is a separate group of wave amplification and damping processes due to resonant mechanisms \\citep{goos,poed}. This means that particular wave modes may be damped (amplified) due to energy transfer (extraction) into or from other kinds of oscillatory motions, even when wave dissipation is excluded from consideration.  Recently, a new kind of interaction between different MHD wave modes, based on a parametric action, has been suggested \\citep{zaq,paper1,zaq0,zaq2,zaq3}. In this case the mechanism of wave interaction originates from a basic physical phenomenon known in classical mechanics as {\\lq parametric resonance\\rq}, occurring when an external force (or oscillation) amplifies the oscillation through a periodical variation of the system's parameters. The mechanical analogy of this phenomenon is a mathematical pendulum with periodically varying length. When the frequency of the length variation is twice the frequency of the pendulum oscillation, the amplitude of the oscillation grows exponentially in time. Such a mechanical system can consist of a pendulum (transversal oscillations) with a spring (longitudinal oscillations). A detailed description of such a mechanical system is given in \\citet{paper1} (hereinafter referred to as Paper I).  Here we present a brief report on our recent results on swing interactions of fast and Alfv\\'en modes in inhomogeneous media (Shergelashvili et al. (2004)). ",
        "conclusions": "The most important characteristic of swing absorption is that the velocity polarization of the amplified Alfv\\'en wave is strictly perpendicular to the velocity polarization (and propagation direction) of fast magnetosonic waves. This is due to the parametric nature of the interaction. For comparison, the well-known resonant absorption of a fast magnetosonic wave can take place only when it does not propagate strictly perpendicular to the magnetic flux surfaces and the plane of the Alfv\\'en wave polarization. In other words, the energy in fast magnetosonic waves propagating strictly perpendicular (i.e.\\ $k_\\bot=0$) to the magnetic flux surfaces cannot be resonantly {\\lq absorbed\\rq} by Alfv\\'en waves with the same frequency polarized in the perpendicular plane. This is because the mechanism of resonant absorption is analogous to the mechanical pendulum undergoing the direct action of an external periodic force. This force may resonantly amplify only those oscillations that at least partly lie in the plane of force. On the contrary, the external periodic force acting parametrically on the pendulum length may amplify the pendulum oscillation in any plane. A similar process occurs when the fast magnetosonic wave propagates across the unperturbed magnetic field. It causes a periodical variation of the local Alfv\\'en speed and thus affects the propagation properties of the Alfv\\'en waves. As a result, those particular harmonics of the Alfv\\'en waves that satisfy the resonant conditions grow exponentially in time. These resonant harmonics are polarized perpendicular to the fast magnetosonic waves and have half the frequency of these waves. Hence, for standing fast magnetosonic waves with frequency $\\omega_n$, the resonant Alfv\\'en waves have frequency ${\\sim}\\omega_n/2$.  In a homogeneous medium all resonant harmonics have the same wavelengths (see Paper~I). Therefore, once a given harmonic of the fast and Alfv\\'en modes satisfies the appropriate resonant conditions (Eqs. (23) and (25) in Paper~I), then these conditions are met within the entire medium. Thus, in a homogeneous medium the region where fast modes effectively interact with the corresponding Alfv\\'en waves is not localized, but instead covers the entire system. However, when the equilibrium is inhomogeneous across the applied magnetic field, the wavelengths of the resonant harmonics depend on the local Alfv\\'en speed. When the medium is bounded along the unperturbed magnetic field (i.e.\\ along the $z$ axis), the resonant harmonics of the standing Alfv\\'en waves (whose wavelengths satisfy condition (\\ref{rescondm}) for the onset of a standing pattern) will have stronger growth rates. This means that the {\\lq absorption\\rq} of fast waves will be stronger at particular locations across the magnetic field. In the previous section we showed numerical solutions of standing fast magnetosonic modes for a polytropic equilibrium ($p_0 \\sim \\rho _0 ^\\gamma$) in which the thermal pressure and magnetic pressure are linear functions of $x$. Further, we performed a numerical simulation of the energy transfer from fast magnetosonic waves into Alfv\\'en waves at the resonant locations, i.e.\\ the regions of {\\it swing absorption}.  The mechanism of swing absorption can be of importance in a variety of astrophysical situations."
    },
    "0410/astro-ph0410088_arXiv.txt": {
        "abstract": "We carry out a detailed numerical investigation of the spatial correlation function of the initial positions of cosmological dark matter halos.  In this Lagrangian coordinate system, which is especially useful for analytic studies of cosmological feedback, we are able to construct cross-correlation functions of objects with varying masses and formation redshifts and compare them with a variety of analytical approaches.   For the case in which both formation redshifts are equal,  we find good agreement between our numerical results and the bivariate model of  Scannapieco \\& Barkana (2002; SB02) at all masses, redshifts, and separations, while the model of Porciani \\etal (1998) does well for all parameters except for objects with different masses at small separations.  We find that the standard mapping between Lagrangian and Eulerian bias performs well for rare objects at all separations, but fails if the objects are highly-nonlinear (low-sigma) peaks.  In the Lagrangian case in which the formation redshifts differ, the SB02 model does well for all separations and combinations of masses, apart from a discrepancy at small separations in situations in which the smaller object is formed earlier and the difference between redshifts or masses is large.  As this same limitation arises in the standard approach to the single-point progenitor distribution developed by Lacey \\& Cole (1993), we conclude that a more complete understanding of the progenitor distribution is the most important outstanding issue in the analytic modeling of Lagrangian bias. ",
        "introduction": "The universe is a complicated place.  Not only does cosmological structure formation involve detailed questions of fluid dynamics, radiative transfer, and chemistry, but most of these issues are non-local.  The impact of high-redshift supernovae (SNe), for example, is not limited to the host-galaxies in which they are born. Theoretical arguments show that the large numbers of SNe formed in high redshift dwarf galaxies drove global winds of such force that they not only reduced further infall onto their hosts, but also suppressed the formation of neighboring galaxies (\\eg Scannapieco, Ferrara, Broadhurst 2000; Thacker, Scannapieco, \\& Davis 2003). Observational support of this idea is strong, with large numbers of massive outflowing starbursts directly detected at $z \\gtrsim 3$ (Pettini \\etal 2001; Frye, Broadhurst, \\& Benitez 2002; Hu \\etal 2002).  Such  observations of outflows are complemented by hints that the intergalactic medium (IGM) surrounding high-redshift starbursts has been excavated by outflows out to distances $\\sim 1$ Mpc (Adelberger \\etal 2003) and that the distribution of intergalactic metals is biased and highly clustered around similar sources (Pichon \\etal 2003).  Furthermore, quasar absorption line studies have uncovered an inhomogeneous distribution of heavy elements that is already pervasive throughout the intergalactic medium by $z \\approx 5$ (\\eg Rauch, Haehnelt, \\& Steinmetz 1997; Songaila 2001; Schaye \\etal 2003; Aracil \\etal 2003).  This large-scale enrichment further complicates the structure formation  process.  Metal-line cooling is of crucial importance to the formation history of large galaxies (\\eg Sutherland \\& Dopita 1993), as low levels of enrichment can produce an order of magnitude change in cooling rates above $10^5$ K. Also, the presence of even trace levels of heavy elements in protogalactic material may be sufficient to drastically affect the shape of the stellar initial mass function (\\eg Bromm \\etal 2001; Schneider \\etal 2002).  But the most highly non-local processes are undoubtedly radiative. The earliest such example is the dissociation of primordial H$_2,$ which regulates the formation of the first generation of galaxies in halos with virial temperatures below $10^4$ K.  As ${\\rm H}_2$ is easily photo-dissociated by 11.2-13.6 eV photons to which the universe is otherwise transparent, even small numbers of forming stars have the potential to suppress dwarf galaxy formation in large regions of space (\\eg Haiman, Rees, \\& Loeb 1997), motivating detailed analysis of ${\\rm H}_2$ formation by X-rays (\\eg Oh 2001) and self-shielding (Yoshida \\etal 2003).  Secondly, there is the question of reionization.  Here the classic picture is of a two-stage process, which begins with sources ionizing their immediate surroundings and ends in a rapid ``overlap'' phase as H~II regions join together, quickly ionizing the remaining neutral regions (\\eg Gnedin 2000).  At these times feedback due to photo-evaporation is of primary importance to galaxy formation in $T_{\\rm vir} = 10^4-10^{4.5}$K halos (\\eg Barkana \\& Loeb 1999), a non-local effect that is only worsened by the fact that similar low-mass galaxies are likely to be the source of the ionizing background in the first place.  This wide range of mechanical, chemical, and radiative feedback processes has necessarily sparked a multifaceted numerical attack on non-local cosmological problems. Recent numerical simulations have included detailed studies of mechanical feedback from galaxy outflows (\\eg Scannapieco, Thacker, \\& Davis 2001; Springel \\& Hernquist 2003), the chemical properties of galaxies and the intergalactic medium (\\eg Thacker, Scannapieco, \\& Davis 2002; Theuns \\etal 2002; Tassis \\etal 2002) and the formation and propagation of cosmological ${\\rm H}_2$ dissociation and ionization fronts (\\eg Ricotti, Gnedin, \\& Shull 2002; Ciardi, Ferrara, \\& White 2003; Yoshida \\etal 2004).  These highly non-local studies, unlike earlier investigations, are far outstripping the power of approximate analytic techniques.  Thus, in our rush to match the dramatic recent observational advances in cosmology with equally sophisticated numerical work, there is a very real threat that we will soon lack the tools to understand  either in a more fundamental context.  The most widely applied analytic method for determining the distribution of halos was first developed by Press \\& Schechter (1974, hereafter PS), and later refined by Bond \\etal (1991), Lacey \\& Cole (1993), and Sheth, Mo, \\& Tormen (2001, hereafter SMT01) among others. Yet, despite its widespread utility, this method is limited in that it can only predict the average number density of virialized dark matter halos, and does not supply any information about their relative positions. However, when used in combination with simulations, this model is exceptionally useful since it allows rapid investigation of parameter space. In the absence of this possibility for studying non-local effects, numerical investigations are thus on a surprisingly shaky footing: any such investigation is computationally intensive, and is able to consider only a limited range of possible scenarios and parameters. Yet, at the same time, most processes depend on a large number of uncertain parameters, almost assuring that the simulations will not be able to bracket the full range of physically plausible scenarios.  Finally, even in the cases in which simulations are able to assess these uncertainties, the lack of an analytical counterpart frequently leaves us  without a basic theoretical understanding.  These limitations are somewhat mitigated by analytical studies of the two-point correlation function,  although only recently  have these approaches been  extended to the point of being useful for  feedback. Considerable effort has gone into to modeling the biased clustering of halos at a given epoch in the Eulerian coordinate system in which they are observed, both in mildly-nonlinear (see Bernardeau \\etal 2002 for a review) and highly nonlinear contexts (\\eg Mo \\& White 1996, hereafter MW96; Seljack 2000; Ma \\& Fry 2000). Limited models have also been developed in the Lagrangian coordinate system in which the halos cell-centers were originally located (Porciani \\etal 1998; Sheth \\& Lemson 1999).  In this  reference frame, the distance between two objects is closely related to the total mass of material between them, rather than their final comoving distance.  The usefulness of such a reference frame is twofold.  Firstly, it is the natural frame for PS-type analytical calculations, which  associate peaks  in the {\\em initial} density field with collapsed objects at various redshifts.  Secondly, it allows for a more accurate treatment of the propagation of disturbances such as shocks (Scannapieco \\etal 2003) or ionization fronts (Iliev \\etal 2004) between objects, as these are largely dependent on the total column depth of material separating two perturbations  rather than their precise distance in physical space.   However, early analytical approaches were unable capture the Lagrangian clustering between halos forming at different epochs, a quantity of crucial importance to modeling non-local effects, as cosmological disturbances take a significant amount of time to propagate from their sources to neighboring objects.  In an attempt to alleviate this shortcoming, we developed in Scannapieco \\& Barkana (2002; hereafter SB02) an approximate analytical model that considered the collapse of two neighboring points of arbitrary mass and formation redshift, separated by an arbitrary Lagrangian distance. Working in the context of the linear excursion-set formalism of Bond \\etal (1991), we showed that while the exact two-point solution could not be obtained analytically, this solution could be reproduced to high accuracy by introducing a simple approximation.  The resulting expressions then interpolated smoothly between all standard analytical limits: reducing for example to the standard halo bias expression described in MW96 in the limit of equal-mass halos at the same redshift, and reproducing the Lacey \\& Cole (1993) progenitor distribution in the limit of different-mass halos at the same position and different redshifts.  However, just as numerical results must be placed in a larger analytical context, the analytical results obtained in SB02 were  left incomplete without an accurate numerical assessment of their strengths and limitations.  In this study, we carry out just such a comparison, computing detailed numerical Lagrangian halo-halo correlations at a variety of masses and redshifts and comparing them directly with SB02 and other analytical models.  The structure of this work is as follows. In \\S 2 we describe our numerical simulation and group finding techniques and compare the resulting mass distributions with standard analytical results.  In \\S 3 we systematically compare the Lagrangian correlation function of groups with the analytical model of SB02 for a  variety of formation redshifts and halo masses. We conclude in \\S 4 with a  brief review of our results and a short discussion. ",
        "conclusions": "Despite the widespread use of single-point ``merger-tree'' approaches in analytical studies of galaxy formation, nonlocal processes abound in cosmology. Starburst galaxies have drove winds strong enough to strip the gas out of neighboring density perturbations.  The resulting inhomogeneous enrichment of the IGM profoundly affected the cooling properties of later-forming galaxies,  perhaps altering their stellar initial mass function. The ultraviolet radiation from  the earliest objects strongly inhibited the formation of similar such sources. Indeed, the moment the universe emerged from the cosmological ``dark ages'' it was plunged into an era of strong, nonlocal feedback that continues to this day.  In this study, we have used high-resolution N-body simulations to examine the Lagrangian dark-matter halo correlation function, a quantity that is uniquely suited to studying nonlocal processes because: (1) it is in the natural coordinate systems for PS-type analytical calculations; (2) it allows for a more accurate treatment of the propagation of cosmological disturbances, which are largely dependent on the total column depth of material  separating two perturbations. Comparing our results with a number of analytic expressions, we find that the bivariate mass function  derived in SB02 can be used to accurately reproduce the  Lagrangian correlation function of halos with different  masses and formation redshifts.  This is a significant step forward,  which confirms that non-local processes can be modeled within the framework of the bivariate mass function.  In the case of single mass, single redshift halos, the SB02 model and the P98 model both reproduce the correlation function behavior well. As  expected we confirm the failure of the MW96 model for low $\\nu$ and small  radii ($r< 4$ Mpc), but the $b_E=b_L+1$ mapping proposed by Jing works well  for all masses and radii, unless $\\nu$\\lsim$1.5$.  At this point the mapping  breaks down, and a more sophisticated technique is needed to deal with  the short-range physics that comes into play.  For objects of different masses at the same redshift, the P98 model and the  SB02 model have close agreement to a limiting radius that  is defined by the turn-over in the correlation function due to the anticorrelation required at zero separation. The P98 model does not correctly reproduce this turn-over, due to the omission of the appropriate barrier treatment for the formation of structure within peaks, while the SB02 model exhibits the correct short-range  behavior for all our studied cases of different masses and redshifts.   The different mass, same redshift Eulerian correlation functions for the SB02 and P98 models exhibit similar performance issues to the same mass, same redshift comparison. The problems are particularly acute for the SB02 model where the barrier treatment introduces a strong kink in the Eulerian correlation function. In these cases, remarkably, the Eulerian MW96 model works best of all three models.  The two redshift, equal mass Lagrangian case also shows the SB02 model to perform well. The growth of halo mass with redshift forces anticorrelation of equal mass halos at different redshifts and small radii.  The correct turnover in the correlation function for small radii is  observed, and the position of the knee is accurately predicted for all  redshift cases.  Unsurprisingly, our most interesting results come from the most difficult  case to study, namely the correlation of halos with different masses and  formation redshifts. In the case in which the larger mass is at the higher redshift, the SB02 model reproduces the expected anticorrelation extremely  accurately. In the more relevant case of the lower mass at higher redshift  the Lacey \\& Cole (1993) progenitor distribution is always reproduced exactly at small radii, and this distribution compares well with the simulations,  provided that the differences in redshift and mass between the halo are not too large ($M_2/M_1<10$ and $z_1-z_2<3$).  For larger differences in mass and redshift, however, both the Lacey \\& Cole (1993) and SB02 models exhibit significant discrepancies from the numerical values measured here and reported elsewhere.  At larger radii, however, the SB02 model reproduces the expected  behavior of the correlation function in all cases.  It is intriguing that the {\\em single-point} case  remains the single most important outstanding issue in our analytic modeling of Lagrangian bias."
    },
    "0410/astro-ph0410041_arXiv.txt": {
        "abstract": "We investigate non-linear scaling relations for two-dimensional gravitational collapse in an expanding background using a 2D TreePM code and study the strongly non-linear regime ($\\bar\\xi \\leq 200$) for power law models. Evolution of these models is found to be scale-invariant in all our simulations. We find that the stable clustering limit is not reached, but there is a model independent non-linear scaling relation in the asymptotic regime. This confirms results from an earlier study which  only probed the mildly non-linear regime ($\\bar\\xi \\leq 40$). The correlation function in the extremely non-linear regime  is a less steep function of scale than reported in earlier studies. We show that this is due to coherent transverse motions in massive haloes. We also study density profiles and find that the scatter in the inner and outer slopes is large and that there is no single universal profile that fits all cases. We find that the difference in typical density profiles for different models is smaller than expected from similarity solutions for halo profiles and transverse motions induced by substructure are a likely reason for this difference being small. ",
        "introduction": "Large-scale structures such as clusters of galaxies are very over-dense compared to the average density of the universe. It is believed that these were formed by the growth of small perturbations via gravitational instability \\cite{2002PhR...367....1B,2002tagc.book.....P,1999coph.book.....P,1980lssu.book.....P}. In general, the equations that describe the growth of density perturbations due to gravitational clustering \\cite{1974ApJ...189L..51P,1980lssu.book.....P} cannot be solved analytically when the over-densities are large and one has to rely on N-Body simulations for detailed predictions. Limited information about non-linear gravitational clustering can be obtained by using approximations or ansatze, e.g., stable clustering \\cite{1980lssu.book.....P} for which it is assumed that the infall velocity and Hubble velocity compensate for each other leading to a stable configuration. Stable clustering allows us to relate the initial spectrum of fluctuations with the final asymptotic spectrum of fluctuations \\cite{1977ApJS...34..425D}. Non-linear scaling relations are a more detailed prescription for relating the initial and final spectrum of fluctuations \\cite{1991ApJ...374L...1H,1994MNRAS.271..976N,1994MNRAS.267.1020P,1995MNRAS.276L..25J,1996MNRAS.280L..19P,1996MNRAS.278L..29P,1996ApJ...466..604P,1997MNRAS.286.1023B,1998ApJ...495...25B,2003MNRAS.341.1311S}. The motivation for studying non-linear scaling relations is to understand key phases in gravitational clustering, and to identify the relevant process in each phase \\cite{1996MNRAS.278L..29P}.  Existence of non-linear scaling relations implies that gravitational clustering does not erase memory of initial conditions. This provides another useful and a more prosaic motivation for studying scaling relations, namely to be able to make predictions about clustering of galaxies for a given power spectrum \\cite{1991ApJ...374L...1H}. Scaling relations can also be used to make predictions for weak lensing observations in a given model \\cite{1992ApJ...388..272K,1997ApJ...484..560J}. These scaling relations are valid for hierarchical models where the initial conditions contain fluctuations at all scales and the amplitude of fluctuations increases monotonically as we go from large-scales to small-scales.  Non-linear scaling relations indicate that there are three prominent regimes in evolution of gravitational clustering. When the amplitude of perturbations is small so that the density contrast ($\\delta = \\rho/{\\bar\\rho} - 1$) is close to zero ($|\\delta| \\ll 1$), mode coupling is not important and the evolution closely follows the predictions of linear perturbation theory \\cite{1974ApJ...189L..51P}. The power spectrum and correlation function evolve without a change in shape in the linear regime. As the density contrast grows and becomes comparable to unity, motions induced by gravitational collapse start to dominate over the expansion of the universe. The quasi-linear regime is dominated by infall onto density peaks and the correlation function grows rapidly in this phase \\cite{1996MNRAS.278L..29P}. Gravitational collapse leads to formation of structures in or close to dynamical equilibrium and further evolution of density contrast is dominated by depletion of average density due to expansion of the universe as the density of collapsed structures remains almost constant. We shall call this the asymptotic regime.  Initial conditions with power law power spectra are very useful for such studies. These power spectra have the form $P(k) \\propto k^n$ with $n$ being the index of the power spectrum. We require $(n+D) > 0$ for hierarchical clustering and N-Body simulations can be used for only this type of models, indeed even for these models there are restrictions on the size of the simulation box in order for the simulation to reproduce the model correctly \\cite{2004astro.ph.10373B}. Here $D$ is the number of spatial dimensions in which gravitational clustering is being studied. If we consider an Einstein-deSitter universe then such a model is characterised by the scale of non-linearity $r_{NL}$ and the index $n$, where $r_{NL}$ is defined as the scale at which rms fluctuations in linearly extrapolated density contrast ($\\sigma_L(r)$) is unity. \\begin{eqnarray} \\sigma_L^2(r) ~ &=& ~ a^2 ~ r^{-\\left(n+D\\right)} \\nonumber \\\\ r_{NL} ~ &=& ~ a^{2/\\left(n+D\\right)} \\end{eqnarray} If the evolution of clustering is self-similar then the non-linear correlation function plotted as a function of $r/r_{NL}$ should have the same shape at all epochs for a given index $n$. Note that the evolution is expected to be self-similar if $r_{NL}$ is the only scale in the problem and it is not expected in cosmologies with non-zero spatial curvature or a non-power-law expansion factor.  Several studies have attempted to understand the nature of the asymptotic regime, mainly with help of N-Body simulations \\cite{1994MNRAS.267.1020P,1995MNRAS.276L..25J,1996ApJ...465...14C,1996ApJ...466..604P,1996MNRAS.280L..19P,1997MNRAS.287..687J,2000ApJ...531...17K,2001ApJ...550L.125J,2003MNRAS.341.1311S}. Key conclusions of these studies about the asymptotic regime may be summarised as follows. \\begin{itemize} \\item Stable clustering is not reached in the range of non-linearities explored by N-Body simulations. However, the departure from stable clustering for most models is small. \\item Gravitational clustering does not erase memory of initial conditions, i.e., the non-linear power spectrum is not independent of the linear power spectrum to the extent it could be probed. There is no {\\it universal} asymptotic regime. \\end{itemize} One of the reasons for the inability to resolve the problem of asymptotic regime has been the limited dynamic range of N-Body simulations. The problem of dynamic range can be circumvented by simulating a two-dimensional system \\cite{1985A&A...144..413B,1999ApJ...512....1E,2000PhRvD..61b3515P} instead of a three-dimensional one, wherein a much higher dynamic range can be achieved with similar computational resources \\cite{1998ApJ...495...25B,1998MNRAS.293L..68M}. Generic features like non-linear scaling relations are likely to be independent of dimension and --- in fact --- the scaling relations in the quasi-linear regime were predicted \\cite{1996MNRAS.278L..29P} well before these could be tested in simulations \\cite{1998ApJ...495...25B}. If we can understand the nature of the asymptotic regime in two dimensions, it will help us solve the problem in three dimensions even if we cannot map the solution directly to the full problem in three dimensions. Previous studies of gravitational clustering in two dimensions concluded that there is no stable clustering \\cite{1998ApJ...495...25B,1998MNRAS.293L..68M}. Both the studies mentioned above were limited to $\\bar\\xi \\leq 40$ and the dynamic range in the non-linear regime was limited compared to the present study (see eqn.(6) for the definition of $\\bar\\xi$). In this work we revisit the same issues, but using the 2D TreePM code \\cite{2004astro.ph..6009R} for N-body simulations. The TreePM code has a better force resolution as compared to a PM code \\cite{1985A&A...144..413B}, therefore in our simulations we have a significantly larger dynamic range over which we can study the asymptotic regime.  A related issue is that of density profiles of haloes surrounding galaxies and clusters. Non-linear scaling relations in the asymptotic regime depend on the kind of dynamical equilibrium that is reached in massive haloes, if an equilibrium is ever reached. Under the assumption of self-similar collapse, the density profiles of massive haloes can be related to the initial conditions, and also to the dynamics within these haloes \\cite{1984ApJ...281....1F,2000ApJ...538..517S}. If these models are true then the slope of density profile at small scales is related to the index of power spectrum as: \\begin{equation} \\rho(r) \\propto r^{-\\eta} ~~~~~ ; ~~~ \\eta ~ = ~ \\frac{D ~ \\left( n + D + 1 \\right)}{2 + D + n} \\end{equation} If self-similar collapse of haloes is the correct description then the final density profile of a halo carries information about the initial profile and hence about the power spectrum \\cite{1986ApJ...304...15B}. It has been claimed that gravitational clustering in hierarchical models leads to a {\\it universal} density profile \\cite{1996ApJ...462..563N} irrespective of the initial conditions, we shall refer to this density profile as the NFW profile. \\begin{equation} \\rho_{_{\\rm NFW}}(r) \\propto \\frac{1}{(r/r_s) ( 1 + r/r_s)^2} \\label{NFW} \\end{equation} Here $r_{s}$ is a scale radius. We will also use the term $r_{200}$ for this. This profile is characterised by a $r^{-3}$ decline at large radii and a cuspy inner profile of the form $\\rho(r) \\propto 1/r$. The claim of universality has been tested in several studies \\cite{1996ApJ...462..563N,1996MNRAS.281..716C,1997ApJ...477L...9F,1997MNRAS.286..865T,1998ApJ...499L...5M,1998ApJ...502...48K,2000ApJ...544..616G,2000ApJ...538..528S,2001ApJ...554..903K,2003MNRAS.338...14P,2004ApJ...606..625F,2004MNRAS.349.1039N} and though there are disagreements, it appears that the NFW profile is {\\it consistent} with the density profiles of massive haloes in N-Body simulations even if it is not the best fit. This essentially implies that there is a scatter in density profiles and one can fit different functional forms to the N-Body data \\cite{2000ApJ...535...30J,2001MNRAS.321..559B}.  If true, the existence of NFW like profile implies that there are some universal aspects of gravitational clustering in an expanding universe. There is no {\\it ab initio} derivation of the NFW profile, but one can argue that if there is a universal profile it should have same asymptotes as the form shown above \\cite{1998MNRAS.293..337S,1999MNRAS.303..685N,2002LNP...602..165P}.  We briefly review non-linear scaling relations in the following section. Numerical simulations used in the present work are discussed in \\S{3}. We conclude in \\S{4}. ",
        "conclusions": "The basic motivation for this work was to improve our understanding of non-linear scaling relations for two dimensional gravitational clustering. Previous studies \\cite{1998ApJ...495...25B,1998MNRAS.293L..68M} established the behaviour for the quasi-linear regime, but could not probe the nature of the asymptotic regime. We used the 2D TreePM code \\cite{2004astro.ph..6009R} which allows us to study the non-linear regime in far greater detail, as compared to the PM code used in earlier studies. The results of our study can be summarised as follows: \\begin{itemize} \\item The evolution of the correlation function for power-law models is scale invariant up to the highest non-linearities reached in our simulations. \\item We reproduce the results of earlier studies in the quasi-linear regime. We confirm that the slope of the non-linear scaling relation in this regime is close to $2$, as predicted on the basis of the infall dominated growth model \\cite{1996MNRAS.278L..29P}. \\item We do not find any difference in the non-linear scaling relation for different power law models. Thus the information about initial conditions is retained even in the extremely non-linear regime. \\item The stable clustering limit ($h \\rightarrow 1$ as $\\bar\\xi \\rightarrow \\infty$) is not reached. (Though it is generally claimed that stable clustering with $h\\to 1$ is a {\\sl natural} asymptotic state, it has been shown that the Davis-Peebles scale-invariant solution \\cite{1977ApJS...34..425D} and the hierarchical model for the three-point function are inconsistent with the standard stable-clustering picture \\cite{2000ApJ...531...17K}. This, of course, is for gravitational collapse in three dimensions.) This can be concluded from the slope of the correlation function, pair velocity, as well as density profiles of massive haloes. \\item Pair velocities in the asymptotic regime are smaller than expected from the stable clustering model. This does not imply that massive haloes are evaporating. \\item The transverse component of pair velocity is more significant for models with a larger spectral index, i.e., for models with more substructure. This suggests that gravitational collisions between substructure orbiting within haloes are responsible for generating coherent transverse motions. \\item We find that there is no universal density profile for massive haloes in two dimensional gravitational clustering. There is a large scatter in inner as well as outer slopes of density profiles. Median value of the inner slope is different for different models, though the difference is much smaller than the scatter in values. Tests show that these clusters are close to dynamical equilibrium. \\item The difference in the median inner slope for different models is smaller than expected from self-similar spherical collapse. We argue that the difference in significance of transverse motions for different models is responsible for this change. \\end{itemize}  Many of the conclusions listed here have been checked independently in large PM simulations ($4096^2$ particles) up to $\\bar\\xi \\leq 100$.  Further work is  required to test our hypothesis that gravitational collisions between substructure are responsible for generating coherent transverse motions. We are carrying out a detailed study using numerical simulations where we tune the amount of substructure to test this proposal.  Considerable work has been done to study the non-linear scaling relations in three dimensions \\cite{1991ApJ...374L...1H,1994MNRAS.271..976N,1994MNRAS.267.1020P,1995MNRAS.276L..25J,1996MNRAS.280L..19P,1996MNRAS.278L..29P,1996ApJ...466..604P,1998ApJ...495...25B,2003MNRAS.341.1311S}. We intend to study these again in light of insight developed using tests in clustering in two dimensions. It has been shown that the scaling relations for gravitational clustering in three dimensions depend on the initial spectrum of fluctuations \\cite{2003MNRAS.341.1311S}. It is important to understand the origin of this index dependence and test whether substructure plays a role in this. The transverse pair velocity function has been introduced in eqn.(\\ref{geqn}), we are studying the three dimensional equivalent of this quantity in order to probe the role of transverse motions in gravitational clustering in the full three dimensional problem. These studies should help us understand the non-linear scaling relations better."
    },
    "0410/astro-ph0410727_arXiv.txt": {
        "abstract": "{ A statistical study of the properties of molecular outflows is performed based on an up-to-date sample. 391 outflows were identified in published articles or preprints before February 28, 2003. The parameters of  position, morphology, mass, energy, outflow dynamics and  central source luminosity are presented for each outflow source. Outflow lobe polarity is known for all the sources,  and 84\\% are found to be bipolar. The sources are divided into low mass and high mass groups according to either the available bolometric luminosity of the central source or the outflow mass. The pace of discovery of outflows over the past seven years has increased much more rapidly than in previous periods. Surveys for outflows are still continuing. The number of high-mass outflows detected (139) has considerably increased, showing that they are commonly associated with massive as well as low mass stars. Energetic mass ejection may be a common aspect of the formation of high mass as well as low mass stars. Outflow masses are correlated strongly with bolometric luminosity of the center sources, which was obtained for the first time. There are also correlations between the central source luminosity and the parameters of mechanical luminosity and the thrust or force necessary to drive the outflow. The results show that flow mass, momentum and energy depend on the nature of the central source. Despite their similarity, there are differences between the high mass and low mass outflows. Low mass outflows are more collimated than high mass outflows. On average, the mass of high mass sources can be more than two orders of magnitude larger than those of low mass outflows. The relation between flow mass and dynamical time appears to differ for the two types of outflows. Low mass sources make up 90\\% of outflows associated with HH objects while high mass outflows make up 61\\% of the sources associated with H$_2$O masers. Sources with characteristics of collapse or infall comprise 12\\% of the entire outflow sample. The spatial distribution of the outflow sources in the Galaxy is presented and the local occurrence rate is compared with the stellar birth rate. ",
        "introduction": "As one of the most exciting discoveries in astronomy in the past three decades, high velocity molecular outflows continue to attract great attention from researchers. Since their first discovery in 1976~\\citep{Zuckermann_KK76, Kwan_S76, Zuckerman_P75}, outflows have been detected at a high rate. 67 had been discovered  by 1984 ~\\citep{Lada85}, 144 by 1989~\\citep{Fukui89}, and 264 by 1996 \\citep{Wu_HH96}.  Outflows were thought to be the earliest observable signatures of star formation. Although evidence of mass loss for young stars had been detected previously in the optical spectrum \\citep{Herbig60, Kuhi64}, a better understanding of such mass loss in early evolutionary phases was only acquired after the discovery of molecular outflows in the material surrounding the embedded young stellar objects (YSOs). Molecular outflows show various characteristics such as radio and optical jets, atomic hydrogen stellar wind, molecular hydrogen jets and interstellar masers. The stars eject mass in energetic flows during formation~\\citep{Lada85, Lizano_HR88, Natta_GP88, Lane89, Felli_PT92, Russell_BP92, Matzner_M99, Bally_JJ01, Reipurth_HM02}. In the 1990s, progress was made in observing and understanding the collapse of protostars \\citep{Zhou92, Evans99, Myers_MT96}. According to theoretical models, collapses should occur before outflow~\\citep{Shu_AL87}. However, nearly all the candidates of collapsing protostars are found in outflow sources.   When advancement of observational techniques allowed the detection of high velocity molecular outflows by high J transitions, extremely high-velocity molecular outflows or highly collimated jet-like outflows were studied extensively \\citep{Koo89, Bachiller_MP91, Choi_EJ93, Narayanan_W96}.  More recently, considerable progress in outflow detection has been made in distant and complex high-mass star formation regions \\citep{Beuther_SS02a, Zhang_HB01, Shepherd_C96a, Wu_YL99}. In the surveys of high-mass star formation regions conducted by \\citet{Beuther_SS02a} and \\citet{Zhang_HB01}, the outflow detection rates reached 80\\% and 90\\%, respectively. Recent studies show that the sources in these surveys are either precursors of ultracompact (UC) HII regions, or are at very early evolutionary stages~\\citep{Molinari_TR02, Beuther_SM02b}.  The outflow parameters and their correlations have been investigated \\citep{Rodriguez_CH82, Bally_L83, Cabrit_B92, Shepherd_C96b, Ridge_M01}. Mechanisms of outflow collimation and the forces driving the outflows have been proposed as well. However, the correlation between morphology, physical parameters and the central source conditions are still not well understood. For example, it is not clear if the collimation is correlated with the driving force. The correlation among bolometric luminosity of the center source, the required force and mechanical luminosity was investigated by \\citet{Bally_L83} and \\citet{Rodriguez_CH82}. As the correlation was discovered to be weak, a larger sample is required.  To further understand the mechanisms of high velocity outflows, we conducted an analysis based on an up-to-date sample. Our calculations and conclusions are presented in this paper.  The used data in this paper are heterogeneous. The outflows were detected at different stages of the young stellar objects: high mass sources from pre-UCHII region to UCHII region and low mass sources from Class 0 to FU Orionis. The observational sensitivity and the calculation method also varied from source to source. However, we believe that our sample is sufficiently large to provide the overall physical characteristics of outflows. Sect. \\ref{catalogue} presents the sample, the catalogue and a brief analysis. The morphology analysis in Sect. \\ref{morph}. Sect. \\ref{phy} discusses the physical parameters. The associated objects or phenomena of outflows are discussed in Sect.~\\ref{associated}. Sect.\\ref{distri} describes the spatial distribution of outflows. We give a summary in Sect. \\ref{sum}. ",
        "conclusions": "}  In this paper, the searches of molecular outflow sources are reviewed. Our Catalogue includes 391 high velocity molecular outflows and their basic parameters. We  classified them into two groups according to the available bolometric luminosity or to the outflow mass. We analyzed the progress of molecular outflow searches and divided the history of outflow search into four periods: I.1976-1984, II. 1985-1989, III. 1990-1995 and IV. 1996-2003. The characters of different periods and the developmental trend are discussed. The morphology, collimation and physical parameters of outflows  were investigated. We also studied the correlation between outflows and their central sources. Finally the spatial distribution and birth rate of young stars are presented and discussed. For convenience, the main statistical results are listed in Table. 4.  Our main results are as follows:   \\begin{enumerate} \\item The catalogued outflows have been classified into high mass and low mass groups according to the bolometric luminosity of the center sources and the outflow masses. The pace of outflow identification has accelerated in the past two decades. The number of outflows detected  per year tends to increase, that is, 7 in Period I but 23 in Period IV. The number of low mass sources increased fast during period II. In period IV, the number of high mass outflows increased rapidly, suggesting that the outflow phenomenon is also common in high mass star formation regions. Data acquired recently suggest  high mass star formation processes may also go through mass accreting and ejecting, similar to those processes in low mass star formation.  \\item Polarity was known for all the 391 sources, 84\\% of which are bipolar. Red mono-polar outflows are found to be  equally abundant compared to the blue ones within the statistical uncertainty. This suggests that  monopolarity of the outflows is intrinsic and is a result of the driving source being partially submerged in the molecular cloud. The number of multiple sources has increased and isolated ones decreased since the last catalogue \\citep{Wu_HH96}, which shows the improvement of the detection technology.  \\item We analyzed the collimation of the outflows. The number of sources with known collimation factors has increased from 26 in  period I to 213 in  period IV. The current average value of the collimation factors is slightly larger than that in period I. Although the resolution of the equipment used in period IV has generally improved, the number of high mass outflows observed increased in this period. As high mass outflows have in general lower collimation factor values, the average collimation value only changed slightly.  \\item The collimation of low mass sources is better than that of high mass ones as a whole. Different outflow driving mechanisms may explain such differences. Moreover, projection effects, physical conditions of the environment and evolutionary states also affect the morphology of outflows.  \\item The relationship between  physical parameters of the outflow and its central source is examined. Analysis shows that outflow masses are correlated strongly with bolometric luminosities of young stellar sources. This relation, a typical example of an indirect correlation of two quantities, is obtained for the first time.   \\item Similar linear correlations exist between flow mechanical luminosity, the force required to drive the outflow and the bolometric luminosity in logarithmic scales. The weak correlation found by \\citet{Bally_L83} in an early survey is confirmed.  The mass, momentum and energy of the outflows depend on the luminosity of the central source. These relationships between the parameters of the outflows and the driving sources suggest that the energetic outflows originate from their central source. Although the radiation pressure could not drive the outflow directly if the photons emitted from the central source  scatter once, these correlations confirm the  relationship between outflows and their central stellar objects.  \\item Outflow mass and dynamic time tends to be correlated for low mass sources while for high mass ones there was no such correlation. This together with the differences between the low mass and high mass group  sources shown by  $L_{\\rm m}$ v.s. $L_{\\rm bol}$ (Fig. \\ref{Lbollm}) and $F$ vs. $L_{\\rm bol}$ (Fig. \\ref{LbolF}) is consistent with the presence of  different driving mechanisms.  \\item Objects associated with outflows were investigated. 136 outflows have associated  HH objects; 10\\% of those are among the high mass group; 190 outflows have water masers, 39\\%  belong to the low mass group. 12\\% of the outflow sources  have ongoing collapse of infall detected.  \\item The spatial distribution of known outflows is not symmetric with respect to the sun because the southern hemisphere has fewer identified sources, largely due to less observation. The occurrence rate within 1kpc from the sun is $74/t ~{\\rm kpc^{-2}yr^{-1}}$, where $t$ is a typical lifetime of outflow in years. For an outflow dynamic time of $5.8\\times 10^4$ years the birth rate is $1.3\\times 10^{-3}~{\\rm kpc^{-2}yr^{-1}}$, close to the local star birth rate. \\end{enumerate}  Higher angular resolution and sensitivity are necessary to further resolve the velocity structure and survey the distant sources. If a large number of outflows with more accurate morphology are found, we may better understand the physical nature of the energetic processe occuring during early stellar evolution."
    },
    "0410/astro-ph0410511_arXiv.txt": {
        "abstract": "To world-wide notice, in 2002 the ATHENA collaboration at CERN (in Geneva, Switzerland) announced the creation of order 100,000 low energy antihydrogen atoms.  Thus, the concept of using condensed antihydrogen as a low-weight, powerful fuel (i.e., it produces a thousand times more energy per unit weight of fuel than fission/fusion) for very deep space missions (the Oort cloud and beyond) had reached the realm of conceivability.  We briefly discuss the history of antimatter research and focus on the technologies that must be developed to allow a future use of controlled, condensed antihydrogen for propulsion purposes. We emphasize that a dedicated antiproton source (the main barrier to copious antihydrogen production) must be built in the US, perhaps as a joint NASA/DOE/NIH project.  This is because the only practical sources in the world are at CERN and the proposed facility at GSI in Germany. We outline the scope and magnitude of such a dedicated national facility and identify critical project milestones. We estimate that, starting with the present level of knowledge and multi-agency support, the goal of using antihydrogen for propulsion purposes may be accomplished in $\\sim50$ years. ",
        "introduction": "In this century, the development of missions  to deeper and deeper space will become an ever increasing priority. To complete a mission within a reasonable time frame, even to the nearest extra-solar system objects of interest, the Oort Cloud or the Alpha Centauri star system (4.3 light years away), the velocity of the spacecraft needs to be high, up to more than 10\\% of the speed of light.  To achieve this one needs the highest energy-density fuel conceivable.  This would be antimatter; a large amount of it and in a compact form.  Antimatter can produce three orders of magnitude more energy per gram than fission or fusion and ten orders of magnitude more energy than the chemical reactions currently used for propulsion. As a result, it is a prime candidate for use in future exploration beyond the solar system. It also is a candidate for future missions to the edge of the solar system, which now require on the order of 15-20 years after launch just to reach Pluto.  In this talk (MMN) we start with a quick review of both the discovery of and also our understanding of antimatter (Section 2).  In Section 3 a description is given of how cold antihydrogen was created in 2002.  We point out why this is the only form of antimatter which is practical for deep space propulsion.  We then take a side trip into current studies using antiprotons for cancer therapy (Section 4). This side trip is important as medical research may help with the funding necessary to yield large amounts of antimatter.  In Section 5 we outline the trail we need to break to obtain the dense antihydrogen that would be needed for deep space travel.  We go on in Section 6 with a discussion of what we can do now to start on this path, providing a roadmap towards the goal.   Our conclusions follow. ",
        "conclusions": "The road we have described is challenging both scientifically and technologically.   Enormous scientific and technological barriers have to be overcome.  But the potential intellectual and societal rewards, even along the way, are enormous.  Antimatter-matter annihilation is one of the prime candidates to achieve the high specific impulse  i) desired for the challenging missions of exploring the Heliopause and  visiting the Oort Cloud, and ii) needed if we plan to attempt a rendezvous with the nearest star systems. While no clear pathway to the necessary technologies exists, experimental development in the normal matter world of laboratory-sized research equipment can help us to  reach these most ambitious goals, {\\it IF} we simultaneously embark on constructing a dedicated low-energy antiproton facility  It behooves us to now embark on extensive, serious work on the possibilities that are before us. To achieve them quickly it is necessary to set ourselves in motion now."
    },
    "0410/astro-ph0410382_arXiv.txt": {
        "abstract": "Using a three dimensional relativistic particle--in--cell code we have performed numerical experiments of plasma shells colliding at relativistic velocities. Such scenarios are found in many astrophysical objects e.g. the relativistic outflow from gamma ray bursts, active galactic nuclei jets and supernova remnants. We show how a Weibel like two--stream instability is capable of generating small--scale magnetic filaments with strength up to percents of equipartition. Such field topology is ideal for the generation of jitter radiation as opposed to synchrotron radiation. We also explain how the field generating mechanism involves acceleration of electrons to power law distributions ($N(\\gamma)\\propto{}\\gamma^{-p}$) through a non--Fermi acceleration mechanism. The results add to our understanding of collisionless shocks. ",
        "introduction": "Many astrophysical objects emit non--thermal radiation when expelled plasma interacts with the surrounding media. These objects include gamma ray bursts (GRBs) and their afterglows, the jets from active galactic nuclei, jets from quasars and supernova remnants. The non--thermal radiation is believed to be emitted in strong collisionless shocks. Despite the vast prevalence and wide astrophysical applicability, collisionless shocks remain poorly understood. The observed non--thermal emission suggests that strong particle acceleration and magnetic field generation takes place in these shocks. It was suggested by \\cite{1999ApJ...526..697M} that a Weibel--like two--stream instability is able to generate a strong magnetic field in the shock transition region. Recently, due to the increase in computer power, this has been verified with particle--in--cell (PIC) simulations (\\cite{bib:astro-ph/0303360,2004ApJ...608L..13F,bib:astro-ph/0408558,bib:medv,2003ApJ...595..555N,bib:astro-ph/0409702,2003ApJ...596L.121S}). Still, todays computer power does not allow us to fully resolve the shock transition region. We can, however, explain the non-linear kinetic plasma dynamics that generate the electromagnetic fields needed for transmission of momentum between the colliding plasma populations. Here, we report on 3D PIC simulations of the shock formation in the counter--streaming region of two colliding plasma shells. ",
        "conclusions": "We have performed numerical experiments of relativistic collisionless plasma shocks using a self-consistent three dimensional particle--in--cell code. We find that the Weibel instability is capable of generating turbulent magnetic fields with a strength up to percents of equipartition. The magnetic field is induced around ion current filaments. These filaments also accelerates electrons to power law distributions. The suggested acceleration scenario does not rule out ion Fermi acceleration but might overcome some of the problems pointed out by \\cite{2004ApJ...613..460B} regarding the apparent contradiction between standard Fermi acceleration of electrons and spectral observations of GRBs.  The microphysics in the field generation and particle acceleration described here is clearly beyond the reach of the magneto hydrodynamic approximation. A parameter study utilizing a PIC code working from first principles is necessary to fully understand the interdependence between the relative bulk Lorentz factors of the colliding plasma shells, the power law index of the non--thermal electron population, $\\epsilon_B$ and in a broader sense the detailed evolution and structure in collisionless shocks."
    },
    "0410/astro-ph0410661_arXiv.txt": {
        "abstract": "We investigate inflation in Randall-Sundrum type II brane scenario with closed Friedman-Robertson-Walker (FRW) brane. We consider only power-law potentials of the scalar field and wide range of powers and parameters for them. For our models we numerically calculate the total number of e-folds, the value of potential at the end of inflation and amplitude and spectral index of scalar perturbations at the epoch when the present Hubble scale leaves the horizon. All these values we calculate for different initial conditions and different values of parameters. Then we compare our theoretical predictions with observation data and set constraints on the parameters of our model. ",
        "introduction": "Brane-world scenarios, after their discovering some years ago~(\\cite{RS1,RS2}) quickly became very popular. There are numerous issues in the field of branes and one of the most actual is brane-world inflation~\\cite{chaot_br}. In this paper we consider brane-world Randall-Sundrum type II (RSII)~\\cite{RS2} scenario with power-law potential of the scalar field and study the possibility for successful inflation in such model. Our current method is similar to method we used when studied inflation in closed Friedman-Robertson-Walker (FRW) models~\\cite{my04}. Namely we calculate the total number of e-folds during inflation, the value of potential at the end of inflation and the amplitude and spectral index of scalar perturbations at the epoch when present Hubble scale leaves the horizon. Then we check if the total number of e-folds is larger than 72 and if it does we say the initial data for this model lead to inflation. This value must be larger than the number of e-folds between the moment when present Hubble scale crosses the horizon during inflation and the moment of the end of inflation. This second value is model-dependent and it depends on some features of theory, such as the way of inflation ends etc.~\\cite{Nhor}. It varies from theory to theory in a range from 55 to 75 approximately and using $N_{hor} = 62$ we use in some sense mean value for it. Our study shows that constraints one can set are slightly dependent on this value if it changes in the range from 55 to 75.  The second test is linked with the value of potential at the end of inflation. We calculate it and compare with observation data~\\cite{V-constr} $$ 3 < \\frac{V_{end}^{1/4}}{10^{15}\\mbox{GeV}} < 29.    \\eqno(1) $$ \\noindent For brane case the situation with the energy density at the end of inflation may be more complicated then in the FRW case (see~\\cite{brane_end_infl} for details).  Also we can calculate index of the scalar perturbations spectrum at the epoch when present Hubble scale leaves the horizon during inflation. Since there are observation constraints on scalar spectral index from WMAP~\\cite{WMAP}, ACBAR~\\cite{acbar}, CBI~\\cite{cbi} and other CMB experimens~\\cite{CMB-other} and large-scale structure~\\cite{LSS}: $$ n_S=0.99 \\pm 0.04 \\quad \\mbox{(WMAP only)}  \\eqno(2) $$  $$ n_S = 0.97 \\pm 0.03   \\eqno(3) $$  \\noindent (WMAP+ACBAR+CBI+2dFGRS+$L_{\\alpha}$-forest) we can compare our predictions with these data and set some constraints on our model.  Finally we calculate the amplitude of scalar perturbations at the epoch when present Hubble scale leaves the horizon and compare it with COBE constraint~\\cite{COBE} $$ A_S \\sim 2\\times 10^{-5}. \\eqno(4) $$  These three small tests help us in setting some constraints on the parameters of our model and on parameters of the potential. But in fact we do not focus on one of parameters~-- {\\it dark radiation}, because it actes only in the beginning of the inflation, so in this sence it behaves itself something like curvature. The influence of the initial curvature on the total number of e-folds and background values for FRW case was studied in~\\cite{my04} and we found that this influence is very weak. One can suspect this influence is weak in brane inflation as well.. But the dark radiation, unlike curvature, is one of the parameters which determine the initial value of the total energy density, so it must act some differently.    The structure of the paper as follows. First, we write down the main equations we used in our work and describe our method. Then, in Section III we describe the situation with quadratic potential, constraints on the parameters for this potential, in Section IV~--  the same but for quartic potential, and linear potential is considered in Section V. Finally in Section VI we discuss our results. ",
        "conclusions": "We have investigated three cases for brane inflation with power-law potential of the scalar field. Namely we considered linear, quadratic and quartic potentials and found constraints on the parameters for these three particulair cases. Constraints on the parameters for these three models are summarized in Table 1. And now let us generalize our results.  \\begin{table} \\caption{\\label{tab:table1}In this table we summarized our constraints on brane inflation models with power-law potentials for three different powers.} \\begin{ruledtabular} \\begin{tabular}{lccc} & $q=1$ & $q=2$ & q=4 \\\\ \\hline $\\lg \\sigma$ & $-4.3 ... -10.3$ & $-4.4 ... -11.7$ & $-4.0 ... -9.8$\\\\ \\hline $\\lg \\lambda$ & $-9.3 ... -12.3$ & $-10.8 ... -12.6$ & $\\sim -17.8$\\\\ \\end{tabular} \\end{ruledtabular} \\end{table}  From comparing Figs. 1, 2 and 3 we can learn that for all three cases we have similair slope of the curve $n_S$. Slopes of lines $N_{total} = 72$ and lines correspondent to Eq.(1) are different due to different powers. But the 'best' behavior (from the model's testability point of view) shows constraint from Eq.(4). Namely increasing of the power leads to decreasing of the maximal value of $\\lambda$ which corresponds to increasing power. This means for $q>4$ we always have an area from intersection of $N_{total}$, $n_S$ and $V_{end}$ constraints and this area always is intersected by $A_S$ line. The only constraint on maximal value of the power one can set from the fact that the maximal value of $\\lambda$ need obey Eq.(10).  Let us note that one can also set some constraints on this model but from other point of view. Say, in~\\cite{0307017,0312162,0407543} some constraints made from some other background. But apart from our results they found that even quartic potential is under strong pressure from observation. Let us also note that in Friedmann case our~\\cite{my04} results and results from~\\cite{0306305} are similair~-- both of them lead to $q~<4$ constraint.  And last our example is linear potential. This is very uncommon potential for inflation, but for brane inflation case it works (see also~\\cite{0312162}). But for case $q<1$ there simple may occur that there will be no intersection between final area of $N_{total}$, $n_S$ and $V_{end}$ with $A_S$ curve: $A_S$ curve seems to 'move' to high values of $\\lambda$ with decreasing of power and at some low enough power it can lie at $\\lambda$'s that higher than any $\\lambda$ of the area of final constraints from $N_{total}$, $n_S$ and $V_{end}$."
    },
    "0410/astro-ph0410457_arXiv.txt": {
        "abstract": "We present a re-analysis of the optical spectroscopic data on SS\\,433 from the last quarter-century and demonstrate that these data alone contain systematic and identifiable deviations from the traditional kinematic model for the jets: variations in speed, which agree with our analysis of recent radio data; in precession-cone angle and in phase. We present a simple technique for separating out the jet speed from the angular properties of the jet axis, assuming only that the jets are symmetric. With this technique, the archival optical data reveal that the variations in jet speed and in precession-cone angle are anti-correlated in the sense that when faster jet bolides are ejected the cone opening angle is smaller.  We also find speed oscillations as a function of {\\em orbital} phase. ",
        "introduction": "\\label{sec:intro}  In a recent paper \\citep{Blu04} we presented the deepest yet radio image of SS\\,433, which revealed an historical record over two complete precession periods of the geometry of the jets. Detailed analysis of this image revealed systematic deviations from the standard kinematic model \\citep{Mar84,Eik01}.  Variations in jet speed, lasting for as long as tens of days, were needed to match the detailed structure of each jet.  Remarkably, these variations in speed were equal, matching the two jets simultaneously.  The Doppler residuals to the kinematic model show little variation with the precessional phase of the jets and this observation rules out variations in jet speed {\\em alone} as the source of the residuals \\citep[e.g.\\, ][]{Kat82a,Eik01}.  Very little phase variation is obtained if the pointing angle jitters \\citep{Kat82a} but there is no evidence excluding symmetric speed variations of the magnitude reported in \\cite{Blu04} superposed on pointing jitter, Fig\\,\\ref{fig:varywithphase}. Thus our findings from the radio image led us to re-analyse the archival optical data.  G.\\ Collins II and S.\\ Eikenberry (with kind permission of B.\\ Margon) made available their compiled datasets, published in \\cite{Col00} and used in \\cite{Eik01}.  We use the Collins' compilation (available at http://www-astro.physics.ox.ac.uk/$\\sim$kmb/ss433/) because of its higher quoted precision, but very similar results are obtained from Margon's. ",
        "conclusions": "Archival optical spectroscopic data on SS\\,433 reveal variations in jet speed, in cone opening angle, and in the phase of the precession. These appear in the plane of redshift residuals (Fig\\,\\ref{fig:eikfigfive}) and through the new (symmetry dependent) technique of combining simultaneously observed redshift pairs for the speed-only ($s$) and angular-only ($a$) characteristics (Fig\\,\\ref{fig:resids_v_time}).  The velocity variations $\\sim 0.014\\,c$ are strongly anticorrelated with cone angle $\\theta$, in the sense that when faster bolides are ejected the cone angle is smaller. We also found smaller amplitude sinusoidal oscillations in speed as a function of orbital phase.  If this is due to ejection speed, perhaps the orbit of the binary is eccentric.  If these 13.08-day oscillations are orbital Doppler shifts the orbital velocity is $\\sim 400\\,{\\rm km\\, s}^{-1}$ (twice that inferred by \\cite{Cra81,Fab90}). If this were the case, the mass of the companion to SS\\,433 would be $> 86\\, {\\rm M}_\\odot$ and if the mass fraction were 0.1, then the mass of the companion would be \\gtsim\\ $100\\, {\\rm M}_\\odot$.  Such masses would be hardly consistent with an A-type companion \\citep{Gei02,Cha04}."
    },
    "0410/hep-ph0410280_arXiv.txt": {
        "abstract": "We investigate the dynamics of thermalization and the approach to equilibrium in the classical $\\phi^4$ theory in $3+1$ spacetime dimensions.  The non-equilibrium dynamics is studied by numerically solving the equations of motion in a light--cone--like discretization of the model for a broad range of initial conditions and energy densities. A smooth cascade of energy towards the ultraviolet is found to be the basic mechanism of thermalization.  After an initial transient stage, at a time scale of several hundreds inverse masses, the squared magnitude of the field spatial gradient becomes larger than the nonlinear term and there emerges a stage of universal cascade, independent on the details of the initial conditions. As the cascade progresses, the modes with higher wavenumbers, but well behind the forefront of the cascade, exhibit weaker and weaker nonlinearities well described by the Hartree approximation, while the infrared modes retain strong selfinteractions. As a consequence, two timescales for equilibration appear as characteristic of two time-dependent wavenumber regions. For $ k^2 \\gtrsim \\overline{\\phi^2}(t) $, we observe an effective equilibration to a time--dependent power--like spectrum with a time scale in the hundreds of inverse masses; cutoff effects are absent and the Hartree approximation holds for $ k^2 \\gg \\overline{\\phi^2}(t) $.  On the other hand, infrared modes with $ k^2 \\lesssim \\overline{\\phi^2}(t) $ equilibrate only by time scales in the millions of inverse masses when cutoff effect become dominant and complete thermalization is setting in. Accordingly, we observe in the field correlator a relatively large and long--lived wavefunction renormalization of nonperturbative character. There corresponds an effective mass governing the long distance behaviour of the correlator which turns to be significantly smaller than the Hartree mass which is exhibited by the modes with $ k^2 \\gtrsim \\overline{\\phi^2}(t)$.  Virialization and the equation of state start to set much earlier than thermalization. The applicability of these results in quantum field theory for large occupation numbers and small coupling is analyzed. ",
        "introduction": "The understanding of the dynamics of thermalization and relaxation in a field theory is a subject of critical importance both in early cosmology as well as in ultrarelativistic heavy ion collisions.  Pioneering work in this topic was initiated by Fermi, Pasta and Ulam\\cite{FPU} for a chain of coupled oscillators.  Since then, this problem has been studied within a variety of models\\cite{varios} with the goal of answering fundamental questions on ergodicity, equipartition and in general the approach to equilibrium in non-linear theories with a large but finite number of degrees of freedom.  By understanding the dynamics of thermalization we mean to uncover in detail the mechanisms that leads to thermal states starting from arbitrary initial conditions as well as to determine the time scales of these phenomena.  In cosmology the inflationary paradigm \\emph{assumes} that after the inflationary stage a period of particle production and relaxation leads to a state of local thermal equilibrium thus merging inflation with the standard hot big bang cosmology\\cite{kolbook,lidlyt}.  Inflationary scenarios lead to particle production either via parametric amplification of fluctuations of the inflaton (in the case of an oscillating inflaton) or by spinodal instabilities during phase transitions\\cite{cosmo,ours,ours2,Ngran}. In both cases the non-equilibrium dynamics is non-perturbative and results in a large population of soft quanta whose dynamics is nearly classical. The non-equilibrium dynamics of particle production and eventual thermalization are non-perturbative and the resulting fluctuations contribute to the evolution of the scale factor, namely the back-reaction from the fluctuations becomes important in the evolution of the cosmological space-time\\cite{cosmo,ours,ours2,Ngran}.  Both parametric amplification or spinodal decomposition lead to non-perturbative particle production of a band of wave-vectors, typically for soft momenta\\cite{cosmo,ours,ours2,Ngran}. This non-perturbatively large population allows a classical treatment of the non-equilibrium evolution. Thermalization may be described in this classical framework or one can implement the quantum 2PI framework in cosmological spacetimes.  \\medskip  In this article we study the non-equilibrium dynamics going towards thermalization in the classical $\\phi^4$ theory in $3+1$ dimensions. As mentioned above the initial stages of non-equilibrium dynamics either in cosmology or in ultrarelativistic heavy ion collisions is mainly \\emph{classical}.  Classical field theory must be understood with an ultraviolet cutoff, or equivalently an underlying lattice spacing, to avoid the Rayleigh-Jeans catastrophe.  The advantage of studying the classical field theory is that the equation of motion can be solved exactly. Contrary to quantum theory, there is no need of making approximations. Classical field theory is expected to be a good approximation to the quantum theory for large occupation numbers and small coupling as we argue in this paper.  We provide a detailed understanding of the \\emph{main dynamical mechanisms} that lead to thermalization and to study the approach to thermalization by several different observables. We address several questions on the equilibrium and non-equilibrium aspects: i) what are the criteria for thermalization in an interacting theory?, ii) what is the mechanism that leads to thermalization?, iii) what is the dynamics for different observables, how they reach thermal equilibrium and which are the relevant time scales?.  We focus our study on these issues within the context of classical field theory, which is an interesting and timely problem all by itself. It is also \\emph{likely} to describe the initial stages of evolution strongly out of equilibrium in quantum theory for the relevant cases of very large occupation numbers and small coupling.  \\bigskip  The main results of this work are the following:  \\begin{itemize}  \\item We implement a discretization of the $\\varphi^4$ model in any number of space-time dimensions which is very accurate and stable, maintains the relativistic symmetry between space and time and conserves energy exactly (that is to machine accuracy on a computer). The same discretization scheme, restricted to space alone, can be used to study the model at canonical equilibrium through Monte Carlo simulations.  \\item We provide the main physical properties of the classical $\\varphi^4$ model in thermal equilibrium using low temperature perturbative expansions as well as Monte Carlo simulations. These simulations show that the leading order perturbative results have a large domain of validity, further extended by the Hartree approximation, except for infrared modes, which retain a strong nonperturbative character.  \\item We extensively investigate the dynamics of the $ \\varphi^4 $ model for a wide range of (infrared supported) initial conditions and energy densities. After a first stage with relatively important fluctuations whose precise structure depends on the details of the initial conditions, the model evolves towards a {\\bf universal stage} where energy transfer from low to high wavenumbers becomes steady and very effective, resulting in a steady smooth {\\bf ultraviolet cascade}. Namely, the power spectrum of the field $\\phi$ and its canonical momentum $\\pi$ acquire support in a monotonic and slow way over larger and larger values of the wavevectors. This ultraviolet cascade leads to a efficient transfer of energy from $\\pi^2$, $\\phi^2$ and the interaction terms $\\phi^4$ to the spatial gradient $\\left( {\\nabla \\phi}\\right)^2 $ which grows monotonically. Therefore, there is a crossover during this stage from a strongly to a weakly interacting theory since the non-linear term $\\phi^4$ becomes much smaller than the spatial gradients. The average wavenumber $\\bk(t)$ of the modes grows monotonically with time approximately as $ t^{1/3} $. The {\\bf universal stage} is well established by a time $t_0 \\simeq 500$ (in inverse mass units) which does not depend on the lattice spacing and depends weakly on the energy density $E/V$ as long as $E/V$ is not very small. The study of the $\\phi$ and $\\pi = {\\dot \\phi} $ correlators allows us to obtain the power spectrum of $\\phi$ and $\\pi$ which exhibits universal scaling properties during the smooth cascade. We find that virialization sets rather fast at times $t\\gtrsim t_0$.  Parallel to virialization, the equation of state approaches the radiation equation. Namely, the ratio of the pressure divided by the energy density approaches $\\frac13$ for energy densities not too small as the nonlinearities weaken. The ultraviolet cascade continues till the lattice cutoff is reached and therefore the universal properties w.r.t. the discretization method are lost. These cutoff effects become important long before full thermalization of the power spectra sets in, but late enough, for small enough lattice spacing, for the window of the universal cascade to be clearly visible.  \\item The infrared modes with $ k^2 \\lesssim {\\overline{\\phi^2}}(t) $ exhibit a much slower dynamics than the rest of the modes. The modes with $\\bk(t)\\gg k^2 \\gg \\phi^2 $ are well described by the Hartree approximation and equilibrate effectively by times $ \\sim t_0 \\simeq 500 $ (in inverse mass units) while the infrared modes reach an equilibrium state only by times $t \\sim 10^6$, when $\\bk(t)$ is very close to the UV cutoff. In this equilibrium state a rather large wavefunction renormalization remains for low wavenumbers. In configuration space the correlation length turns to be $ 1/M_{\\rm eff}(t) $ where the effective mass of the infrared modes $ M_{\\rm eff}(t) $ is substantially smaller than the Hartree mass.  Therefore, there are {\\bf two} scales for equilibration: the shorter one $ t_0 $ characterizes the UV cascade evolution, the local magnitudes as the time average $ {\\overline{\\phi^2}}(t) $ and the modes with $ k^2 \\gtrsim {\\overline{\\phi^2}}(t) $, while the longer scale governs the evolution of the infrared modes and the wavefunction renormalization for $ k^2 \\lesssim {\\overline{\\phi^2}}(t)$.  \\item The thermalization process in $3+1$ dimensions is quite different from that in $1+1$, studied with the same light--cone--like discretization approach in ref.\\cite{uno}, although a universal smooth ultraviolet cascade is present in both cases. In the $1+1$ case the cascade of the $\\pi$ power spectrum is characterized by a single universal shape function with the time evolution reducing to a scale transformation on such function. On the contrary, in $D=3+1$, one can still define a shape function for the cascade but it turns out to be time--dependent; in particular, there exist a window at intermediate values of the scaled wavenumber $k/\\bk(t)$ with a power--like behaviour, but the power depends markedly on time even when cutoff effects are fully negligible. Most remarkably, the shape function is almost flat near the origin in $1+1$ dimensions, allowing to consistently define an effective thermalization which, starting from the deep infrared where power is concentrated at zero time, progresses to higher wavenumbers with a well defined effective temperature monotonically decreasing in time. In $D=3+1$ instead, the deep infrared is the last to thermalize and the mode equilibration in the bulk of the cascade, where power--like spectra are observed, cannot be characterized only by a time--dependent effective temperature.  Nonetheless, the notion of a time--dependent effective temperature keeps its sense for specific (simple) observables also in $ 3+1 $ dimensions. For instance, from the late time behaviour of $ \\pi^2 $, $ \\phi^2, \\; \\phi^4 $ and $ \\left( {\\nabla \\phi}\\right)^2 $ one can read off an effective temperature which monotonically decreases with time approximately as $ \\sim t^{-\\frac13} $. This effective temperature reaches the proper nonzero limit in the lattice for very late times showing that true thermalization has been achieved. In the continuum this effective temperature would always vanish for infinite time.  \\item Although all our present work (and in ref.\\cite{uno}) deals with classical field theory, some of the main features as the ultraviolet cascade and the slow thermalization dynamics of the infrared modes should be relevant in quantum field theory for large occupation numbers and small coupling.  \\end{itemize}  The main message to early cosmology and ultrarelativistic heavy ion collision physics from the present work is that thermalization proceeds very slowly in classical field theory (at least for unbroken symmetry in scalar models).  However, a state of effective equilibration is soon reached and is characterized by a steady, smooth and relatively simple cascade of energy flowing towards the ultraviolet. The natural interpretation is that the ``fast'' degrees of freedom have indeed reached some sort of statistical equilibrium defined by few macroscopic parameters (such as the temperature, but not the temperature alone) which are slowly varying in time. On the continuum, without any cutoff effect for arbitrary long time, this effective temperature should become the only relevant slow--varying quantity for late enough time; it would eventually vanish for infinite time in classical theory, while in quantum theory it should reach a finite nonzero value necessarily involving $\\hbar$.  Thermalization has been reached in quantum field theories by numerical studies in $3+1$ and $2+1$ dimensions in refs. \\cite{ber,berR} and \\cite{gre}, respectively.  In these works the 2PI expansion to lowest order and the analogous Kadanoff-Baym approximation are used, respectively. The couplings considered are quite strong and do not allow a direct connection with the classical dynamics investigated in the present paper.  The numerical studies previously reported on the dynamics of classical and quantum field theories \\cite{marcelo}-\\cite{ast} have not yet focused on studying the mechanism of energy transfer from long to short wavelengths as a function of time.  In ref.\\cite{marcelo} the emergence of spatio-temporal structures is reported for a symmetry breaking $ \\phi^4 $ model. Refs. \\cite{ABW} study the evolution of the classical $1+1$-dimensional $\\phi^4$ model. For further work in this domain see refs.\\cite{A,B} and \\cite{lisref}.  It would indeed be interesting to study if the universal cascade found in the classical theory\\cite{uno} remains at least during some early and intermediate stages in the quantum theory.  Ref. \\cite{mic} investigates the classical in the $3+1$-dimensional massless $ \\phi^4 $ model. The ultraviolet cascade is present in their numerical results.  In addition, an analytic scenario for the scaling behaviour is proposed following the transport treatment of weak wave turbulence \\cite{zakh}. However, this heuristic treatment does not fully capture the thermalization dynamics found in $1 + 1$ dimensions in \\cite{uno} as well as in $3+1$ dimensions in the present paper.  \\bigskip  This paper is organized as follows: in sec. II we present the classical $\\phi^4$ model in the continuum and in the lattice and our way to perform suitable coarse-grainings. Sec. III discuss the classical $\\phi^4$ model in thermodynamic equilibrium and its properties both from perturbative as well as Monte Carlo calculations.  Sec. IV is the core of the article where the dynamics of thermalization is presented. We study the time evolution of local observables as well as the correlators whose Fourier transforms yield the power spectra of the field $\\phi$ and its conjugate momentum $\\pi$. The early virialization and the late thermalization of the infrared modes is discussed.  Sec. V contains discussions and conclusions while four appendices deal with more specific topics. ",
        "conclusions": "The present work shows that a UV cascade enjoying universal properties is the basic mechanism of thermalization both in $3+1$ and $1+1$ dimensions\\cite{uno}. Although we are not in a position to derive the properties of this cascade from the microscopic equation of motion (\\ref{eqnofmot}), it is easy to see from eq.(\\ref{eqnofmot}) that energy should flow towards higher wavenumbers. Assume we start from a plane wave initial condition $$ \\phi(0,\\vx) = A \\cos(\\vk \\cdot \\vx) \\quad , \\quad {\\dot \\phi}(0,\\vx) =  A \\sin(\\vk\\cdot\\vx)\\; , $$ where $A$ is a constant. The nonlinear term $ \\phi^3 $ in Eq. (\\ref{eqnofmot}) will immediately generate higher harmonics: $ \\cos(3 \\; \\vk\\cdot\\vx) , \\; \\cos(9 \\;\\vk\\cdot\\vx),  \\; \\cos(5 \\;\\vk\\cdot\\vx) $ etc. That is, energy is moving to higher wavenumbers.  Moreover, if we consider a superposition of plane waves as initial condition, $$ \\phi(0,\\vx) = \\sum_{a=1}^n A_a \\cos(\\vk_a\\cdot\\vx) \\quad , \\quad {\\dot \\phi}(0,\\vx) = \\sum_{a=1}^n A_a \\sin(\\vk_a\\cdot\\vx)\\; , $$ the interaction term  $ \\phi^3 $ generates sums and differences: $ \\vk_a \\pm \\vk_b\\pm \\vk_c $ etc, for $ a,b,c = 1, \\ldots , n $. This implies energy flowing {\\bf both} for increasing {\\bf and} for decreasing wavenumbers.  Notice that this mode mixing takes place very fast, at the microscopic time scale. That is, a unit time scale in dimensionless variables. The UV cascade we observe happens in a much longer time scale. This means that these fast microscopic processes combine in a nontrivial way resulting on a  UV cascade with a slow time scale.  \\bigskip  It is important to estimate in which extent the present results can be applied in quantum theory. A necessary condition for  the validity of the classical approximation is that the occupation numbers must be large. The relevant occupation numbers decrease during the UV cascade due to the modes flow towards unoccupied high wavenumber slots. Hence, if the classical approximation is valid for late times, it will be valid earlier (as it is the case $1+1$ dimensions \\cite{uno}). A simple criterion for the validity of the classical approximation in QFT at thermal equilibrium is that \\begin{equation}\\label{condQ} T_p \\gg \\omega_p(k_p) \\end{equation} where $\\omega_p(k_p)= \\sqrt{k_p^2+ M_{p, eff}^2}$ are the dimension-full frequencies.  In dimensionless variables eq.\\eqref{condQ} takes the form, \\begin{equation}\\label{conQd} \\frac{T}{\\lambda} \\gg \\sqrt{k^2+ M_{\\rm eff}^2} >  M_{\\rm eff} \\end{equation} Following eq.(\\ref{mef2}) we obtain as an estimate for $M_{\\rm eff}^2$ from our numerical results \\be \\label{meff} M_{\\rm eff}^2 \\sim 1 + A \\; \\sqrt{\\rho} \\ee where $ \\rho = \\frac{E}{V} $ and $ A \\sim 0.2 - 0.5 $.  For $ \\rho \\gtrsim 1 $ we thus find that the classical approximation applies for \\be\\label{cond1} \\lambda \\ll  (2a)^3 \\;  \\rho^{\\frac34} \\ee where we used that $ T = (2a)^3 \\;  \\rho $ for small $a$ [eq.(\\ref{TV})].  For $ \\rho \\lesssim 1 $ we instead find that the classical approximation applies for \\be\\label{lamax} \\lambda \\ll  (2a)^3 \\;  \\rho \\ee Both eq.(\\ref{cond1}) and (\\ref{lamax})  leads to small values of the coupling $\\lambda$ since the spacing $ a $ must be itself small to avoid lattice effects.  However, in inflationary theories very small values of $\\lambda \\sim 10^{-12} $ are customary (in order to agree with the smallness of the CMB anisotropy) which leaves room for the use of the classical approximation.  \\bigskip  As stated above the validity of the classical approximation decreases with time. Let us estimate the time where it ceases to be valid.  The total energy density can be estimated in terms of the average occupation number $  \\overline{ n_k } $ (where the average is here over the $\\vec k$ modes) as follows $$ \\rho \\sim  \\overline{k}^3 \\;  \\overline{ n_k } \\; \\sqrt{\\overline{k}^2 + M_{\\rm eff}^2 }  \\; . $$ The classical approximation holds if the occupation numbers are large for the relevant modes. Therefore, a necessary condition is $$ \\frac{\\overline{ n_k }}{\\lambda}  \\gg 1  \\; , $$ since the coupling $ \\lambda $ has been absorbed in the field redefinition Eq. (\\ref{adim}).  In the ultrarelativistic regime $ \\overline{k} \\gg  M_{\\rm eff} $ we have \\be \\label{ur} \\overline{k}^4 \\ll \\frac{\\rho}{\\lambda} \\quad \\mbox{and~therefore,} \\quad \\overline{k} \\ll \\left(\\frac{\\rho}{\\lambda}\\right)^{\\frac14} \\; , \\ee while in the non-relativistic domain  $ \\overline{ k} \\ll  M_{\\rm eff} $ we have, $$ M_{\\rm eff} \\; \\overline{ k}^3 \\ll \\frac{\\rho}{\\lambda} $$ and hence using eq.(\\ref{meff}) \\be\\label{56} \\overline{k}^3 \\ll \\frac{\\rho^{\\frac34}}{\\lambda} \\quad \\mbox{and~therefore,} \\quad \\overline{k} \\ll \\frac{\\rho^\\frac14}{\\lambda^\\frac13} \\; . \\ee Since $ \\lambda \\ll 1 $, $$ \\frac1{\\lambda^\\frac13} \\gg \\frac1{\\lambda^\\frac14} \\quad , \\quad {\\rm thus} \\quad \\left(\\frac{\\rho}{\\lambda}\\right)^{\\frac14} \\ll \\frac{\\rho^\\frac14}{\\lambda^\\frac13} \\; , $$ and the ultrarelativistic condition eq.(\\ref{ur}) is more stringent than the non-relativistic bound  eq.(\\ref{56}). We can therefore always use the condition eq.(\\ref{ur}) for the validity of the classical approximation.  Using now eq.(\\ref{expk}) for $ \\overline{ k}(t) $ yields in both regimes that the classical approximation is valid for times $t$ \\be \\label{tmax} t\\ll  t_{max} \\simeq  \\frac{\\rho^{\\frac34}}{\\lambda^{\\frac34} \\; k_0^3} \\; . \\ee The conditions eq.(\\ref{lamax}) and (\\ref{ur}) are compatible since $ \\overline{k} < \\frac{\\pi}{2a} $. Therefore, eq.(\\ref{ur}) implies $$ \\frac{\\pi}{2a} \\ll  \\left(\\frac{\\rho}{\\lambda}\\right)^{\\frac14} \\quad \\mbox{and ~ then} \\quad (2a)^3 \\; \\rho^{\\frac34} \\gg \\pi^3 \\; \\lambda^{\\frac34} \\; \\mbox{or}  \\; \\frac{\\rho}{\\lambda} \\gg \\left(\\frac{\\pi}{2a}\\right)^4  \\; , $$ which is more stringent than eq.(\\ref{cond1}).  In summary, the classical approximation is valid for times earlier than  $ t_{max} $ [given by eq.(\\ref{tmax})] provided \\be \\label{conval} \\frac{\\rho}{\\lambda} \\gg \\left(\\frac{\\pi}{2a}\\right)^4  \\; . \\ee That is, {\\bf high} density (large occupation numbers) and/or {\\bf small} coupling. This condition constraints the initial conditions which fix $ \\rho $. Notice that the coupling $\\lambda$ {\\bf does have} an intrinsic meaning in the quantum theory.  A condition for the validity of the classical dynamics in QFT is derived for high-density in ref. \\cite{berR} in the context of the 2PI-1/N approach. For further studies about the validity of the classical approximation in different contexts see \\cite{A,preter,lisref}.  \\bigskip  Let us now comment about the character of the universal stage. The dimensionality of the space plays a crucial r\\^ole in this phenomena. In one space dimension effective thermalization takes place much faster in analogous conditions\\cite{uno}. Here, in $3+1$ dimensions a second and even longer time scale appears characterizing the thermalization of infrared modes with $ k < \\sqrt{\\overline{\\phi^2}}(t) $. The infrared modes keep interacting for far longer times than the higher $k$ modes. In $3+1$ dimensions the infrared modes thermalize the last while in $1+1$ dimensions they thermalize the first\\cite{uno}.  The UV cascade is clearly less  efficient in three space dimensions to fill the higher $k$-modes pumping energy from the low  $k$-modes. Clearly, in one space dimension the phase space is dramatically small making the UV cascade very efficient.  The same phase space effect ($k^2$) in three space dimensions makes the classical statistical mechanics of the $ \\phi^4 $ theory divergent due to the ultraviolet catastrophe.  \\medskip  All this suggest that thermalization is reached in quantum theory too as recent works (including memory effects) indicate \\cite{ber,berR,gre}. Moreover, for large initial occupation numbers and small coupling the classical regime should correctly describe the evolution of the theory till a time $ t_{max} $ [see eq.(\\ref{tmax})].  \\bigskip  A remarkable feature of the thermalization mechanism is that even starting from a completely classical regime (large occupation numbers at low wavenumbers), the UV cascade depletes these modes and fills the high ones. Therefore, at some time $ \\sim t_{max} $ (that could be very long) quantum physics unavoidable shows up. This is the out of equilibrium counterpart to the fact that classical statistical mechanics is ill defined due to the ultraviolet catastrophe which can only be cured by the quantum treatment.  Thermalization implies to forget everything about the initial conditions except (obviously) the conserved quantities like energy, momentum, angular momentum. The thermalization is therefore expected to substantially change for integrable theories where the number of conserved quantities equals the number of degrees of freedom.  \\bigskip  In this paper we solve the exact microscopic dynamics and then we averaged on time intervals and space in order to derive the slow dynamics we are interested in. This is perfectly correct but a lot of information is first obtained and then dumped in the averaging process. Alternatively a transport approach could be followed deriving equations for the observables in macroscopic time scales and then solve them. We mean to derive transport equations of the Boltzmann or Fokker-Planck type.  \\bigskip  {\\bf Acknowledgments:} Its a pleasure to thank D. Boyanovsky for useful discussions. The numerical calculations were performed on clusters of personal computers at LPTHE-Paris and at the Physics Department of Milano-Bicocca; the simulations with largest lattices were performed on the Avogadro cluster of CILEA.  \\appendix"
    },
    "0410/astro-ph0410310_arXiv.txt": {
        "abstract": "The D/H ratio in interstellar gas varies on scales of a few hundred pc in the local Milky Way, with D/H values ranging from $\\sim$7$\\,$ppm to $\\sim$22$\\,$ppm. The reduction in D/H relative to the primordial value of D/H$\\,\\approx26\\,$ppm is usually attributed to ``astration'' -- conversion of D into other elements by nuclear fusion in stars. However, it is shown here that astration has difficulty accounting for the observations because the expected associated variations in O/H are not seen.  The lower D/H values are instead likely due to ``depletion'' of the D onto dust grains. Polycyclic aromatic hydrocarbons (PAHs) are a possible repository for the missing D, and it appears possible for gas-grain reactions to achieve extreme deuteration of carbonaceous grain material. Grain destruction will release D from the grains; the gas phase abundance of D should therefore be positively correlated with the gas phase abundances of other elements that exhibit strong depletions, such as Mg, Si, Ti, and Fe, which also will be returned to the gas by grain destruction. ",
        "introduction": "In recent years it has become clear that the gas-phase D/H ratio varies significantly from one sightline to another within a few hundred pc of the Sun (see, e.g., Moos et al.\\ 2002; Steigman 2003; Wood et al.\\ 2004; H\\'ebrard 2005). This has often been taken to result from regional variations in the extent to which nucleosynthesis in stars, followed by return of matter to the interstellar medium via stellar winds and ejecta, has lowered the interstellar D/H ratio from the primordial value established by nucleosynthesis in the Big Bang.  I argue here that the observed local variations in the gas-phase D/H ratio are not due to variable astration. The low values of (D/H)$_{\\rm gas}$ seen on some sightlines are instead attributed to incorporation of D into dust grains. I discuss processes acting to add D to dust grains, and show that these can strongly deuterate PAH material in the interstellar medium. This hypothesis predicts a positive correlation of (D/H)$_{\\rm gas}$ with gas-phase abundances of strongly-depleted metals, and deuteration of carbonaceous interstellar grains captured by the {\\it Stardust} mission. ",
        "conclusions": "\\begin{enumerate} \\item The lack of correlation between (D/H)$_{\\rm gas}$ and (O/H)$_{\\rm gas}$ in Figure \\ref{fig:DvsO} argues strongly against variations in astration as the explanation for the observed variation in (D/H)$_{\\rm gas}$. \\item It appears possible for dust grains to remove D from the gas phase in amounts sufficient to account for the observed variations in (D/H)$_{\\rm gas}$ from one sightline to another. \\item Collisions of D with PAH cations, and collisions of D$^+$ with PAHs, are expected to result in incorporation of D into the PAH.  The rate of such collisions is sufficient to deplete D from the gas on time scales of $\\sim$2~Myr in cool H~I clouds. \\item Grain destruction should return D to the gas along with other elements found in refractory grains. There should therefore be an observable correlation between (D/H)$_{\\rm gas}$ and the gas-phase abundances of Mg, Si, Ca, Ti, Fe, and other elements that exhibit strong depletions. This prediction appears to be confirmed by a recent study finding a positive correlation between (D/H)$_{\\rm gas}$ and (Fe/H)$_{\\rm gas}$ (Linsky et al.\\ 2005). \\item Carbonaceous interstellar material collected by {\\it Stardust}, if representative of grains from the local interstellar medium, should be strongly deuterated. \\item Laboratory studies can test whether the proposed deuteration reactions have the required branching ratios. \\end{enumerate}"
    },
    "0410/astro-ph0410583_arXiv.txt": {
        "abstract": "We present a set of evolutionary computations for binaries with a B-type primary at birth. Some liberal computations including loss of mass and angular momentum during binary evolution are added to an extensive grid of conservative calculations. Our computations are compared statistically to the observed distributions of orbital periods and mass ratios of Algols. Conservative Roche Lobe Over Flow (RLOF) reproduces the observed distribution of orbital periods decently but fails to explain the observed mass ratios in the range $\\in$ $\\lbrack$0.4-1$\\rbrack$. In order to obtain a better fit the binaries have to lose a significant amount of matter, without transferring too much angular momentum. ",
        "introduction": "\\citet{Eggleton} introduced the denomination \"liberal\" to make a distinction between binary evolution with mass and angular momentum loss and the conservative case where no mass and consequently no angular momentum leave the system. \\citet{Refsdal et al.} showed that the binary AS Eri is the result of liberal binary evolution; the amounts of mass and angular momentum lost by the system are however uncertain. \\citet{Sarna} showed that only 60 $ \\%$ of the mass lost by the loser of $ \\beta$ Per was captured by the gainer and that 30 $ \\%$ of the initial angular momentum was lost during RLOF. Hence it is clear that the liberal evolutionary scenario is  important for binary evolution calculations. However the amount of mass and angular momentum that has to be removed from interacting systems is far from obvious.  With the Brussels simultaneous evolution code (a description is given in \"The Brightest Binaries\", \\citet{Vanbeveren et al.}) we calculated a representative grid of conservative evolution of binaries with a B type primary at birth. Application of the criterion of  \\citet{Peters} allows then to determine for each of the evolutionary sequences the beginning and ending of various Algol-stages. This criterion states that in the semi-detached system: \\begin{flushleft} $\\bullet$~~~The less massive star fills its Roche lobe (RL)\\\\ $\\bullet$~~~The most massive star does not fill its RL and is still on the main sequence\\\\ $\\bullet$~~~The less massive star is the coolest, the faintest and the largest\\\\ \\end{flushleft}  The grid allows to determine an expected distribution of orbital periods and mass ratios of Algols. These can be compared to the well established observed distributions (\\citet{Budding et al.}).  Considering the examples mentioned above (\\citet{Refsdal et al.}, \\citet{Sarna}) it may be expected that the match between the observations and the conservative results are far from satisfactory. Therefore we add liberal calculations to our conservative library. The detailed evolutionary tracks can be found in \\verb\"http://www.vub.ac.be/astrofys/\" ",
        "conclusions": ""
    },
    "0410/astro-ph0410060_arXiv.txt": {
        "abstract": "Short-duration Gamma-ray bursts (GRBs) ($\\leq 2{\\rm s}$) have remained a mystery due to the lack of afterglow detection until recently. The models to interpret short GRBs invoke distinct progenitor scenarios. Here we present a generic analysis of short GRB afterglows, and calculate the optical lightcurves of short GRBs within the framework of different progenitor models. We show that all these optical afterglows are bright enough to be detected by the Ultraviolet and Optical Telescope (UVOT) on board the {\\em Swift} observatory, and that different models could be distinguished with a well-monitored lightcurve. We also model the afterglow data of the recently discovered short burst GRB 040924. We find that the limited data are consistent with a low medium density environment which is consistent with the pre-concept of the compact-star merger progenitor model, although the models with a collapsar progenitor are not ruled out. ",
        "introduction": "In the past several years great advances have been been made in revealing the nature of Gamma-ray Bursts (GRBs) of relatively long duration, i.e. $T_{90}>2{\\rm s}$ (e.g. M\\'{e}sz\\'{a}ros 2002; Zhang \\& M\\'{e}sz\\'{a}ros 2004 for recent reviews). However, another category of GRBs, i.e. those with short durations (i.e. $T_{90}<2{\\rm s}$), which comprise about 1/3 of the total GRB population, have remained as mysterious as long GRBs were before 1997. This has been mainly due to the lack of afterglow detections for short GRBs, until very recently.  The leading progenitor model for short GRBs invokes merger of two compact objects (e.g. NS-NS merger or BH-NS merger, Eichler et al. 1989; Paczy\\'{n}ski 1991; Narayan et al. 1992;  M\\'{e}sz\\'{a}ros \\& Rees 1992), which has been found suitable to interpret many short GRB properties (Ruffert et al. 1997; Popham et al. 1999; Perna \\& Belczynski 2002; Rosswog et al. 2003;  Aloy et al. 2004). In this scenario, the burst site is expected to have a large offset from the host galaxy due to asymetric kicks during the birth of NSs (Bloom et al. 1999; but see Belczynski et al. 2002), so that the number density of the external medium in the GRB environment is low, typically $\\sim 10^{-2}{\\rm cm^{-3}}$. Alternatively, with the increasing evidence that long GRB progenitors are collapsars, it has been suggested that short GRBs may also be associated with collapsars, with either a less energetic jet (i.e. short emerging model, Zhang, Woosley \\& MacFadyen 2003) or a jet composed of many subjets seen by an off-axis observer looking into one or a few subjet(s) (subjets model, Yamazaki, Ioka \\& Nakamura 2004). If this is the case, the environment around the progenitor should be similar to that of long GRBs, which is either a constant density medium (e.g. Panaitescu \\& Kumar 2002; Yost et al. 2003) with ISM number density $n\\sim 1{\\rm cm^{-3}}$ or a prestellar wind (e.g. Chevalier \\& Li 2000). Other possibilities for the origin of short GRBs have been proposed within the cylindrical jet model (Wang et al. 2004) and the Poynting-flux dominated GRB model (Lyutikov \\& Blandford 2003).  Within the standard afterglow model and adopting a typical compact star merger environment, the forward shock afterglow emission of short GRBs have been calculated by Panaitescu et al. (2001), Perna \\& Belczynski 2002, and Li et al. (2003). Panaitescu et al. (2001) have shown that the afterglows of short GRBs are faint, and they are likely to be most easily detected in the X-ray band. Li et al. (2003) considered possible $e^\\pm$ pair loading and evaluated its possible observational signature. In this work, we present a generic treatment of short GRB optical afterglows which differs from the previous ones by including both the forward and the reverse shock emission, a crucial ingredient for characterizing the early afterglow light curve and the spectrum. The model is applied to various progenitor models and sample lightcurves are calculated which are compared against Swift UVOT sensitivity (\\S2). Lately, a short, soft burst GRB 040924 was located by HETE-2, which led to the discovery of its optical afterglow (Fox \\& Moon 2004). We also apply the model to fit the afterglow data of this burst (\\S3). ",
        "conclusions": "We have modeled the typical optical afterglow lightcurves for short bursts within the context of the leading progenitor models. Both the forward and reverse shock emission components are considered. With typical parameters, the early afterglows should be detectable by the Swift UVOT, and a well-monitored lightcurve can help to identify the progenitors of short bursts.  The optical afterglow data collected so far for the recent bright short burst GRB 040924 can be modeled well with an isotropic fireball expanding into a low density medium with $n\\sim 10^{-2}{\\rm cm^{-3}}$. The wind model is found to be less favored. The resulting  parameters are consistent with the pre-concept of the compact star merger model. Other models such as a collapsar progenitor with low-density environment, however, cannot be ruled out at this stage. In principle, if GRB 040924 came from a collapsar, a lightcurve flattening is expected within weeks resulting from either the supernova component or the central core component for the subjet model. The non-detection of such a feature so far presents a further constraint on the collapsar model.  GRB 040924 is a relatively soft event. It may not be a good representative of the traditional short-hard bursts. Swift will locate more short-hard bursts, and our analysis could be directly utilized to discuss their nature."
    },
    "0410/astro-ph0410256_arXiv.txt": {
        "abstract": "The HI Parkes All-Sky Survey (HIPASS) gives an unprecedented view of the local large-scale structures in the southern sky. I will review the results from the HIPASS Bright Galaxy Catalog (BGC, Koribalski et al. 2004) and the first version of the deep catalog (HICAT, Meyer et al. 2004) with particular emphasis on galaxy structures across the Zone of Avoidance. Some previously hardly noticed galaxy groups stand out quite distinctively in the HI sky distribution and several large-scale structures are seen for the first time. Because HI surveys preferentially detect late-type spiral and dwarf galaxies, their distribution is less clustered than that of optically selected galaxies. HI surveys also significantly enhance the connectivity of large-scale filaments by filling in gaps populated by gas-rich galaxies. ",
        "introduction": "The HIPASS Bright Galaxy Catalog (BGC, Koribalski et al. 2004) contains the 1000 HI-brightest galaxies in the southern sky as obtained from the HI Parkes All-Sky Survey. These were selected by their HI peak flux density ($S_{\\rm peak} \\ga 116$ mJy) measured in the spatially integrated HIPASS spectrum. We note that a subset of $\\sim$500 galaxies is complete in integrated HI flux density ($F_{\\rm HI} \\ga 25$ Jy\\,km\\,s$^{-1}$). The HIPASS BGC is highly reliable (S/N $\\ga$ 9) and contains accurate HI properties of all sources as well as optical identifications were available. Sources without cataloged optical (or infrared) counterparts, i.e. $\\sim$10\\% of the BGC, fall into two categories: (1) gas-rich galaxies in regions of high Galactic extinction or high stellar density, and (2) late-type dwarf and low surface brightness (LSB) galaxies. The latter are generally missed in optical catalogs because they are small in size or rather faint, but they are easily identified given the HI position. For details of the newly cataloged galaxies in the HIPASS BGC, 70\\% of which lie in the Zone of Avoidance ($|b| < 10^{\\circ}$), see Ryan-Weber et al. (2001). We found only one extragalactic HI cloud, HIPASS J0731--69 ($M_{\\rm HI} \\sim 10^9$ M${\\odot}$), located at a projected distance of 180 kpc from the asymmetric spiral galaxy NGC~2442 (Ryder et al. 2001).  \\begin{figure}[!ht] \\plottwo{stg1new.eps}{stg2new.eps} \\caption{Distribution of galaxies in the southern sky, divided into several velocity ranges (indicated at the top right of each panel). The filled black circles and triangles indicate HIPASS BGC sources, while the plus signs and crosses indicate HICAT sources. Within each velocity bin we distinguish galaxies in the lower half (circles and plus signs) and the upper half (triangles and crosses). Galactic latitudes of $b = \\pm10^{\\circ}$ are marked to emphasize the many new structures seen across the (optical) Zone of Avoidance. We use a stereographic projection onto the South Pole; Right Ascension tick marks on the outer circle (declination $\\delta(J2000) = 0^{\\circ}$) are given every hour, starting at the righthand side and increasing counter-clockwise.} \\end{figure} \\setcounter{figure}{0} \\begin{figure}[!ht] \\plottwo{stg3new.eps}{stg4new.eps} \\plottwo{stg5new.eps}{stg6new.eps} \\caption{--- continued.} \\end{figure}  Fig.~1 shows the distribution of all sources in the HIPASS BGC and HICAT, divided into velocity bins of $\\Delta v$ = 1000 km\\,s$^{-1}$. Fig.~2 shows the histogram of local group velocities in both catalogs. The deep HIPASS catalog (HICAT) contains 4315 HI sources, including a large fraction of the BGC. For the first time, we can trace local large-scale structures over the whole southern sky, unobscured by dust or foreground stars. A beautiful network of filaments, loops, groups and voids is revealed, creating a cellular web that gradually changes with redshift. In particular, HIPASS uncovers previously obscured galaxy structures across the Zone of Avoidance. The most striking feature in the first velocity bin ($<$1000 km\\,s$^{-1}$) is the --- in this projection --- nearly horizontal, well defined narrow filament which corresponds to the near-side of the Supergalactic Plane (SGP). It is very long, reaching from the Sculptor Group on the righthand side all the way past the Cen\\,A Group on the left, clearly continuing through the optical Zone of Avoidance (ZOA). The voids on both sides of the plane help to delineate it so clearly. Also visible is the near-side of the Fornax Wall whose full extend is seen in the next velocity bin (1000 -- 2000 km\\,s$^{-1}$), revealing another ZOA crossing. The criss-crossing of filaments (or walls) creates a remarkable cellular structure depicting regions of high galaxy density and voids. The local large-scale structures are further described in Section~3. Overall, we find that the filaments seen in HI appear well defined and continuous while optically outlined structures are generally much more clumpy (Koribalski et al. 2004). In other words, the correlation length of HI rich galaxies, most of which are late-type spiral galaxies (see Fig.~3) is much smaller than that of optically selected galaxies, which includes many more elliptical and early-type galaxies. This decrease of clustering from early- to late-type galaxies (Giuricin et al. 2001) affects our view of the local large-scale structure such that HI surveys reveal many previously unknown structures and enhance connectivity within known structures.  Fig.~4 also shows the distribution of galaxies on the southern sky, but here symbols indicate the galaxy morphological type. While optical identification of the HIPASS BGC sources revealed $\\sim$90\\% cataloged galaxies (see Koribalski et al. 2004), cross-identification for HICAT is still under way. The left side of Fig.~4 shows the optically identified BGC sources with their mean morphological type as obtained from the Lyon Extragalactic Database (LEDA) as well as all newly cataloged sources. A preliminary match of all HICAT sources with velocities below 4000 km\\,s$^{-1}$ was also carried out using LEDA; the result is shown on the right side of Fig.~4. ",
        "conclusions": ""
    },
    "0410/astro-ph0410706_arXiv.txt": {
        "abstract": "We report the discovery of a new metal-rich DAZ white dwarf, GD 362. High signal-to-noise optical spectroscopy reveals the presence of spectral lines from hydrogen as well as Ca~\\textsc{i}, Ca~\\textsc{ii}, Mg~\\textsc{i}, and Fe~\\textsc{i}.  A detailed model atmosphere analysis of this star yields an effective temperature of $\\Te=9740\\pm50$~K, a surface gravity of $\\logg=9.12 \\pm 0.07$, and photospheric abundances of $\\log{\\rm Ca/H}=-5.2\\pm0.1$, $\\log{\\rm Mg/H}=-4.8\\pm0.1$, and $\\log{\\rm Fe/H}=-4.5\\pm0.1$. White dwarf cooling models are used to derive a mass of 1.24 \\msun\\ for GD 362, making it the most massive and metal-rich DAZ uncovered to date. The problems related to the presence of such large metal abundances in a nearby ($d\\sim 25$~pc) white dwarf in terms of an accretion scenario are briefly discussed. ",
        "introduction": "The most common type of white dwarfs are those that show hydrogen absorption lines in their optical spectra, the DA stars. A fraction of these, the so-called DAZ stars, also display absorption features from heavier elements, most notably the Ca~\\textsc{ii} H and K lines. A recent spectroscopic survey at high dispersion of about 120 cool ($\\Te<10,000$~K) DA stars by \\citet{zuckerman03} reveals that $\\sim 25$\\% of the objects in their sample show metallic features. The DAZ phenomenon thus seems much more common than previously believed \\citep{billeres97}.  The origin of these metals in the atmospheres of DAZ stars is still poorly understood, however. Indeed, as most DAZ stars are comparatively cool and old, it is expected that all elements heavier than hydrogen would have settled rather quickly at the bottom of the atmosphere due to the strong gravitational field present in these stars. Therefore, any metals found in the atmosphere must be provided by an exterior source, such as the interstellar medium or even comets. We refer the reader to \\citet{zuckerman03} for a more detailed discussion of the various theories that have been put forth to account for the presence of metals in DAZ stars.  We have recently undertaken a spectroscopic survey of white dwarfs drawn from the catalog of \\citet{mccook99}. Our aim is to derive the atmospheric parameters for all stars, as well as to confirm the spectroscopic classifications given in the catalog using modern CCD spectroscopy. As part of this survey, we have discovered that GD 362 (WD 1729+371, PG) is a unique DAZ white dwarf, with metallic features stronger than in any other member of its class, the Ca~\\textsc{ii} H and K lines in particular. GD 362 had initially been classified as a DA white dwarf by \\citet{greenstein80} with a note that it could possibly constitute a composite spectrum or a dK subdwarf in a common proper motion binary system. It was later rediscovered in the Palomar-Green survey \\citep{green86} where it was classified \\textit{sd}, which does not imply a subdwarf type, but rather ``a lower signal-to-noise observation in which two or three Balmer absorption lines of moderate gravity are visible''  We present in this {\\it Letter} our spectroscopic data for GD 362 in which, besides the usual hydrogen Balmer lines, we have identified spectral lines of Ca~\\textsc{i}, Ca~\\textsc{ii}, Mg~\\textsc{i} and Fe~\\textsc{i}. By fitting the optical spectrum with a grid of synthetic spectra appropriate for these stars, we determine the effective temperature and surface gravity of GD 362, as well as the abundances of all the elements observed in this star. ",
        "conclusions": "Using the evolutionary models of \\citet{fontaine01} for DA white dwarfs, we derive a mass of 1.24 \\msun.  Our derived atmospheric parameters thus make GD 362 the most metal-rich and massive DAZ white dwarf studied thus far, and somewhat hotter than the majority of DAZ stars discovered by \\citet{zuckerman03}. In Figure \\ref{fg:f3}, we show the calcium abundance as a function of effective temperature for the white dwarfs from \\citet{zuckerman03}. We see from its location in Figure \\ref{fg:f3} how unique GD 362 truly is. In fact, if we compare with the values listed in Table 4 of \\citet{zuckerman03}, the iron and magnesium abundances are nearly solar!  A comparison of our calcium abundance determination for GD 362 with those of other DAZ stars from \\citet{zuckerman03} with similar effective temperatures reveals that the abundance observed in GD 362 is at least a factor of $10^{3}$ higher (see Fig.~\\ref{fg:f3}). Since the mass of GD 362 is also higher than average, we need to verify whether this is the reason for the unusual observed photospheric abundances. To do so, we compare the expected steady-state calcium abundance in a massive star at 1.2 \\msun\\ to that of a normal star with a mass of 0.6 \\msun.  The steady-state abundance, $X_{\\rm ss}$, is given by equation (4) of \\citet{dupuis93},  \\begin{equation} X_{\\rm ss}=\\frac{\\theta \\dot{M}}{\\Delta M_{\\rm cz}}\\ \\ , \\end{equation}  \\noindent where $\\theta$ is the diffusion time scale of calcium at the bottom of the hydrogen convection zone, $\\dot{M}$ is the accretion rate of that element, and $\\Delta M_{\\rm cz}$ is the mass of the convection zone. If we assume the Bondi-Hoyle accretion rate, then $\\dot{M}$ is proportional to $M^2$, where $M$ is the mass of the white dwarf.  Then the expected relative abundance of calcium between a 1.2 and a 0.6 \\msun\\ white dwarf can be expressed as  \\begin{equation} \\frac{X_{1.2}}{X_{0.6}} \\propto \\frac{\\theta_{1.2}}{\\theta_{0.6}} \\ \\ \\frac{M_{1.2}^{2}}{M_{0.6}^{2}} \\ \\ \\frac{\\Delta M_{{\\rm cz}, 0.6}}{\\Delta M_{{\\rm cz}, 1.2}}\\ \\ , \\end{equation}  \\noindent where we assume that both stars are moving with the same velocity through an interstellar medium of similar density. Using values of $\\theta$ and $\\Delta M_{\\rm cz}$ generously provided to us by G.~Fontaine (private communication), we obtain that  \\begin{displaymath} \\frac{X_{1.2}}{X_{0.6}} \\propto \\frac{1}{36} \\cdot 4 \\cdot 8.5 \\approx 1\\ \\ . \\end{displaymath}  \\noindent We see that even though a more massive star accretes at a higher rate and has a thinner convection zone --- implying a smaller dilution factor, the diffusion time scale at the bottom of the hydrogen convection zone is so much shorter that comparable abundances are predicted.  Additionally, the accretion model suffers from the fact there is little interstellar matter within 100 pc from the Sun, as \\citet{aanestad93} aptly pointed out. The distance to GD 362 can be estimated in two ways. First, the $\\Te$ and $\\logg$ values inferred for GD 362 together with the photometric calibration of \\citet{bwb95} yield an absolute visual magnitude of $\\mv=14.31$ (or $M_B=14.48$), which combined with the photographic magnitude of $B_{\\rm ph}=16.15$ taken from \\citet{green86}, yields a distance estimate of $\\sim 22$~pc. Alternatively, we can use the relation between the monochromatic fluxes from our optical spectrum, $f_{\\lambda}$, and the theoretical Eddington fluxes, $H_{\\lambda}$, given by  \\begin{equation} f_{\\lambda}=4\\pi (R/D)^{2}H_{\\lambda}\\ \\ , \\end{equation}  \\noindent where $R$ and $D$ are the radius of the star and its distance from Earth, respectively, to obtain an independent distance estimate of $\\sim 26$~pc, consistent with our previous estimate. In either case, GD 362 seems to be well within the local bubble and far from any high concentrations of interstellar matter. We are therefore at odds to explain the very high metal abundances observed.  Another possibility would be that GD 362 is an unresolved degenerate binary composed of a normal DA star and a helium-rich DZ star of comparable luminosities. If this were the case, however, the intrinsic absorption features of each star would be even stronger than observed here, because of the line dilution factor, and the effective temperature inferred from the Balmer lines would thus be grossly {\\it underestimated} (Balmer lines get stronger with increasing $\\Te$ in this temperature range). In such unresolved systems, the slope of the theoretical spectrum is incompatible with that of the observed spectrum \\citep[see, e.g.,][]{bergeron90}. We see on the contrary from Figure \\ref{fg:f2} that the slope of the theoretical spectrum is in perfect agreement with the observed slope of GD 362, consistent with the presence of a single star.  It is perhaps not too suprising that GD 362 had not been classified as a DAZ white dwarf by \\citet{greenstein80} or \\citet{green86}.  The very strong Ca~\\textsc{ii} H and K lines seen on their photographic spectrum, in addition to the narrow Balmer lines, have led these investigators to classify GD 362 as composite or as a lower gravity object. This raises the interesting possibility that many such objects exist and have been overlooked both in the past and in modern surveys such as the Sloan Digital Sky Survey.  We thank Gilles Fontaine and Pierre Brassard for useful discussions. We would also like to thank the director and staff of Steward Observatory for the use of their facilities, as well as the referee for useful suggestions. This work was supported in part by the NSERC Canada and by the FQRNT (Qu\\'ebec).  \\clearpage"
    },
    "0410/astro-ph0410195_arXiv.txt": {
        "abstract": "{ We summarize the Sloan Digital Sky Survey data acquisition and processing steps, and describe {\\it runQA}, a pipeline designed for automated data quality assessment. In particular, we show how the position of the stellar locus in color-color diagrams can be used to estimate the accuracy of photometric zeropoint calibration to better than 0.01 mag in 0.03 deg$^2$ patches. Using this method, we estimate that typical photometric zeropoint calibration errors for SDSS imaging data are not larger than $\\sim0.01$ mag in the $g$, $r$, and $i$ bands, 0.02 mag in the $z$ band, and 0.03 mag in the $u$ band (root-mean-scatter for zeropoint offsets). ",
        "introduction": "Modern large-scale digital sky surveys, such as SDSS, 2MASS, and FIRST, are opening new frontiers in astronomy. Due to large data rates, they require sophisticated data acquisition, processing and distribution systems. The quality of data is of paramount importance for their scientific impact, but its quantitative assessment is a difficult problem -- mainly because it is not known very precisely what measurement values to expect. In the majority of large surveys to date, data quality assessment has not been fully automated and has required substantial human intervention -- a significant shortcoming and resource sink when the data rate and volume are large.  The SDSS (Sloan Digital Sky Survey, York et al.~2000, Abazajian et al.~2003) is an optical imaging and spectroscopic survey, which aims to cover one quarter of the sky, and obtain high-quality spectra for 100,000 quasars, a similar number of stars, and a million galaxies. The SDSS observing, data processing, and data dissemination are highly automated operations -- occasionally, SDSS is referred to as a {\\it science factory} (to date, over a thousand journal papers are based on, or refer to, SDSS). We provide a brief summary of these operations in Section 2.  As has been the case for other surveys with large data rates, the data quality assessment and assurance has been a hard problem for SDSS. Recently, we have developed a set of tools, organized in the {\\it runQA} pipeline, that are used to automatically assess the quality of SDSS imaging data products, and report all instances where it is substandard. The main ideas and results are described in Section 3, and further discussed in Section 4. ",
        "conclusions": "SDSS is an excellent example of a modern astronomical survey -- it produces unprecedentedly accurate data ($\\sim$10 times more accurate than previous large-scale optical sky surveys, such as POSS, see Sesar et al. 2004), with a large peak data rate (20 GB/hr), and is built upon sophisticated data acquisition, processing and distribution systems ($\\sim$million lines of code), that were developed by a large number of collaborators ($\\sim$50). The success of SDSS provides encouragement that even more ambitious surveys, such as LSST (Tyson 2002), which is expected to produce and immediately process 20 TB of data per observing night, may also be successful endeavors.  SDSS also demonstrated, as described in this contribution, that a quantitative and efficient data quality assessment can be designed and implemented even when the ``true answers'' are not known for individual measurements. However, the method described here is not universal -- it applies only to the wavelength range accessed by SDSS (0.3--1 $\\mu$m), and to Galactic latitudes more than $\\sim$15 degree from the galactic plane. Despite these shortcomings, it is a robust automated method that tracks the accuracy of SDSS photometric zeropoint calibration to better than 0.01 mag. We estimated, using this method, that typical photometric zeropoint calibration errors for SDSS imaging data are not larger than $\\sim0.01$ mag in the $g$, $r$, and $i$ bands, 0.02 mag in the $z$ band, and 0.03 mag in the $u$ band.    \\vskip 0.1in"
    },
    "0410/astro-ph0410476_arXiv.txt": {
        "abstract": "The preliminary analysis of the optical data of SN1999cw show that this object has the  photometric and  spectroscopic behavior of a Type Ia supernova similar to SN1991T, reaching an apparent magnitude at maximum $ B_{max}  = 14.30 $ and a  $\\Delta m_{15}(B)=0.94$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410689_arXiv.txt": {
        "abstract": "FUSE spectroscopy has proved that extremely hot hydrogen-deficient post-AGB stars (PG1159 stars) display matter on their surface that usually remains hidden in the region between the H- and He-burning shells of the former AGB star. Hence, the spectral analysis of PG1159 stars allows to study directly the chemistry of this intershell region which is the outcome of complicated burning and mixing processes during AGB evolution. Detailed abundance determinations provide constraints for these processes which are still poorly understood. With FUSE we have discovered high neon and fluorine overabundances.  There is also a significant iron deficiency, which may be caused by  s-process neutron capture transforming iron into heavier elements. ",
        "introduction": "PG1159 stars are hot hydrogen-deficient (pre-) white dwarfs ($T_{\\rm eff}$ between 75\\,000 and 200\\,000\\,K, $\\log g$=5.5--8; Werner 2001). They are probably the outcome of a late He-shell flash, a phenomenon that drives the currently observed fast evolutionary rates of three well-known objects (FG~Sge, Sakurai's object, V605 Aql). Flash-induced envelope mixing produces a H-deficient stellar surface (Herwig et\\,al.\\ 1999). The photospheric composition then essentially reflects that of the region between the H- and He-burning shells in the precursor AGB star. The He-shell flash transforms the star back to an AGB star (``born-again AGB star'') and the subsequent, second post-AGB evolution explains the existence of Wolf-Rayet central stars of planetary nebulae and their successors, the PG1159 stars.  PG1159 stars provide the unique possibility to study the chemistry in the intershell region between the H- and He-burning shells that is created after complicated and still poorly-understood burning and mixing processes during the AGB phase. Usually the intershell material remains hidden within the stellar interior. During the third dredge-up on the AGB, however, intershell material can get mixed into the convective surface layer and appears on the stellar surface, though in rather diluted abundances. Nevertheless, this process defines the role of AGB stars as contributors of nuclearly processed matter to the Galaxy. Our motivation to study PG1159 stars is based on the fact that these objects directly display their intershell matter. However, the quantitative interpretation of the abundance analyses is still premature because evolutionary calculations through a final He-shell flash including a full nuclear network are not available, yet.  Before the advent of FUSE with its outstanding FUV capabilities we have performed spectral analyses of almost all known PG1159 stars based on optical and HST-UV spectra. We determined $T_{\\rm eff}$, $\\log g$, and abundances of the dominant elements, namely He, C, and O. Typical values are He=33\\%, C=50\\%, O=17\\% (by mass). Traces of nitrogen (1\\%) or considerable amounts of residual hydrogen (about 25\\%) were found in a few stars, and in a few cases, optical as well as Chandra X-ray identifications of neon were successful (2\\%).  \\begin{figure} \\epsfxsize=\\textwidth \\epsffile{werner_fig1.ps} \\caption{ Discovery of a neon line (left panels) and a fluorine line (right panels) in the hydrogen-deficient PG1159-type central star PG\\,1520+525 (top panels) and in the hydrogen-normal central star of NGC\\,1360 (bottom panels). The neon and fluorine abundances in the PG1159 star (given as mass fractions in the panels) are strongly enhanced, namely 20 times and 250 times solar, respectively, whereas they are solar in NGC\\,1360. Note the strong Fe\\,VII line (not included in the models) at 1141.4\\AA\\ in NGC\\,1360, which indicates a solar iron abundance (Hoffmann et\\,al.\\ 2005). It is not detectable in the PG1159 star, probably due to a subsolar Fe abundance.} \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410640_arXiv.txt": {
        "abstract": "% We obtained new optical spectra of 13 \\hii\\ regions in WLM with EFOSC2; oxygen abundances are derived for nine \\hii\\ regions. The temperature-sensitive \\othreea\\ emission line was measured in two bright \\hii\\ regions HM~7 and HM~9. The direct oxygen abundances for HM~7 and HM~9 are 12$+$log(O/H) = $7.72 \\pm 0.04$ and $7.91 \\pm 0.04$, respectively. We adopt a mean oxygen abundance of 12$+$log(O/H) = $7.83 \\pm 0.06$. This corresponds to [O/H] = $-0.83$~dex, or 15\\% of the solar value. In \\hii\\ regions where \\othreea\\ was not measured, oxygen abundances derived with bright-line methods are in general agreement with direct values of the oxygen abundance to an accuracy of about 0.2~dex. In general, the present measurements show that the \\hii\\ region oxygen abundances agree with previous values in the literature. The nebular oxygen abundances are marginally consistent with the mean stellar magnesium abundance ([Mg/H] = $-0.62$). However, there is still a 0.62~dex discrepancy in oxygen abundance between the nebular result and the A-type supergiant star WLM~15 ([O/H] = $-0.21$). Non-zero reddening values derived from Balmer line ratios were found in \\hii\\ regions near a second \\hi\\ peak. There may be a connection between the location of the second \\hi\\ peak, regions of higher extinction, and the position of WLM~15 on the eastern side of the galaxy. ",
        "introduction": "Dwarf galaxies are thought to be the building blocks in the assembly of more massive galaxies within the hierarchical picture of structure formation. These galaxies are also very important venues in which questions about cosmology, galaxy evolution, and star formation may be answered. Dwarf irregular galaxies are relatively low-mass, gas-rich, metal-poor, and are presently forming stars as shown by their \\hii\\ regions, whereas low-mass dwarf spheroidal galaxies are gas-poor and no longer host present-day star-forming events. The properties of these galaxies may be similar to those found in the early universe, and dwarf irregulars may possibly be sites out of which damped Lyman-$\\alpha$ absorber systems form at high redshift (e.g., \\citealp{calura03,prochaska03}). An important question which has yet to be fully explained is the relationship between dwarf irregular and dwarf spheroidal galaxies (e.g., \\citealp{ggh03,scm03,evan_ic1613,vanzee04}, and references therein). That streams have been observed within the Galaxy and M~31 (e.g., \\citealp{yanny03,martin04,zucker04a,zucker04b}) has been taken as evidence of ongoing accretion and of representing past merging of dwarfs by the more massive galaxies. However, work presented by \\cite{tolstoy03} and \\cite{venn_dsph} have shown that stars in present-day dwarf spheroidals cannot make up the dominant stellar populations in the halo, bulge, or the thick disk of the Galaxy, although the merging of dwarf galaxies at very early times cannot be ruled out.  The measurements of element abundances provide important clues to understanding the chemical history and evolution of galaxies. In star-forming dwarf galaxies, the analysis of bright nebular emission lines from the spectra of \\hii\\ regions is used to derive abundances of $\\alpha$-elements (i.e., oxygen) in the ionized gas (see e.g., \\citealp{dinerstein90,skillman98,garnett04}). However, a limited number of elements can be studied by comparison to the number of elements found in the absorption spectra of stars. For a more complete picture, additional elements should be included, since various elements arise from different sites and involve different timescales. Oxygen and other $\\alpha$-elements are created in very massive progenitor stars before being returned to the interstellar medium (ISM) on short timescales, when these stars explode as Type II supernovae. Iron is an element produced by explosive nucleosynthesis in Type~I supernovae from low-mass progenitor stars on longer timescales, and is also produced in Type~II supernovae. Because of the varying timescales for stars of different masses, the $\\alpha$ element-to-iron abundance ratio, [$\\alpha$/Fe],\\footnote{ We use the notation: [X/Y] = log(X/Y) $-$ log(X/Y)$_{\\odot}$. } is tied very strongly with the star formation history (e.g., \\citealp{gw91,matteucci03}). Interestingly, [$\\alpha$/Fe] values for three dwarf irregular galaxies are near or at solar, which indicates that stars have been forming at a very low rate and/or the last burst of star formation occurred long ago \\citep{venn01,venn03,kaufer04}. \\cite{it99} claim that O/Fe is elevated in low metallicity blue compact dwarf galaxies ([O/Fe] = $+0.32 \\pm 0.11$). However, their analysis does not account for potential depletion of Fe onto dust grains, and the Fe abundance is only measured in Fe$^{+2}$, requiring very large and uncertain ionization correction factors (ICFs). \\cite{rodriguez03} finds that the adopted ICFs underestimate the total Fe abundance by factors larger than the elevated abundance ratio claimed by \\cite{it99}. Thus, it is prudent to assume that the nebular Fe abundances in these galaxies, and thus the nebular O/Fe ratios, are quite uncertain \\citep{garnett04}. At present, reliable O/Fe ratios will need to be obtained from stellar abundances. While a complete discussion of $\\alpha$/Fe values is beyond the scope of the present work, brief reviews of stellar abundances in external galaxies have recently been presented by \\cite{tv04} and \\cite{venn_ociw}.  High efficiency spectrographs on 8- and 10-m telescopes have made possible the spectroscopic measurements of individual stars in extragalactic systems. In particular, bright blue supergiants have been observed in galaxies at distances of about 1~Mpc. These hot young massive stars allow us to measure simultaneously present-day $\\alpha$- and iron-group elements. The important advantage of these measurements also allow for the direct comparison of stellar $\\alpha$-element abundances with nebular measurements, as massive stars and nebulae are similar in age and have similar formation sites. Oxygen abundances derived from the spectroscopy of blue supergiants have been obtained in nearby dwarf irregular galaxies NGC~6822, WLM, and Sextans~A \\citep{venn01,venn03,kaufer04}.  The relative ease with which spectra of \\hii\\ regions have been obtained in dwarf irregular galaxies has led to establishing: (1) the metallicity-luminosity relation, thought to be representative of a mass-metallicity relation for dwarf irregular galaxies (e.g., \\citealp{skh89,rm95,lee03field}); and (2) the metallicity-gas fraction relation, which represents the relative conversion of gas into stars, and may be strongly affected by the galaxies' surrounding environment (e.g., \\citealp{lee03field,lee03virgo,scm03}). It is assumed that nebular oxygen abundances are representative of the present-day ISM metallicity for an entire dwarf galaxy, where there is often only a single \\hii\\ region present. In fact, spatial inhomogeneities or radial gradients in oxygen abundances have been found to be very small or negligible in nearby dwarf irregular galaxies (e.g., \\citealp{ks96,ks97}), although recent observations have cast uncertainty about the assumption in NGC~6822 and WLM \\citep{venn01,venn03}. Here we will focus on oxygen abundances, and the comparison between stellar and nebular determinations. For the remainder of this paper, we adopt 12$+$log(O/H) = 8.66 as the solar value for the oxygen abundance \\citep{asplund04}.  \\subsection{WLM}  WLM (Wolf-Lundmark-Melotte) is a dwarf irregular galaxy at a distance of 0.95~Mpc \\citep{dolphin00} and is located in the Local Group. The galaxy was discovered by \\cite{wolf10}\\footnote{ On 1909 October 15, Wolf observed the galaxy in two hours with a Waltz reflector at the Heidelberg Observatory atop K\\\"onigstuhl. He submitted a short description of his observations with the title ``\\\"Uber einen gr\\\"osseren Nebelfleck in Cetus'' (On a larger hazy spot in Cetus) to Astronomische Nachrichten on 1909 November 16. }, and independently rediscovered by Lundmark and Melotte \\citep{melotte26}. WLM is relatively isolated, as the nearest neighbor about 175~kpc distant is the recently discovered Cetus dwarf spheroidal galaxy \\citep{whiting99}. Basic properties of the galaxy are listed in Table~\\ref{table_wlm}.  A number of observations are summarized here. \\cite{jl81} identified two planetary nebulae in the galaxy, and \\cite{sc85} identified the brightest blue and red supergiant stars, including over 30 variable stars. Ground-based optical photometry of stars were obtained by \\cite{ferraro89}, and \\cite{mz96,mz97}. The presence of a single globular cluster was established, and \\cite{hodge99} showed that the properties of the globular cluster are similar to those of Galactic globular clusters. In independent \\halpha\\ imaging programs, \\cite{hhg93} detected two small shell-like features, and \\cite{hm95} cataloged and measured \\halpha\\ fluxes for 21 \\hii\\ regions in the galaxy. \\cite{tomita98} presented \\halpha\\ velocity fields for the brightest \\hii\\ regions in WLM, and showed that the southern \\hii\\ ring is expanding at a speed of 20 km~s$^{-1}$ and that the kinetic age of the bubble is 4.5~Myr. Recent studies of the resolved stellar populations with the Hubble Space Telescope ({\\em HST\\/}) have been carried out by \\cite{dolphin00} with the Wide Field Planetary Camera~2 and by \\cite{rejkuba00} with the Space Telescope Imaging Spectrograph. \\cite{dolphin00} found that over half of the stars were formed about 9~Gyr ago, and that a recent burst of star formation has mostly occurred in the central bar of the galaxy. \\cite{rejkuba00} identified the horizontal branch, also confirming the presence of a very old stellar population. In the carbon star survey by \\cite{bd04}, they found that WLM contained the largest fraction of carbon-to-M stars for the dwarf galaxies surveyed, and showed that WLM is an inclined disk galaxy with no evidence of an extended spherical stellar halo. \\cite{tk01} searched for molecular gas in WLM, but only upper limits to the CO intensity and subsequent H$_2$ column densities were determined. Recent 21-cm measurements with the Australia Telescope Compact Array have shown that there are two peaks in the \\hi\\ distribution, and that the measured \\hi\\ rotation curve is typical for a disk \\citep{jackson04}.  The spectroscopy of the brightest \\hii\\ regions were reported by \\cite{stm89} and \\cite{hm95}. The resulting nebular oxygen abundances were found to be 12$+$log(O/H) $\\simeq 7.74$, or [O/H] $\\simeq -0.92$. \\cite{venn03} measured the chemical composition of two A-type supergiant stars in WLM, and showed that the mean stellar magnesium abundance was [Mg/H] = $-0.62$. However, the oxygen abundance in one of the stars was [O/H] = $-0.21$, which is about 0.7 dex or almost five times larger than the nebular abundance. This presents a vexing question: how can the young supergiant be significantly more metal-rich than the surrounding ISM from which the star was born?  The research reported here is part of a program to understand the chemical evolution from the youngest stellar populations in the nearest dwarf irregular galaxies (e.g., \\citealp{venn01,venn03,kaufer04}). The motivations are: (1) to obtain a homogeneous sample of abundance measurements for \\hii\\ regions presently known in WLM; (2) to measure the temperature-sensitive \\othreea\\ emission line, derive direct oxygen abundances, and compare the present set of measurements with those in the literature; and (3) to examine whether the present measurements show any inhomogeneities in oxygen abundances across the galaxy. This is the first of two papers of our study; the measurements and analyses for \\hii\\ regions in NGC~6822 will be discussed in the next paper (H. Lee et al., in preparation). The outline of this paper is as follows. Observations and reductions of the data are presented in Sect.~\\ref{sec_obs}. Measurements and analyses are discussed in Sect.~\\ref{sec_analysis}, and nebular abundances are presented in Sect.~\\ref{sec_abund}. Our results are discussed in Sect.~\\ref{sec_discuss}, and a summary is given in Sect.~\\ref{sec_concl}. ",
        "conclusions": "% \\label{sec_concl}  Optical spectra of 13 \\hii\\ regions were obtained in WLM, and oxygen abundances were derived in nine \\hii\\ regions. \\othreea\\ was measured in bright \\hii\\ regions HM~7 and HM~9. The resulting direct oxygen abundance for HM~7 is in agreement with previously published values. Our $\\simeq 5\\sigma$ detection of \\othreea\\ in HM~9 confirms the lower signal-to-noise measurements reported by \\cite{stm89} and \\cite{hm95}. For the remaining \\hii\\ regions, oxygen abundances derived with bright-line methods are accurate to about 0.2~dex. We adopt for WLM a mean nebular oxygen abundance 12$+$log(O/H) = $7.83 \\pm 0.06$, which corresponds to [O/H] = $-0.83$, or 15\\% of the solar value. The upper end of the range of derived nebular oxygen abundances just agrees with the mean stellar magnesium abundance reported by \\cite{venn03}, but the present mean nebular result is still 0.62~dex lower than the oxygen abundance derived for the A-type supergiant WLM~15. Significant reddening values derived from observed Balmer emission-line ratios were found in \\hii\\ regions on the eastern side of the galaxy near one of the \\hi\\ peaks discovered by \\cite{jackson04}. There may be a relationship between the location of the east \\hi\\ peak, regions of redder color (higher extinction), large reddenings derived from Balmer emission-line ratios in \\hii\\ regions along the eastern side of the galaxy, and the location of WLM~15."
    },
    "0410/astro-ph0410148_arXiv.txt": {
        "abstract": "Space-born gravitational-wave interferometers such as {\\it LISA} will detect the gravitational wave (GW) signal from the inspiral, plunge and ringdown phases of massive black hole binary mergers at cosmological distances.  From the inspiral waves, we will be able to measure the masses of the binaries' members; from the ringdown waves, we will be able to measure the mass of the final merged remnant.  A subset of detected events allow the identification of both the inspiral and the ringdown waveforms in a given source, and thus allow a measurement of the total mass-energy lost to GWs over the coalescence, $M_{\\rm GW}$.  We define ``golden'' binary mergers to be those with measurement errors likely to be small enough for a physically useful determination of $M_{\\rm GW}$.  A detailed sensitivity study, combined with simple black hole population models, suggests that a few golden binary mergers may be detected during a 3-year {\\it LISA} mission lifetime.  Any such mass deficit measurement would constitute a robust and valuable observational test of strong-field relativistic gravity. An extension of this concept to include spin measurements may allow a direct empirical test of the black hole area theorem. ",
        "introduction": "\\label{sec:intro}  The Laser Interferometer Space Antenna ({\\it LISA}), a joint ESA and NASA enterprise, is due for launch around the year 2013.  One of its main science objectives is the direct detection of gravitational waves (GWs) from coalescing massive black holes (MBHs) at cosmological distances\\footnote{{See {\\tt http://lisa.nasa.gov/}}}.  Although the mission concept has not yet been entirely finalized, it is already possible to estimate the accuracy with which {\\it LISA} will measure a variety of observables related to MBH coalescences. For example, several groups have made predictions for the rate of MBH mergers detectable by {\\it LISA} under various assumptions for the underlying cosmological model of BH mass assembly (see, e.g., Haehnelt 1994; Menou, Haiman \\& Narayanan 2001; Wyithe \\& Loeb 2003; Islam, Taylor \\& Silk 2004; Sesana et al. 2004a,b).  These measurements will provide rich information about the binary generating the GW signal.  Hughes (2002) has presented detailed calculations of the exquisite precision with which {\\it LISA} will measure MBH masses in equal-mass mergers, as a function of redshift. Vecchio (2004) has shown that Hughes' estimates are in fact fairly pessimistic: by taking into account spin-induced precessional effects in the binary coalescence, mass measurements can be improved.  It should even possible to extract useful information about the spins of the binary's members.  In this analysis, we focus on a specific class of MBH binaries that we label ``golden'' binaries.  These are binaries for which {\\it LISA} will witness the entire coalescence, from inspiral to plunge and ensuing ringdown.  In these events, the masses of the binary's BH components before merger and of the final merged remnant can each be measured to high precision.  Because GWs carry mass-energy from the system, the system's final mass $M_f$ will be less than the initial mass $M_i$ it has when the binary's members are widely separated.  For golden binaries, the masses are determined so precisely that the total mass deficit due to GW emission, $M_{\\rm GW} \\equiv M_i - M_f$, can potentially be determined observationally.  Such a measurement would constitute a very robust test of a strikingly strong-field prediction of general relativity.  At present, theory does not tell us us too much about $M_{\\rm GW}$ in general, though we have information about some idealized cases, such as extreme mass ratios (e.g., Davis et al 1971, Nakamura, Oohara, \\& Kojima 1987). Much effort is being directed towards improving this situation, both via large scale numerical efforts (see Baumgarte \\& Shapiro 2003 for a review, or Br\\\"ugmann, Tichy, \\& Jansen 2004 for a snapshot of recent progress) and analytical calculations (Damour 2001).  It is thus obviously of great interest to understand what direct observational constraints can be obtained about this process as it occurs in nature.  We show that for the golden binaries {\\it LISA} has the capability to measure mass deficits in cosmic MBH mergers to high enough accuracy that strong-field general relativity can be tested.  A particularly noteworthy feature of this measurement is that our calculation relies only on parameters measured during the ``inspiral'' (when the binary's members are widely separated and slowly spiraling towards one another due to the backreaction of GW emission) and during the ``ringdown'' (the last dynamics of the system, generated after the binary's holes have merged and are settling down to a quiescent Kerr black hole state).  No information about the highly dynamical final plunge and ``merger'' process is needed.  In this sense, measurement of $M_{\\rm GW}$ from golden binaries is very robust, and will thus constitute a particularly simple but powerful probe of strong-field gravity.  We argue that it is not unreasonable to expect {\\it LISA} to perform such measurements for a 3-year mission life-span.  In {\\S\\ref{sec:model}}, we describe the parameterized gravitational waveform upon which we base our analysis.  We describe in some detail the forms we use for the inspiral and ringdown waves, outlining how they depend upon and thus encode the binary's masses.  It's worth noting that the functional form of the inspiral waves that we use does not take into account precessional effects that arise due to spin-orbit and spin-spin interactions.  A substantial improvement in mass determination during the inspiral can be obtained by taking this effect into account; the spins of the binary's members can thereby be measured as well.  This comes, however, with a substantial increase in waveform complexity and thus computational cost (Vecchio 2004).  It is worth emphasizing this point since it is likely that our estimates of how accurately $M_{\\rm GW}$ is measured are somewhat pessimistic due to our reliance on this relatively crude parameterization.  In {\\S\\ref{sec:measure}}, we describe the calculation we perform to estimate how accurately $M_{\\rm GW}$ can be measured.  What is clear from the modeling that has been done to date is that $M_{\\rm GW}$ is a fairly small fraction of the binary's total mass --- at most, $10-20\\%$ of the total binary mass $M$ is radiated away over the entire coalescence.  In order for a determination of $M_{\\rm GW}$ to be meaningful, the error $\\delta M_{\\rm GW}$ must be substantially smaller than $M_{\\rm GW}$ itself.  We (somewhat arbitrarily) define the regime of golden binaries to be those for which $\\delta M_{\\rm GW}/M \\simeq 5\\%$ or less.  We construct the error $\\delta M_{\\rm GW}$ from errors in the mass parameters measured by the inspiral [the ``chirp mass'' and reduced mass, defined in Eq.\\ (\\ref{eq:insp_masses})] and the ringdown (the final mass of the system).  We outline how this error is calculated using a standard maximum-likelihood measurement formalism (Finn 1992).  We then perform a large number of Monte-Carlo simulations to assess how well $M_{\\rm GW}$ can be measured as a function of source masses and redshift.  We find that binaries with total mass $M$ in the (very rough) range ${\\rm several}\\times 10^5\\, M_\\odot \\lesssim M \\lesssim {\\rm a\\ few}\\times 10^6\\,M_\\odot$ have the potential to be golden out to a redshift $z \\sim 3 - 4$.  In {\\S\\ref{sec:rates}}, we calculate event rates for golden binary mergers.  The specific population models we use are based on a merger tree describing dark matter halo evolution in a concordance $\\Lambda$-CDM cosmology.  This scenario suggests that {\\it LISA} should plausibly be able to measure several golden events during its mission lifetime.  We conclude in {\\S\\ref{sec:conclude}} with a discussion of the limitations of and possible extensions to our work. ",
        "conclusions": "\\label{sec:conclude}  Within a certain mass range, {\\it LISA} will be able to measure the GWs that come from all epochs of the binary black hole coalescence process.  By precisely measuring the initial and final masses of the system, {\\it LISA} will determine how much of the system's mass is radiated away.  Such a measurement will constitute an extremely interesting and robust probe of strong-field general relativity.  As we have emphasized, a particularly interesting feature of these ``golden'' binaries is that the mass deficit can be determined without a detailed model of the highly dynamical, strong field plunge and merger epoch.  Although it is almost certain that our theoretical understanding of this epoch will be much improved by the time {\\it LISA} flies, we should not be surprised if substantial uncertainties remain in the {\\it LISA} era.  A simple and robust measure of strong-field physics such as $M_{\\rm GW}$ can play a particularly important role in developing a phenomenology of binary black hole merger physics.  Such an observable may be useful as a kind of global calibration on the process, even if the waveform models for the plunge/merger regime can not be computed with good accuracy.  We also emphasize again that our estimates for the accuracy with which $M_{\\rm GW}$ will be measured are, in all likelihood, somewhat pessimistic.  Especially for the larger masses, much of our error $\\delta M_{\\rm GW}$ is due to a (relatively) poorly determined value of the binary's reduced mass $\\mu$.  Parameter measurement analyses of post-Newtonian waveforms have shown such large uncertainties are typically due to correlations between $\\mu$ and so-called ``spin-orbit'' and ``spin-spin'' parameters --- parameters which depend upon couplings between the binary's orbital angular momentum and black hole spin vectors (Cutler \\& Flanagan 1993, Poisson \\& Will 1995).  Vecchio (2004) has shown that these correlations can be broken by carefully taking into account spin-induced modulations in the orbital motion of the binary and their influence upon the gravitational waveform.  In some cases, the error $\\delta\\mu$ can be reduced by an order of magnitude or better.  The range of golden binaries thus likely extends to somewhat larger masses than we have estimated, improving the likelihood that {\\it LISA} will measure these events.  This effect may also allow binaries to be measured with ``golden'' precision even if our description of {\\it LISA}'s low-frequency noise turns out to be too optimistic.  Taking into account spin-induced precessions not only reduces the error with which $\\mu$ is measured; it also makes possible a determination of the binary's individual black hole spins, ${\\bf S}_1$ and ${\\bf S}_2$.  As we mentioned in {\\S\\ref{sec:model}}, ringdown waves allow us to determine the merged remnant's spin magnitude $|{\\bf S}_f|$ as well as its mass; in fact, spin and mass are measured with essentially the same relative precision (Finn 1992).  This suggests another extremely interesting measurement that might be possible with these binaries: demonstrating that the initial and final configurations of the system satisfy the {\\it area theorem} (Hawking 1971).  Although the mass of the system must decrease because of radiative losses, \\begin{equation} M_i = M_1 + M_2 > M_f\\;, \\end{equation} the area of all event horizons in the system must {\\it increase} during the coalescence: \\begin{equation} A_i = A_1 + A_2 < A_f\\;. \\end{equation} The area of a black hole's event horizon is determined totally by its mass and angular momentum: \\begin{equation} A = 8\\pi M^2\\left(1 + \\sqrt{1 - a^2}\\right)\\;. \\end{equation} (Here, $a = |{\\bf S}|/M^2$ is the dimensionless Kerr parameter.)  This measurement would demonstrate observationally a fundamental law of black hole physics."
    },
    "0410/astro-ph0410462_arXiv.txt": {
        "abstract": "We carry out a general-relativistic global linear stability analysis of the amassed carbon fuel on the surface of an accreting neutron star to determine the conditions under which superbursts occur.  We reproduce the general observational characteristics of superbursts, including burst fluences, recurrence times, and the absence of superbursts on stars with accretion rates $\\dot{M} < 0.1 \\dot{M}_{\\mathrm{Edd}}$, where $\\dot{M}_{\\mathrm{Edd}}$ denotes the Eddington limit. By comparing our results with observations, we are able to set constraints on neutron star parameters such as the stellar radius and neutrino cooling mechanism in the core. Specifically, we find that accreting neutron stars with ordered crusts and highly efficient neutrino emission in their cores (due to direct URCA or pionic reactions, for example) produce extremely energetic ($> 10^{44}$ ergs) superbursts which are inconsistent with observations, in agreement with previous investigations. Also, because of pycnonuclear burning of carbon, they do not have superbursts in the range of accretion rates at which superbursts are actually observed unless the crust is very impure.  Stars with less efficient neutrino emission (due to modified URCA reactions, for example) produce bursts that agree better with observations. Stars with highly inefficient neutrino emission in their cores produce bursts that agree best with observations. Furthermore, we find that neutron stars with large radii ($R \\sim 16$ km) produce very energetic superbursts that conflict with observations, even if the core neutrino emission mechanism is highly inefficient. Superburst characteristics are quite sensitive to several other parameters as well, most notably the composition of the accreted gas, concentration of carbon in the ignition region, and degree of crystallization of the crust. All systems that accrete primarily hydrogen and in which superbursts are observed show evidence of H- and He-burning delayed mixed bursts. We speculate that delayed mixed bursts provide sufficient amounts of carbon fuel for superbursts and are thus a prerequisite for having superbursts.  We compare our global stability analysis to approximate one-zone criteria used by other authors and identify a particular set of approximations that give accurate results for most choices of parameters. ",
        "introduction": "Superbursts are energetic thermonuclear flashes on the surfaces of accreting neutron stars thought to be caused by unstable carbon burning \\citep{WT76,TP78}.  They are similar to their hydrogen- and helium-burning Type I X-ray burst counterparts (which we refer to as ``normal'' bursts) in that they have thermal spectra, fast rises,  and roughly exponential decays (see Strohmayer \\& Bildsten 2004 for a review).  Superbursts distinguish themselves, however, by their much larger fluences and longer recurrence times.  Superbursts are a fairly new observational phenomenon:  The first superburst was discovered in the system 4U 1735-444 by \\citet{CHKV00}.  Since then, eight more superbursts have been discovered in six other sources (see Kuulkers 2004 for a review).  All nine superbursts had observed energies of $\\sim 10^{42}$ ergs and were detected in systems with accretion rates between $10 \\%$ and $30 \\%$ of the Eddington limit.  Recently, several more superburst candidates were observed in GX 17+2, which accretes at a rate near the Eddington limit \\citep{intZCC04,intZetalHEAD04}.  Superburst recurrence times are not well constrained, though three have been observed within 4.7 years from the system 4U 1636-536 \\citep{W01,SM02,Ketal04}.  Previous theoretical studies of superbursts \\citep{BB98,CB01,SB02} have been quite successful at modeling the general properties of superbursts.  These models demonstrate that a large carbon-enriched layer on the surface of an accreting neutron star can indeed produce a thermonuclear burst due to unstable carbon burning if the accretion rate $\\dot{M} \\gtrsim 0.1 \\dot{M}_{\\mathrm{Edd}}$.  The resulting energetics and recurrence times of the thermonuclear flashes are roughly consistent with observations.  However, these models employ an approximate criterion to estimate the column depth at which the unstable carbon ignition will occur.  Specifically, they define ignition to occur at the depth at which the condition $\\mathrm{d} \\epsilon_{\\mathrm{nuc}} / \\mathrm{d} T > \\mathrm{d} \\epsilon_{\\mathrm{cool}} / \\mathrm{d} T$ is satisfied, where $\\epsilon_{\\mathrm{nuc}}$ is the nuclear energy generation rate and $\\epsilon_{\\mathrm{cool}}$ is an approximation to the cooling rate.  Furthermore, these models integrate the stellar structure and time evolution equations only down to the superburst ignition region, applying  physically-motivated but not necessarily self-consistent boundary conditions at both ends of their computational domains.  However, the long characteristic accretion timescale for superbursts means that the thermal diffusion depth is quite large, often down to the core.  Thus, to accurately model superbursts, one must solve for the thermal and hydrostatic profiles of the entire crust, well below the ignition region.  Recently, \\citet{B04} constructed an improved model that integrates the equations all the way to the core and solves for the inner temperature boundary condition using a method that is more self-consistent than previous studies.  However, he uses the same approximate ignition criterion described above and assigns a fixed outer temperature boundary condition.  Some authors have attempted to go beyond these simple approximations in modeling normal Type I X-ray bursts. \\citet{FL87} were the first to discuss global perturbations of the steady-state configuration of the accreted layer.  Their analysis was rather crude, however, since it assumed a constant temperature perturbation as a function of depth and did not treat perturbations below the accreted layer. \\citet[][hereafter Paper I]{NH03} developed this approach much further by carrying out a full global linear stability analysis to determine the ignition conditions of normal Type I X-ray bursts, thereby putting thermonuclear burst theory on a more rigorous footing.  However, their code has several limitations that, while relatively inconsequential for the study of normal bursts, are inadequate for the application to superbursts.  Explicitly, the model omits carbon burning, is not fully general-relativistic, and is unable to integrate the stellar structure and time evolution equations below the depth at which neutron drip occurs.  In the present investigation, we remove these limitations by significantly improving and expanding the burst model of Paper I. We then apply the new model to superbursts.  The motivation of this present study is to develop a mathematically more rigorous and self-consistent stability analysis to determine the physical conditions under which superbursts occur.  Our goals are to compare and contrast our results to preceding theoretical work, determine the effects of a wide range of neutron star parameters upon superburst characteristics, and put the theory of superburst ignition on a more rigorous footing to enable quantitative comparisons with observations.  Some of the results of this research were previously reported in \\citet{CN04}.  We begin the paper in \\S2 with a description of our numerical model, highlighting the additions to and improvements of the original model of Paper I.  In \\S3, we discuss the thermal profile of the accreted layer and crust of the neutron star.  In particular, we illustrate the effects on the thermal profile of varying several physical parameters such as the accretion rate, accreted gas composition, accreted layer composition, impurity concentration, degree of ion crystallization in the crust, and neutrino emission mechanism in the core.  The consequences of these parameters upon the resulting superburst characteristics, notably the energetics and recurrence time, are significant, as we describe in \\S4.  In \\S5 we compare our results with observations.  We compare and contrast the results of our global linear stability analysis to those of the one-zone approximation in \\S6, and we conclude with a summary in \\S7. ",
        "conclusions": "In this investigation, we have carried out the first self-consistent global linear stability analysis of the carbon fuel on accreting neutron stars and determined the physical conditions under which superbursts occur.  Our model reproduces the general observed superburst features, including burst energies, recurrence times, and the range of accretion rates at which superbursts occur.  In contrast to normal Type I X-ray bursts, the observational characteristics of superbursts are very sensitive to the thermal profile of the entire crust of the neutron star.  Consequently, superbursts can be useful probes to study neutron star interiors.  Because our theoretical model evaluates the thermal profile self-consistently, we are able to explore a wide range of neutron star parameters and study the effects of each on the resulting superburst energies and recurrence times.  By comparing our results with observations, we are able to set constraints on various neutron star parameters.  We find that accreting neutron stars with highly efficient neutrino emission in their cores produce extremely energetic superbursts which are inconsistent with observations. Such neutron stars also do not have superbursts in the range of accretion rates at which superbursts are observed unless the crust is very impure.  Stars with less efficient neutrino emission produce bursts that agree better with observations, while stars with highly suppressed neutrino emission in their cores, e.g., because of superfluidity, produce bursts that agree best with observations.  If the neutron star crust is disordered, the thermal profile of the crust, in particular the region in which superbursts are triggered, is insensitive to the core temperature.  Therefore, superburst energetics and recurrence times from neutron stars with highly efficient core neutrino cooling mechanisms may still be consistent with observations if the crust is completely disordered.  However, such neutron stars will have superbursts at accretion rates lower than the observed cutoff rate.  The cutoff at $\\dot{M} \\sim 0.1 \\dot{M}_{\\mathrm{Edd}}$ thus requires some other effect, e.g.\\ delayed mixed bursts.  Neutron stars with large radii ($R \\sim 16$ km) produce very energetic superbursts that are inconsistent with observations, even if the core neutrino emission mechanism is highly inefficient.  Better constraints on neutron star parameters could be made with improvements to the theoretical model and with more observational data.  All systems in which superbursts are observed and that accrete predominantly hydrogen have exceptionally high $\\alpha$ values (see \\S5). Delayed mixed bursts, which are normal bursts that burn a significant amount of hydrogen and helium stably to carbon before the instability is triggered (Paper I), occur in systems for which $0.1 \\lesssim \\dot{M}/\\dot{M}_{\\mathrm{Edd} } \\lesssim 0.3$ and produce $\\alpha$ values consistent with those observed in these systems.  Since the delayed mixed burst instability occurs above a layer of carbon that has been already produced by stable burning, it is possible that they leave behind sufficient amounts of unburnt carbon to fuel superbursts. For systems in which delayed mixed bursts do not occur, it is likely that the rp-process burns nearly all the available fuel to elements heavier than carbon, leaving insufficient carbon for superbursts.  We thus speculate that delayed mixed bursts are a necessary prerequisite for the occurrence of superbursts in systems that accrete predominantly hydrogen.  More work is needed to verify this suggestion.  Previous theoretical investigations have used approximate one-zone ignition criteria to determine the conditions under which superburst occur.  We find that the results of these approximate criteria generally agree well with our global stability analysis if the parameters used in the one-zone criteria are consistent with those derived in the equilibrium configuration calculation.  Specifically, calculations in which the carbon mass fraction is kept constant do not reproduce the absence of superbursts at lower accretion rates, and calculations which set $\\mathrm{d} \\ln \\epsilon_{\\mathrm{C}} / \\mathrm{d} \\ln T = 26$ produce inaccurate results for neutron stars with highly efficient neutrino emission mechanisms in their cores. For calculations in which these quantities are consistent with those derived in the equilibrium configuration calculation, the only significant discrepancy we have found is for systems that accrete at a rate close to the critical accretion rate below which superbursts do not occur.  At these accretion rates, a substantial amount of carbon burns stably before a superburst is triggered.  Consequently, the one-zone approximation drastically underestimates the superburst recurrence times."
    },
    "0410/astro-ph0410238_arXiv.txt": {
        "abstract": "Various scenarios have been proposed to explain the origin of the Galactic high-velocity clouds, predicting different distances and implying widely varying properties for the Galaxy's gaseous halo.  To eliminate the difficulties of studying the Galactic halo from within, we have embarked on a program to study anomalous neutral gas in external galaxies, and here we present the results for two nearby, face-on spiral galaxies, M 83 and M 51. Significant amounts of anomalous-velocity \\hi\\ are detected in deep VLA 21-cm observations, including an extended, slowly rotating disk and several discrete \\hi\\ clouds.  Our detection algorithm reaches a limiting \\hi\\ source mass of 7\\eex{5} \\msun, and it allows for detailed statistical analysis of the false detection rate.  We use this to place limits on the HVC mass distributions in these galaxies and the Milky Way; if the HVC populations are similar, then the Galactocentric HVC distances must be less than about 25 kpc. ",
        "introduction": "The high-velocity clouds (HVCs) of the Milky Way galaxy were discovered over 40 years ago by \\cite{Mulleretal63}, and have been extensively studied in \\hi\\ emission, \\halpha\\ emission, and absorption, but their nature remains puzzling.  These clouds of \\hi\\ gas, deviating from the disk velocity by up to 300 \\kps, exhibit a tremendous variation in size, morphology and velocity.  The largest complexes span many degrees on the sky and possess complicated spatial structure, while the compact high-velocity clouds \\citep[CHVCs; ][]{BraunBurton99} are of small angular extent.  The distances to the HVCs remain largely unknown, except for a few with direct determinations \\citep[\\eg, absorption against background halo stars,][]{Wakker01} or indirect determinations \\citep[\\eg,][]{Bland-Hawthornetal98,Putmanetal03}.  A number of formation theories posit either galactic or extragalactic origins for the HVCs, predicting different physical and dynamical properties.  Among the galactic models, the most favored is the galactic fountain, a cyclic phenomenon whereby hot ($T \\sim \\eez{6}$ K) gas is ejected by supernova explosions into the lower halo, where it cools and falls back down as observable \\hi\\ \\citep{ShapiroField76,Bregman80}.  Models which point to an extragalactic origin include stripping from companion galaxies, accretion from the inter-galactic medium \\citep[IGM; \\eg,][]{Oort66,Oort70,Oort81}, or a population of gaseous, dark-matter-dominated mini-halos, either concentrated within $\\sim 200$ kpc of our galaxy or scattered throughout the Local Group \\citep[\\eg,][]{Oort66,Blitz99}.  Since the inferred mass of a 21cm-emitting \\hi\\ cloud goes as the distance squared, the mass distribution of the HVC ensemble is quite different under the various scenarios, as summarized in Table \\ref{tab:scenarios}.  \\begin{table}[t] \\caption{Predictions of HVC Production Scenarios \\label{tab:scenarios}} \\smallskip \\begin{center} \\begin{tabular}{lrrrrr} \\tableline \\noalign{\\smallskip} scenario & distance & $|\\vdev|$ & size & $M_{\\rm cloud}$ & $Z$ \\\\ & (kpc) & (\\kps) & (kpc) & (\\msun) & (\\zsun) \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} IGM infall          & $< 3$        & $< 150$  & $< 1$     & $<$\\eez{4} & 0.1--0.3 \\\\ galactic fountain   & $< 10$       & $< 150$  & $< 1$  & \\eez{4}--\\eez{5} & $\\ge$ 1 \\\\ companion stripping & 5--100       & $< 300$  & $\\sim 1$  & \\eez{5}--\\eez{6} & 0.1--1 \\\\ circumgalactic DM halos & $\\sim 150 $  & $< 300$  & $\\sim 1$ & \\eez{6} & 0.1--0.3 \\\\ Local Group DM halos    & $\\sim 750 $  & $< 300$  & 1--10 & \\eez{8} & 0.1--0.3 \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} \\multicolumn{6}{l}{\\footnotesize Sizes and masses shown are typical values based on the median observed parameters      } \\\\ \\multicolumn{6}{l}{\\footnotesize of the HVC ensemble.  Masses for the first three scenarios assume \\hi\\ is the primary  } \\\\ \\multicolumn{6}{l}{\\footnotesize constituent of the cloud, while the mini-halo scenarios assume an \\hi\\ to dark matter } \\\\ \\multicolumn{6}{l}{\\footnotesize ratio of 0.1.  Velocities for the galactic fountain and IGM infall models assume an } \\\\ \\multicolumn{6}{l}{\\footnotesize adiabatic corona of fully ionized gas with $T \\sim \\eez{6}$ K. } \\\\ \\end{tabular} \\end{center} \\end{table}  The distance and mass distributions provide fundamental demarcations for the formation scenarios, and along with metallicity they can select one or another.  Mass depends on distance, however, and HVC distance is difficult to determine for clouds of primarily neutral hydrogen.  This problem can be largely overcome by observing external galaxies of known distance and of similar type to the Milky Way.  Evidence of anomalous \\hi\\ sources in spiral galaxies has been observed for over a decade \\citep[\\eg][]{vanderHulstSancisi88,Kamphuisetal91,KamphuisBriggs92,KamphuisSancisi93}, and a recent study has revealed for the first time an extensive population of discrete \\hi\\ clouds around an external galaxy \\citep[M31;][]{Thilkeretal04}.  This contribution describes the first results of a study of nearby spiral galaxies to search for HVC counterparts and other anomalous \\hi.  Deep VLA observations of M 83 and M 51 have revealed several discrete \\hi\\ sources. Comparison of the derived source mass distributions to the expected Galactic HVC distribution allows us to discriminate between the various HVC production scenarios.  Full details of the analysis are presented by \\citet{Miller03} and in papers currently in preparation. ",
        "conclusions": "Discovery of a ``beard'' and limits on the HVC mass distribution in M 83 and M 51 constrain the formation models, but as one might expect, multiple scenarios remain viable, and perhaps required.  Under a galactic fountain, the slowly rotating and (presumably) vertically-extended \\hi\\ disk would represent condensing material beginning to fall back on to the star-forming disk.  The lagging rotation is in line with expectations of the model, as material ejected from the disk would flow outward in radius and slow in rotation.  A possible radial inflow discovered by \\citet{Fraternalietal02a} in NGC 2403 supports this condensing argument as well.  In addition, a simple estimate of the mass exchange rate from an assumed cooling time is $\\mdot \\sim 1$ \\msun yr$^{-1}$, a value similar to the star formation rates of the galaxies under study.  The discrete \\hi\\ sources are consistent with a galactic fountain.  The on-disk HVCs, which are real sources of mass $> \\eez{6}$ \\msun\\ in both galaxies, could be condensed, free-falling clouds or cool gas entrained in superbubbles breaking out of the galaxy.  The appearance of holes in the cold \\hi\\ disk at some of these locations \\citep{Miller03} supports the latter idea, unless the infalling clouds are sufficiently dense to punch through the disk.  Tidal stripping clearly is active in M 51, with large tidal tails of \\hi\\ and streams of stars.  Some of the HVC candidates to the west of M 51 might result from this as well, as it is difficult to reconcile these with a galactic fountain.  The extended disk and arms of M 83 could be relics of a recent encounter, of which there is evidence in the form of a faint, arc-like optical companion \\citep{KarachentsevaKarachentsev98}.  The other extragalactic scenarios remain viable, although this study has placed constraints on them.  The anomalous \\hi\\ disk could be accretion from the IGM; metallicity estimates would be invaluable to determine this. The discrete AVCs could be dark matter mini-halos, although we detect only a handful at the $M_{\\rm HI} \\sim \\eez{6}$ \\msun\\ level.  If they are about 50 kpc from the Galaxy,  we would have seen nearly five times as many detections in M 83 and M 51.  Much farther than this, however, and many would have fallen outside of our field of view.  One thing made clear by this meeting is that most galaxies are found to contain ``anomalous'' gas upon deep observation.  Warps, extended disks, and extra-planar clouds are typical of spiral galaxies, and it is the completely quiescent and regular systems that are the anomalies."
    },
    "0410/astro-ph0410524_arXiv.txt": {
        "abstract": "In this short note we discuss recent observation of linear correlation on log-log scale between distribution of dark and visible mass in gravitationally bound systems. The coefficient of such correlation appears to be essentially the same for various systems of dramatically different scales such as spiral galaxies of different luminosities and galaxy clusters. We briefly touch possible interpretations of this observation and implications for the mass of dark matter particle. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410597_arXiv.txt": {
        "abstract": "It has been assumed that intermediate-mass black holes (IMBHs) in globular clusters can only reside in the most centrally concentrated clusters, with a so-called `core-collapsed' density profile.  While this would be a natural guess, it is in fact wrong.  We have followed the evolution of star clusters containing IMBHs with masses between $125 \\le M_{BH} \\le 1000$ $M_{\\odot}$ through detailed $N$-body simulations, and we find that a cluster with an IMBH, in projection, appears to have a relatively large `core' with surface brightness only slightly rising toward the center.  This makes it highly unlikely that any of `core-collapsed' clusters will harbor an IMBH.  On the contrary, the places to look for an IMBH are those clusters that can be fitted well by medium-concentration King models.  The velocity dispersion of the visible stars in a globular cluster with an IMBH is nearly constant well inside the apparent core radius.  For a cluster of mass $M_C$ containing an IMBH of mass $M_{BH}$, the influence of the IMBH becomes significant only at a fraction $2.5 M_{BH}/M_C$ of the half-mass radius, deep within the core, where it will affect only a small number of stars.  In conclusion, observational detection of an IMBH may be possible, but will be challenging. ",
        "introduction": "Over the last few years, four lines of evidence have accumulated pointing to the possible presence of a $\\sim10^3M_\\odot$ black hole in some globular clusters.  The first hint follows from an extrapolation of the $M_{bh}$--$M_{bulge}$ relation found for super-massive black holes in galactic nuclei \\citep{Magorrianetal1998}, which leads to a prediction of a typical central black hole mass $\\sim10^3M_\\odot$ for globular clusters \\citep{KormendyRichstone1995, vanderMarel1999}. The empirical $M_{bh} - M_C$ relation also comes out naturally from rapid mass segregation and the Spitzer instability applied to a standard IMF in young, dense star clusters \\citep{Gurkanetal2004}.  The second hint is related to the discovery of a new class of ultra-luminous, compact X-ray sources (ULXs).  Their high luminosities and strong variability suggest that they are IMBHs, rather than binaries containing a normal stellar-mass black hole, and they may occur preferentially in young star clusters \\citep{Zezasetal2002}.  The third hint stems from an analysis of the central velocity dispersions of specific globular clusters.  \\citet{Gerssenetal2002, Gerssenetal2003} and \\citet{Gebhardtetal2002} have published evidence for black holes in M15 and G1 with masses of order of $10^3$ and $10^4$ $M_\\odot$, respectively (since M31's G1 is an order of magnitude more massive than typical globular clusters in our galaxy, both values fall on the $M_{bh}$--$M_{bulge}$ relation).  The fourth hint is based on detailed $N$-body simulations by \\citet{PortegiesZwartetal2004} of the evolution of a young ($\\sim 10$ Myrs) star cluster in M82, the position of which coincides with an ULX with luminosity $L > 10^{40}$ ergs/sec.  They found that runaway merging of massive stars could have led to the formation of an IMBH of $\\sim10^3M_\\odot$.  Since globular clusters in their youth may have resembled this type of star cluster, it is altogether likely that at least some globulars harbor IMBHs.  None of these four hints in itself carries enough weight to be convincing. Given our lack of understanding of the formation process of globulars, there is no strong reason to expect the $M_{bh}$--$M_{bulge}$ relation to carry over to globular clusters.  ULXs may just be unusual forms of X-ray binaries containing a massive but still stellar-mass black hole.  The velocity dispersion profiles of M15 and G1 can be reproduced by simulations without central black holes \\citep{Baumgardtetal2003a, Baumgardtetal2003b}. Still, the fact that the four arguments are so different in character does suggest that we have to take the possible existence of IMBHs very seriously.  The question arises: which globular clusters contain IMBHs?  The intuitive answer would be: clusters with a steep central luminosity profile, both because a higher density might suggest an easier formation of a large black hole, and because such a black hole could be expected to draw more stars inward.  The main message of this letter is: both of these arguments are wrong. Since the most plausible formation scenario of an IMBH is run-away merging in the first $10^7$ years after the formation of a cluster, dynamical relaxation makes a comparison with current conditions irrelevant, and thus invalidates the first intuitive argument.  More importantly, dynamical $N$-body simulations reported in this paper clearly show that the second intuitive argument is false as well.  We find that IMBHs, whenever they are formed, quickly puff up the core to a size far larger than that of the so-called `core collapsed' clusters. In fact, we show that globular clusters with IMBHs have the appearance of normal King model clusters, except in the central regions. We discuss the observational implications in section 4, after describing our simulation methods in section 2 and our numerical results in section 3. ",
        "conclusions": ""
    },
    "0410/astro-ph0410304_arXiv.txt": {
        "abstract": "Some massive binaries should contain energetic pulsars which inject relativistic leptons from their inner magnetospheres and/or pulsar wind regions. If the binary system is compact enough, then these leptons can initiate inverse Compton (IC) $e^\\pm$ pair cascades in the anisotropic radiation field of a massive star. $\\gamma$-rays can be produced in the IC cascade during its development in a pulsar wind region and above a shock in a massive star wind region where the propagation of leptons is determined by the structure of a magnetic field around the massive star. For a binary system with specific parametres, we calculate phase dependent spectra and fluxes of $\\gamma$-rays escaping as a function of the inclination angle of the system and for different assumptions on injection conditions of the primary leptons (their initial spectra and location of the shock inside the binary). We conclude that the features of $\\gamma$-ray emission from such massive binaries containing energetic pulsars should allow to obtain important information on the acceleration of particles by the pulsars, and on interactions of a compact object with the massive star wind. Predicted $\\gamma$-ray light curves and spectra in the GeV and TeV energy ranges from such binary systems within our Galaxy and Magellanic Clouds should be observed by future AGILE and GLAST satellites and low threshold Cherenkov telescopes such as MAGIC, HESS, VERITAS or CANGAROO III. ",
        "introduction": "It seems obvious that some close massive binary systems should contain young neutron stars able to accelerate leptons and possibly ions to relativistic energies. In fact, numerical simulations of the evolution of neutron stars show that the fraction of systems with not accreting pulsars (during so called ejector phase) may become several percent of the total number of massive binaries (Lipunov~1990). However, due to their proximity to the high density winds of the massive stars most of such systems do not show clear modulation of their radio signal with pulsar rotational period. Therefore, only in the case of broad or highly eccentric binaries the pulsars with periods significantly shorter than 1 second have been discovered (e.g. PSR B1259-63 with the period of 47.8 ms or A0538-66 with the period of 69.2 ms). TeV $\\gamma$-ray emission, modulated with the period of $\\sim 12.59$ ms, has been also claimed from the compact binary Cyg X-3 (Brazier et al.~1990), but this result has not been confirmed by any other experiment. Other binaries are suspected to contain short period pulsars (e.g.  LSI +61 303, Maraschi \\& Treves~1981), or stellar mass black holes ejecting relativistic particles highly anisotropically (e.g. Cyg X-1, Bednarek et al.~1990).  In fact, observations of high energy X-ray and $\\gamma$-ray sources to be coincident with the locations of some binary systems support the hypothesis that the high energy processes play an important role in these sources. For example, EGRET sources have been found in directions towards e.g. LSI +61 303 (2EG J0241+6119, Thompson et al.~1995), Cyg X-3 (2EG J2033-4112, Mori et al.~1997), and LS 5039 (3EG J1824-1514, Paredes et al.~2000). However, no TeV $\\gamma$-rays from the above mentioned sources have been observed up to now (except early claims during the 80's which were not confirmed by subsequent more sensitive observations; see e.g. Weekes 1992). The limits on TeV emission from PSR B1259-63, recently observed by the CANGAROO group at $\\sim$47 days and at $\\sim$157 days after the periastron (Kawachi et al.~2004), are above the theoretical predictions (Kirk et al.~1999, Kawachi et al.~2004, Murata et al.~2003). However, very recently the HESS group reported positive detection of this binary system on the level of $5\\%$ of the Crab emission at energies above $\\sim 400$ GeV with the power law spectrum and spectral index -2.8 (Schlenker et al.~2004). The TeV emission has been observed a few days before and after the periastron passage and declines towards the periastron moment. Observations by the Whipple telescope of other four binary systems containing young pulsars (Hall et al.~2003) have not shown any positive detection of a steady or modulated TeV signal. Another binary system, Cen X-3, reported more recently in the TeV $\\gamma$-rays (Chadwick et al.~1998, 1999, Atoyan et al.~2002), and also in the GeV $\\gamma$-rays (Vestrand et al.~1997), contains slowly rotating and accreting neutron star.  In spite of not completely convincing high energy observational results, massive binaries are still often considered as likely sources of $\\gamma$-rays produced mainly in two general scenarios: (1) the anisotropic injection of particles from neutron stars or black holes, which interact with the radiation or matter inside the binary system (a massive companion, an accretion disk and its corona or a blob, see e.g. Cheng \\& Ruderman~1991, Aharonian \\& Atoyan~1991, Levinson \\& Blandford~1996, Aharonian \\& Atoyan~1999, Romero et al.~2001, Atoyan et al.~2002, Georganopoulos et al.~2002, Romero et al.~2002, Romero et al.~2003, Bosch-Ramon \\& Paredes~2004a,b, Orellana \\& Romero~2004), or (2) the interaction of particles accelerated by a compact object or a shock wave created in collisions of the pulsar and stellar winds (e.g. Vestrand \\& Eichler~1982, Harding \\& Gaisser~1990, Tavani et al.~1994, Tavani \\& Arons~1997, Kirk et al.~1999, Ball \\& Kirk~2000, Murata et al.~2003) or two stellar winds (e.g. Eichler \\& Usov~1993, Benaglia \\& Romero~2003). In fact, the injection of relativistic particles and high energy $\\gamma$-rays in dense radiation field expected close to the accretion disks around compact objects (Carraminana~1992, Bednarek~1993) and massive companion stars (Protheroe \\& Stanev~1987, Moskalenko et al.~1993) should result in copious $\\gamma$-ray production. In very compact binaries and compact objects surrounded by luminous accretion disks, the primary particles can initiate IC $e^\\pm$ pair cascades triggered due to the large optical depths. The $\\gamma$-ray spectra are then produced with the characteristic cut-offs at a few tens of GeV (Bednarek~1997, Bednarek~2000). It was assumed in the latter works that the secondary $e^\\pm$ pairs produced in the cascade process are isotropised inside the binary system by the random component of the magnetic field.  In the present paper we consider in detail a more specific scenario in which the volume of the binary system is separated by the shock wave into two regions with different properties. The shock appears as a result of collisions of the pulsar and massive star winds. We analyze very compact binary systems which contain a young pulsar and a massive companion star of the OB or WR type able to create soft radiation field in which the optical depths for relativistic leptons are much larger than unity. In fact, young pulsars in binary systems can be responsible for quite different phenomena such as the appearance of jets called microquasars (e.g Cyg X-3), binary radio pulsars (e.g. PSR 1259-63, SAX J0635+0533) or a long period accreting objects (e.g. A0535+26). It is very difficult to observe some of such systems due to dense stellar winds produced by early type massive stars. However, binaries hidden dense winds should manifest themselves by the presence of large nonthermal luminosities which can, as we argue below, peak in the $\\gamma$-ray energies. In Sect.~2 we define general conditions inside the massive binary in which the termination shock is created in collisions of the pulsar and stellar winds. The injection of relativistic leptons (electrons or $e^\\pm$ pairs) by the pulsar, their subsequent propagation and interaction with the radiation of the massive star is considered. In Sect.~3 we calculate production of $\\gamma$-rays in IC $e^\\pm$ cascades initiated by these leptons in the pulsar wind zone (PWZ) by applying the Monte Carlo method. Because of partial absorption of these $\\gamma$-rays in the massive star wind region (MSWR), the secondary $e^\\pm$ pairs are created and initiate there the second part of the cascade. This cascade differ significantly from the cascade in the PWZ due to the presence of ordered magnetic fields. The spectral and angular features of $\\gamma$-rays escaping from the PWZ and MSWR are discussed in Sect.~5. As an example we consider the binary system with the parameters derived for the WR star in the Cyg X-3 binary in which there are some observational evidence of the existence of a very young and energetic pulsar with the period of several milliseconds. ",
        "conclusions": "We have analyzed in details the anisotropic IC $e^\\pm$ cascade inside the massive binary system containing young pulsar.  We have considered the situation in which leptons are injected mono-directionally or isotropically from a place (identified with the pulsar) which differs from the center of the isotropic source of soft radiation (identified with the massive stellar companion). The simplest geometrical situation has been analyzed in which a compact object is on a circular orbit around the massive star. As an example, the parameters of the massive binary in the Cyg X-3 system have been applied. We concentrate on the situation in which a pulsar is energetic enough to create strong shock in collision with the wind of the massive star. As a consequence, two regions for the development of the cascade has to be distinguished which differ in basic properties, i.e. the pulsar wind zone (PWZ) and the massive star wind region (MSWR). In the PWZ the cascade develops radially from the pulsar and therefore can be considered only in one direction. In the MSWR the cascade develops in complex magnetic field of the massive star. Therefore, complicated propagation paths of secondary $e^\\pm$ pairs in this magnetic field have to be taken into account. However, for typical values of the magnetic fields in the winds of the massive stars and parameters of Cyg X-3 binary, the energy density of radiation from the massive star is larger than the energy density of the magnetic field. Hence, the IC energy losses of leptons dominate over their synchrotron losses and the influence of synchrotron process on the energy budget of leptons can be neglected in the first approximation.  For this specific binary system we have calculated the optical depths for leptons and $\\gamma$-rays and found that they are much larger than unity (Fig.~\\ref{fig2}). Therefore, if relativistic leptons are injected inside the binary they should initiate efficient cascades. Such cascades are analyzed in the case of leptons injected by the compact object with: (a) the monoenergetic spectrum and energies $10^6$ MeV and, (b) the power law spectrum of the type predicted by the calculations of the cascade processes in the inner pulsar magnetosphere (Hibschman \\& Arons~2001).  The main results of our calculations are following:  \\begin{enumerate}  \\item Most of the energy of primary leptons is transfered to the secondary $\\gamma$-rays ($\\sim 80 - 90\\%$, see Table~1) for both the monoenergetic and the power law primary spectra, and for the shock defined by $\\eta = 0.3$ which correspond to the case of the massive star with the wind parameters from the middle of the estimated range (Sect.~2). This percentage of energy does not depend strongly on the location of the shock for the power law spectrum of primary leptons (compare $\\eta = 0.3$ and 0.06 in Table~1), but is significantly lower for the shock closer to the pulsar, determined by $\\eta =0.06$ in the case of monoenergetic primary leptons. The difference between the power in $\\gamma$-rays arriving to the shock (marked by {\\it shock}) and $\\gamma$-rays escaping to the observer (marked by {\\it esc}) gives part of the power of $\\gamma$-rays produced in the PWZ which is converted into the secondary $e^\\pm$ pairs in the MSWR. These $e^\\pm$ pairs can initiate additional cascade in the MSWR and take typically between one third and half of the $\\gamma$-ray power escaping to the observer from the PWZ (see Table~1). Therefore, the contribution of $\\gamma$-rays produced in the MSWR is usually lower than $\\gamma$-rays produced in the PWZ. However, if the monoenergetic primary leptons are injected into the binary with the shock located relatively close to the pulsar (e.g. for $\\eta = 0.06$) the $\\gamma$-ray luminosity produced in the MSWR can be even a factor of two or three larger than $\\gamma$-ray luminosity from the PWZ. This is due to lower optical depths in relatively small PWZ and so inefficient cascading precess initiated by high energy monoenergetic leptons.  \\item The $\\gamma$-ray light curves for photons escaping from the PWZ at lower energies (energy range $0.1 - 10$ GeV) are anticorrelated with the $\\gamma$-ray light curves at higher energies ($> 10$ GeV) for both discussed primary spectra of leptons independently on localizations of the shock within the binary system (see e.g. the $\\gamma$-ray light curves in Figs.~\\ref{fig4} and \\ref{fig8}). This basic feature is also clearly seen when comparing the $\\gamma$ spectra escaping to the observer at different phases of the pulsar and inclinations of the binary system (Figs.~\\ref{fig5} and \\ref{fig9}). It is easily understood since the intense $\\gamma$-ray fluxes at energy range $0.1-10$ GeV are expected in directions of large optical depths, when the binary system is viewed from the direction close to the massive star. In contrast, high level $\\gamma$-ray fluxes at energies $>10$ GeV are expected when the compact object is in front of the massive star. For considered parameters of the binary system, the optical depth is high enough for efficient production of $\\gamma$-rays above 10 GeV but is too low for efficient cascading and production of lower energy $\\gamma$-rays. Moreover, the $\\gamma$-ray spectra extend to higher energies for the monoenergetic primary leptons but are less intense at lower energies than expected for the power law primary leptons (see Figs.~\\ref{fig5} and \\ref{fig9}) since the optical depths are lower for higher energy particles (see Figs.~\\ref{fig2}). Therefore, it is expected that detailed results in the $\\gamma$-ray spectra from specific binaries (PSR B1259-63, SAXJ 0635+0533), for which the parameters of the massive star are well known, should allow to obtain information on the spectrum of primary leptons accelerated by the pulsar. This problem is difficult to investigate directly in the case of isolated pulsars for which the surrounding soft radiation field in the pulsar wind zone is relatively weak.  \\item The complex magnetic field in the MSWR has significant effect on the directions of propagation of  secondary $e^\\pm$ pairs and, as a consequence, on the distribution of $\\gamma$-rays on the sky. Our detailed calculations show that the angular distribution of $\\gamma$-rays, produced in the cascade processes in the MSWR by secondary $e^\\pm$ pairs from absorption of $\\gamma$-rays from PWZ, strongly depends on the injection direction of primary leptons (see maps in Figs.~\\ref{fig5a} and \\ref{fig5b}). For example, if the primary particles are injected perpendicular to the plane of the binary system then the $\\gamma$-rays produced in the cascade are collimated within the cone around the direction of the local magnetic field lines, i.e. in the outward direction with respect to the massive star (Figs.~\\ref{fig5a}a and \\ref{fig5b}a). This is due to the strong radial component of the massive star magnetic field in the main volume of the binary system (see Eq.~\\ref{eq9}). Therefore, in the case of highly anisotropic injection of primary particles by the pulsar, the strong $\\gamma$-rays fluxes may appear in quite different directions. In contrast, if the primary particles are injected in the directions not far from the the stellar limb, this radial component of the magnetic field tends to focus the $\\gamma$-rays in directions tangent to the stellar limb. A significant part of these $\\gamma$-rays impinge onto the surface of the massive star (see Figs.~\\ref{fig5a}c and \\ref{fig5b}c). It is expected that these $\\gamma$-rays can excite observable nuclear $\\gamma$-ray line fluxes due to their interactions with the matter on the massive star surface. This line emission should be correlated with the high energy $\\gamma$-ray emission ($> 10$ GeV) and anticorrelated with the low energy $\\gamma$-ray emission (between $0.1 - 10$ GeV).  The distribution of $\\gamma$-rays produced in the MSWR is also strongly anisotropic for the case of isotropic injection of primary leptons by the pulsar (see Fig.~\\ref{fig5c}). There is a strong concentration of produced secondary $\\gamma$-rays in directions close to the limb of the massive star. The number of secondary photons produced by the primary leptons with the power law spectrum is by a factor of a few lower than for monoenergetic primary leptons but general emission pattern on the sky does not differ significantly. These features are due to the fact that leptons with lower energies produce $\\gamma$-rays with energies above 100 MeV less efficiently.  \\item The $\\gamma$-ray fluxes produced in the MSWR are weakly dependent on the phase of the pulsar and the inclination angles of the binary system, except for the case of eclipsing binaries (see Figs.~\\ref{fig6} and \\ref{fig10}). This is due to the propagation effects of secondary $e^\\pm$ pairs in the magnetic field of the MSWR which result in a significant isotropisation of the secondary $e^\\pm$ pairs with respect to the directions of propagation of primary leptons. Therefore, when the pulsar is in front of the massive star (phase $\\varphi = 0^o$), the main contribution to the escaping low energy $\\gamma$-ray flux comes from the MSWR and the anticorrelation between the lower energy ($0.1 - 10$ GeV) and higher energy ($>10$ GeV) $\\gamma$-rays produced in the PWZ is partially suppressed (compare Figs.~\\ref{fig4} with \\ref{fig6} and \\ref{fig8} with \\ref{fig10}). The $\\gamma$-ray spectra escaping to the observer from the MSWR are flatter in the GeV energy range (spectral index close to -2) than the $\\gamma$-ray spectra from the PWZ (spectral index close to -1.5) (compare Figs.~\\ref{fig5} with \\ref{fig7} and \\ref{fig9} with \\ref{fig11}). This is due to the difference in the optical depth for particles propagating in the PWZ and the MSWR. Since the cascade is not efficient enough in the PWZ, the secondary $e^\\pm$ pairs are not able to influence lower energy part of the IC $\\gamma$-ray spectrum formed mainly by cooling of primary leptons. However, secondary $e^\\pm$ pairs in the MSWR are cooled to energies low enough to produce a lot of secondary GeV photons and thus creating steeper spectra. Due to the same reasons, the $\\gamma$-ray spectra from the MSWR extend usually to lower energies and the predicted fluxes above $\\sim 10$ GeV are lower.  \\item The $\\gamma$-ray spectra escaping from the MSWR depend strongly also on the localization of the shock inside the binary. This is caused by the change in relative volume for the development of cascades in the PWZ and MSWR. If the shock is closer to the pulsar then the $\\gamma$-ray spectra from the MSWR extend to higher energies (see Figs.~\\ref{fig7} and \\ref{fig11}) since they are produced by on average more energetic secondary $e^\\pm$ pairs originated in the absorption of more energetic $\\gamma$-rays escaping from the PWZ. The $\\gamma$-ray fluxes at energies $< 10$ GeV, produced in the MSWR, do not depend strongly on the localization of the shock. But the $\\gamma$-ray fluxes at energies $>10$ GeV are on a much lower level for the case of primary leptons injected with the power law spectrum (see Figs.~\\ref{fig6} and \\ref{fig10}) from the same reasons as discussed  in this item above.  \\end{enumerate}  In summary, detailed investigation of the phase resolved IC continuous $\\gamma$-ray spectra and the $\\gamma$-ray light curves from the compact massive binaries, in which energetic pulsar is responsible for the acceleration of relativistic leptons, should allow extraction of information on the primary spectra of leptons (monoenergetic, power law ?) and conditions for collisions of the pulsar and stellar winds. The fluxes and spectra of escaping $\\gamma$-rays should be also sensitive to the acceleration site of the primary leptons, inside the inner pulsar magnetosphere (power law spectrum ?), close to the pulsar light cylinder (monoenergetic ?), or during the propagation in the pulsar wind zone (see e.g. re-acceleration model of Contopoulos \\& Kazanas~2002). The effects on the $\\gamma$-ray spectra caused by leptons which are additionally accelerated in the pulsar wind zone have been recently investigated with the application to the PSR B1259-63/Be binary system (Sierpowska \\& Bednarek~2004b), but they need further more detailed calculations.  Taking into account sensitivities of the future and already operating $\\gamma$-ray telescopes in the energy range below a few hundred GeV (which are estimated on $\\sim 3\\times 10^{-13}$ erg cm$^{-2}$ s$^{-1}$ for the GLAST and $\\sim 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ for the MAGIC), the type of the massive binary considered in this paper with the $\\gamma$-ray luminosity of the order of $\\sim 10^{37}$ erg sr$^{-1}$ s$^{-1}$ (similar to the Cyg X-3 massive binary) should be detected at any location within our Galaxy and also within the Magellanic Clouds.  In the forthcoming paper (Sierpowska \\& Bednarek 2004a) we intend to apply such general model for the $\\gamma$-ray production in specific massive binaries in which case there are evidences of the existence of energetic pulsars (e.g. PSR B1259-63, SAX J0635) or compact objects with jets (see the list of sources in Table~1 in Mirabel \\& Rodriguez~1999). Note that in the case of less compact binaries than considered here, the secondary cascade $e^\\pm$ pairs arrive to the shock with significant energies since the optical depths for leptons are much lower. These pairs move then along the shock plane. Their contribution to the total escaping $\\gamma$-ray fluxes has to be also taken into account. It has not been considered in this paper since, for the parameters of the massive star in the Cyg X-3, these pairs take only a relatively small part of energy of primary leptons (typically less than $\\sim 20\\%$ of the total injected energy, see Table~1). Therefore, relative importance of processes occurring in the PWZ, the shock region, and the MSWR can differ significantly for the very close binaries with the pulsars on circular orbits (as considered here) from the binaries of the PSR 1259-63 type, i.e. with a broad, highly eccentric orbit and a massive Be star producing non-spherical stellar wind.   We also plan to investigate possible importance of the cascade processes inside the binary systems of two very massive stars which create strong shock as a result of their stellar wind collisions, e.g. $\\gamma$-2 Vel (WC8+O7.5). Energetic $e^\\pm$ pairs can appear inside such binaries as a result of interaction of hadrons, accelerated by the shock, with radiation field of the massive star."
    },
    "0410/astro-ph0410132_arXiv.txt": {
        "abstract": "We review in this paper the main results recently obtained on the identification and study of very high-z galaxies using lensing clusters as natural gravitational telescopes. We present in detail our pilot survey with ISAAC/VLT, aimed at the detection of z$>7$ sources. Evolutionary synthesis models for extremely metal-poor and PopIII starbursts have been used to derive the observational properties expected for these high-$z$ galaxies, such as expected magnitudes and colors, line fluxes for the main emission lines, etc. These models have allowed to define fairly robust selection criteria to find z$\\sim 7-10$ galaxies based on broad-band near-IR photometry in combination with the traditional Lyman drop-out technique. The first results issued from our photometric and spectroscopic survey are discussed, in particular the preliminary confirmation rate, and the global properties of our high-z candidates, including the latest results on the possible z=10.0 candidate A1835-1916. The search efficiency should be significantly improved by the future near-IR multi-object ground-based and space facilities. However, strong lensing clusters remain a factor of $\\sim 5-10$ more efficient than blank fields in the z $\\sim 7-11 $ domain, within the FOV of a few arcminutes around the cluster core, for the typical depth required for this survey project. ",
        "introduction": "Considerable advances have been made in the exploration of the high-z Universe during the last decade, starting with the discovery and detailed studies of redshift $z\\sim 3$ galaxies (Lyman break galaxies, LBGs), mostly from the pioneering work of Steidel and collaborators (cf.\\ Steidel et al.\\ 2003, Shapley et al.\\ 2003), the $z \\sim$ 4--5 galaxies found from different deep multi-wavelength surveys, to galaxies at $z \\sim$ 6--7, close to the end of the reionisation epoch of the Universe (e.g. Kodaira et al.\\ 2002, Hu et al.\\ 2002, Cuby et al.\\ 2003, Kneib et al. 2004, Stanway et al.\\ 2004, Bouwens et al.\\ 2004). To extend the present searches beyond $z\\ge$ 6.5 and back to ages where the Universe was being re-ionized (cf.\\ Fan et al.\\ 2002), it is mandatory to move into the near-IR bands. According to WMAP results, the first generation of stars may exist at $z \\sim$ 14--20 (Bennet et al. 2003; Spergel et al. 2003).  At z$>7$, the most relevant signatures are expected in the near-IR $\\lambda \\gtapprox 1 \\mu$m window. Most feasibility studies during the last years have been motivated by JWST, which should be able to observe galaxies up to redshifts of the order of $z \\sim$ 20-30. However, the delay of JWST and the availability of well suited near-IR facilities in ground-based 8-10m telescopes hasten the development of observational projects targeting $z \\gtapprox 7$ sources. The first studies on the physical properties of these sources can be started nowadays with instruments such as ISAAC/VLT, and pursued with improved efficiency using the future multi-object spectrographs to come, such as EMIR at GTC ($\\sim$ 2006), or KMOS for the second generation of VLT instruments ($\\sim$ 2009). In other words, the efficiency of this research is closely related to the instrumental developments in the near-IR domain.  We present here the first results from our ongoing pilot project with ISAAC/VLT, aimed at the detection of $z>7$ sources through lensing clusters used as Gravitational Telescopes (GTs). Our goal is to build up the first spectroscopic sample of galaxies at $z>7$ starting with GTs, in preparation of future massive surveys in the field. Such sample should allow to start a detailed study on the physical properties of these objects (SFR, extinction, metallicity, IMF, ...), and also start constraining their global properties (number counts, luminosity function, clustering properties, ...).  The plan of the paper is as follows. We review in section~\\ref{identification} the techniques and main results recently obtained on the identification of high-z galaxies in the optical bands, both in the field and using lensing clusters. In Sect.~\\ref{z7-10} we summarize the observational properties expected for galaxies at 7 $\\ltapprox$ z $\\ltapprox$ 11. Our pilot survey with ISAAC/VLT, aimed at the detection of z$>7$ sources, is presented in details in Sect.~\\ref{VLTsurvey}. The first results issued from our photometric and spectroscopic survey are discussed in this section, in particular the preliminary confirmation rate, and the global properties of our high-z candidates, including the latest results on the possible $z=10.0$ candidate A1835-1916. A discussion on the preliminary results and constraints derived from the 7$\\le$ z $\\le$ 10 photometric candidates is given in Sect.~\\ref{discussion}, together with future prospects. Throughout this paper, we adopt the cosmological parameters: $\\Omega_{\\Lambda}=0.7$, $\\Omega_{m}=0.3$ and $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "The bottom line in the present searches for z $\\sim 7-11 $ galaxies is efficiency. The different techniques used to identify high-$z$ candidates (narrow-band or broad-band searches) are complementary. They have to be improved and optimized in the near future to enhance the confirmation rates. Both the identification procedures and the individual high-$z$ sources issued from these procedures should be validated, if possible through complementary observations of the same field.  The search efficiency should be significantly improved by the future near-IR multi-object ground-based and space facilities. However, strong lensing clusters remain a factor of $\\sim 5-10$ more efficient than blank fields in the z $\\sim 7-11 $ domain, within the FOV of a few arcminutes around the cluster core, for the typical depth required for these survey projects."
    },
    "0410/nlin0410028_arXiv.txt": {
        "abstract": "We address the occurrence of narrow planetary rings under the interaction with shepherds. Our approach is based on a Hamiltonian framework of non--interacting particles where open motion (escape) takes place, and includes the quasi--periodic perturbations of the shepherd's Kepler motion with small and zero eccentricity. We concentrate in the phase--space structure and establish connections with properties like the eccentricity, sharp edges and narrowness of the ring. Within our scattering approach, the organizing centers necessary for the occurrence of the rings are stable periodic orbits, or more generally, stable tori. In the case of eccentric motion of the shepherd, the rings are narrower and display a gap which defines different components of the ring. ",
        "introduction": "Narrow rings have impressed us ever since their observation by Galileo. With the discovery of the rings of Uranus in 1977 by stellar--occultation measurements, the view of the rings as a special feature of Saturn changed completely. This went on with the information obtained in particular by the Voyager missions, which provided for the first time a close view into the complexity that takes place in the ring systems. We now know that rings appear in all giant planets. They display different morphologies, sizes, masses and physical properties. Different relevant physical phenomena takes place on them: Surface--wave phenomena, electromagnetic forces acting on micrometric charged dust grains and gravitational confinement due to nearby satellites are some examples. Different aspects of ring structure, dynamics and open problems are reviewed in Refs.~\\cite{rings,Nicholson91,Esposito02}. New information is becoming available nowadays by the Cassini mission.  After the discovery of Uranus' rings, which are extremely narrow in comparison to those of Saturn, models including confinement due to nearby moons were introduced. In particular, we mention the shepherding model introduced by Goldreich and Tremaine~\\cite{Goldreich79}, where two (shepherd) moons around the Uranian rings are proposed to bound and to define the sharp edges of the rings. The mechanism to maintain the ring involves angular momentum transfer between the shepherd moons and the ring particles as well as viscous damping due to interparticle collisions. While the full scenario for shepherding has not been fully understood, the presence of dissipation seems to be essential for the argument, which is in addition formulated assuming that the ring boundary is located at a lower--order resonance (see~\\cite{Esposito02}). In this paper, we propose a purely Hamiltonian formulation of the shepherding mechanism that confines the ring, focusing in particular in the phase--space structure that characterizes it. We do this by assuming, in a first approximation, an independent--particle model, i.e., where the ring particles do not interact among their selves. Within our approach, we succeed in explaining the confinement, on the one hand, and also long standing problems like the narrowness of the ring and its sharp edges.  We choose a rather simple but unrealistic example to illustrate the relevant structures in phase space and understand its implications. We emphasize that the example we define here, which is based on a billiard system with an arbitrary large number of non--interacting particles, is simply a rough cartoon of the real interaction between the particles of the ring and the shepherd moons. We are not proposing by any means that the planetary rings are due to real collisions between the particles of the ring and the shepherds, or that the interactions among particles of the ring are unimportant. Our goal is to point out that, within the independent particle picture considered here, the conclusions we can draw from the phase--space structure are {\\it generic} and therefore of interest in the context of planetary rings. The reason is that the relevant phase--space structures are independent of the precise details of the system. This is the reason we cautiously refer to it as an {\\it example} or {\\it toy} model rather than a {\\it model} for the (shepherded) narrow rings. The reader should not be misguided by this, which simply avoids the complications of working explicitely with $1/r$ potentials.  The paper is organized as follows: In Sec.~\\ref{sec_toy} we define the toy model within a Hamiltonian formulation, using an inertial reference frame and, what turns to be more general and convenient, using rotating--pulsating coordinates. We consider only the influence of what it would be interpreted as one shepherd moon for simplicity. In Sec.~\\ref{sec3} we consider the occurrence of rings in the case of circular motion of the shepherd. New results including analytical estimates of the bounds of the ring are presented here. In Sec.~\\ref{sec4} we consider the case of small but non--zero eccentricity, and show how the arguments carry on in this case. Here, we present indirect evidence for the existence of trapped motion in the system under consideration. Finally, in Sec.~\\ref{concl} we present our conclusions and some outlook of the work. ",
        "conclusions": "\\label{concl}  In this paper we have discussed generic mechanisms in phase space for the occurrence of narrow eccentric planetary--like rings due to the interaction with shepherd moons. Our approach is based on a Hamiltonian formulation and assumes non--interacting particles. The non--interacting character of our approach is only a first approximation in the same sense that the mean--field theory is for nuclear physics. We have used as a simple but unrealistic example for the argumentation: An open (scattering) billiard system on a Kepler orbit. We have focused on the structure of phase space, and extended previous results to the case of an eccentric Kepler orbit. Although the general arguments carry over, the generalization is not straightforward: The enlarged dimensionality introduces new possibilities in the stability properties of the fixed points and on the bifurcation scenario involving now at least two parameters. We find that the eccentric character of the orbit indeed influences the structure of the ring, which turns out to be narrower in comparison to the circular counterpart and also displays a gap. We have no satisfactory explanation for the latter aspect yet, and do not know whether it is an artifact of our model example or something with deeper significance. We have also provided a more detailed analytical framework to study the circular case, in particular providing estimates for the width of the ring. We emphasize that these new results are by no means a trivial extension: The dimensionality of the system implied through the time--dependence makes this a subtle task. To find rings in the present toy example which has no attractive interaction shows the genericity of the arguments.  The relevant structures in phase space correspond to those of pure trapped motion, which are organized around stable periodic orbits, or more generally, around stable tori. In phase space these regions are bounded by the invariant manifolds of unstable orbits. Within these regions no escape is possible. If there exists such a set of pure bounded motion in phase space, it is strictly of non--zero measure.  An ensemble of (non--interacting) particles defined by their initial conditions within these regions yields in position space a ring. The motion of particles with initial conditions in these regions must be contrasted with those outside: In the latter case, the particles are ejected from the system after few periods of rotation. This situation leads naturally to sharp--edges of the ring. This scenario is {\\it robust}, in the sense that it does not depend on the particular interaction between the particles of the ring, the shepherd moons and the central planet: The rotation of the shepherd moon by itself can build the conditions to have trapped motion. Once there is pure trapped motion the existence of such a ring follows immediately. That is, within a more realistic description, non--keplerian effects like planet oblateness can also be accounted for. The rings obtained display sharp--edges, are eccentric and narrow. We emphasize that these properties do not assume any specific resonance condition between the particle motion and the shepherd moons, as it is often considered or assumed. It is the existence of stable motion and its organizing centers in phase space what determines the existence of such a ring. We believe that this scenario may provide an explanation to the observed feature of apse alignment, which remains an open problem within the usual shepherding model.  This example can be taken over to consider the interaction with two disks on non--overlapping keplerian orbits, in order to model the more realistic situation of two shepherd moons. Using the same ideas, it is possible to define a (geometrical) criterion essentially given by the intersection of the ring (due to the outer disk) and the orbit of the inner one. This criterion may display a specific dependence on the actual position of the outer ring in the case of eccentric orbits. A possible generalization to the case with smooth potentials and in particular the $1/r$ case~\\cite{Benet01}, should maybe incorporate the results related to the stability of the organizing center, and then provide estimates of the width of the ring.  Interparticle interactions do indeed play a crucial role in the dynamics of narrow rings. Our current point of view is that they can be incorporated within the present modeling, at least for a dilute ring, as some chaotic diffusion (Brownian motion) in phase space, with probably some constrains. Such diffusive process induces small fluctuations in the value of the Jacobi integral and in the instantaneous direction of motion of the particles. One may thus impose that such a collision process conserves the total angular momentum of the colliding particles. This will effectively make that the dimension of phase space incorporates the number of particles. Future research will be carried along this direction."
    },
    "0410/astro-ph0410227_arXiv.txt": {
        "abstract": "A long standing problem with asymptotic giant branch (AGB) star models has been their inability to produce the low-luminosity carbon stars in the Large and Small Magellanic Clouds. Dredge-up must begin earlier and extend deeper. We find this for the first time in our models of LMC metallicity. Such features are not found in our models of SMC metallicity.  The fully implicit and simultaneous stellar evolution code {\\sc{stars}} has been used to calculate the evolution of AGB stars with metallicities of $\\mathrm{Z}=0.008$ and $\\mathrm{Z}=0.004$, corresponding to the observed metallicities of the Large and Small Magellanic Clouds, respecitively. Third dredge-up occurs in stars of 1\\ms\\ and above and carbon stars were found for models between 1\\ms\\ and 3\\ms. We use the detailed models as input physics for a population synthesis code and generate carbon star luminosity functions. We now find that we are able to reproduce the carbon star luminosity function of the LMC without any manipulation of our models. The SMC carbon star luminosity function still cannot be produced from our detailed models unless the minimum core mass for third dredge-up is reduced by 0.06\\ms. ",
        "introduction": "Carbon stars are stars which show the features of carbon rich molecules such as C$_{2}$ and CN in their spectra indicating they have a surface carbon-to-oxygen ratio (by number) greater than unity. They are typically identified in photometric surveys using the colours from narrow band filters centred near $7800$\\AA\\ and $8100$\\AA\\ (Cook \\& Aaronson 1989). The first filter focuses on a TiO absorption feature whilst the latter centres on a CN absorption feature. A census of carbon stars has been performed in the Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC). In the LMC a total of 7750 carbon stars have been found; in the SMC 1707 are known (Groenewegen 2004). The distances to the LMC and SMC are well determined and so we are able to construct luminosity functions for the carbon stars in both these locations. We are thus presented with an opportunity to test our models of stellar evolution.  It is believed that two populations of carbon stars exist which reflect two different formation routes. These populations are referred to as {\\it intrinsic} and {\\it extrinsic}. The intrinsic carbon stars are believed to be thermally pulsing asymptotic giant branch (TP-AGB) stars. It has long been known that the unstable double-shell burning in AGB stars leads to the phenomenon of third dredge-up (TDUP). The ashes of helium burning are mixed into the convective envelope enhancing the surface abundance of carbon (Iben \\& Renzini 1983). In low-mass stars repeated occurence of TDUP gives rise to carbon stars because the products of helium burning are carbon plus a little oxygen. However, for stars of around 4\\ms\\ or greater (the limit varying with metallicity), the temperature at the base of the convective envelope can become high enough for the CN cycle to operate. Carbon is converted into nitrogen via proton captures (e.g. Iben \\& Renzini 1983). This prevents the conversion to a carbon star. If sufficient mass is stripped from the envelope of such a star then the hot-bottom burning can cease and a carbon star can then form. The extrinsic carbon stars are believed to be formed by the accretion of carbon-rich material from a more evolved donor (Van Eck et al. 1998).  There is a long-standing problem with forming intrinsic carbon stars. Detailed stellar evolution models have so far proved unable to produce dredge-up at low enough core masses and hence at low enough luminosities. When Iben formulated the carbon star mystery (Iben 1981) calculations showed that dredge-up did not occur for core masses below about 0.6\\ms\\ (e.g. Sackmann 1980, Wood \\& Zarro 1981). This is also borne out by more recent calculations but leads to problems in reproducing the observed luminosity functions of the LMC and the SMC. Using the results of Karakas et al. (2002), Izzard et al. (2004) determined that the core masses at dredge-up would have to be lower by 0.07\\ms.  Most modern stellar evolution codes do not solve the equations of stellar structure, nuclear burning and mixing simultaneously. Mixing is typically solved for after the structure and burning have been determined. Examples include the codes used by Straniero et al. (1997) and Herwig (2000). Karakas et al. (2002) employ a partial simultaneous approach. A structure iteration is followed by a mixing iteration before another structure iteration is made. In this way a model is converged. Initial work on a fully implicit and simultaneous solution of these equations (whereby the structure, burning and mixing are all treated together in the {\\it same} iteration) together with an improved treatment of diffusive mixing, on the TP-AGB was done by Pols \\& Tout (2001). More recently Stancliffe, Tout \\& Pols (2004) found that it was possible to obtain dredge-up at lower core masses than was previously thought.  Following on from that work, the evolution of stars of between 1 and $6\\,$\\ms\\ at $\\mathrm{Z}=0.008$ and $\\mathrm{Z}=0.004$ is investigated. The outcome of this detailed modelling is used as the input physics to a population synthesis code in order to compare the luminosity functions predicted by the data with the observed LMC and SMC carbon star luminosity functions (CSLF). In section 2 we discuss the stellar evolution code used and the details of the models. The results of the these calculations are presented in section 3. The population synthesis code is briefly discussed in section 4 where we present our luminosity functions. ",
        "conclusions": "We have made evolution calculations, with a fully implicit and simultaneous code, of the thermally pulsing asymptotic giant branch evolution of low- and intermediate-mass stars at LMC and SMC metallicities. For the LMC we find more efficient dredge-up at lower core masses than others. The SMC models are comparable to those of Karakas et al. (2002). Using a population synthesis code calibrated to the detailed models we find we are able to reproduce the carbon star luminosity function of the LMC. The SMC CSLF can still only be reproduced if the minimum core mass required for TDUP is reduced by 0.06\\ms.  The fact that these calculations can reproduce, for the first time, the LMC CSLF is of great interest.  It is not clear why the present calculations succeed while others have failed, but the simultaneous solution of structure and mixing may be important. We are currently exploring the reasons for the difference between our results and those of others."
    },
    "0410/astro-ph0410361_arXiv.txt": {
        "abstract": "We present an updated timing solution for the  young, energetic pulsar \\psr\\ based on 21.3 years of radio timing data and 7.6 years of X-ray timing data. No glitches have occurred in this time span, in contrast to other well-studied young pulsars, which show frequent glitches. We report a measurement of the third frequency derivative of $\\nudotdotdot=(-1.28 \\pm 0.21) \\times 10^{-31}\\,$s$^{-4}$. This value is 1.65 standard deviations from, i.e. consistent with, that predicted by the simple constant magnetic dipole model of pulsar spin-down. We measured the braking index to be $n=2.839 \\pm 0.003$ and show that it varies by 1.5\\% over 21.3\\,yr due to contamination from timing noise. Results of a low-resolution power spectral analysis of the significant noise apparent in the data yield a spectral index of $\\alpha = -4.6 \\pm 1.0$ for the red noise component. ",
        "introduction": "\\label{sec:intro} Radio pulsars are powered by their rotational kinetic energy: as they emit electromagnetic radiation, their rotation rates decrease. Thus, measuring the rotational evolution of pulsars is a probe of the physics underlying these fascinating objects. Young, energetic pulsars, such as \\psr\\  are especially interesting due to their large spin-down luminosities and rapid spin-down rates. High rates of spin-down allow for the measurement of higher-order frequency derivatives, providing a deeper probe into the physics of spin-down. This slow-down can be described by \\begin{equation} \\dot{\\nu} = -K{\\nu}^n      , \\end{equation} where ${\\nu} \\equiv 1/P$ is the pulse frequency, $\\dot{\\nu}$ is its derivative and $n$ is the braking index. For braking by magnetic-dipole radiation in a vacuum, $K$ is related to the dipole magnetic moment of the pulsar and $n=3$ \\citep{mt77}. Differentiating (1) shows that $n$ is given by \\begin{equation} n=\\frac{ {\\nu}{\\ddot{\\nu}}}{\\dot{\\nu}^2}  . \\end{equation} The braking index is only measurable for the youngest pulsars and values not entirely dominated by noise processes have been measured for only 5 objects. In all cases, the measured values are less than the canonical value of $n=3$ \\citep{lps93,kms+94,lpgc96,dnb99,ckl+00}. Implications of observed braking indices less than 3 are still being discussed \\citep[see for example, ][]{br88} and include magnetic field growth or alignment with the spin axis, effects due to higher-order magnetic moments and the interaction of the pulsar with its relativistic particle wind.  Most pulsars do not have measurable braking indices due to a combination of contamination from low-frequency noise processes and the relatively slow spin-down of middle-aged pulsars. Because of the form of the spin-down power law given in equation 1, successive frequency derivatives go approximately as powers of the first frequency derivative. Thus most pulsars, which have values of $\\dot{\\nu}$ of $\\sim 10^{-15}$\\,s$^{-2}$, should exhibit $\\ddot{\\nu}$ on the order of $\\sim 10^{-30}$\\,s$^{-3}$, which would require hundreds of years of observation to measure. Many younger pulsars that could conceivably have measurable $\\ddot{\\nu}$ exhibit significant timing noise which masks the deterministic value of $\\ddot{\\nu}$ due to dipole spin-down, resulting in measured values orders of magnitude larger than predicted, often with the wrong sign. Furthermore, these values are not stable and consequently do not accurately predict pulse arrival times. \\par An important check on the validity of the spin-down law is the measurement of deterministic higher-order frequency derivatives. An expression for $\\nudotdotdot$ is given by taking an additional derivative of equation (1), \\begin{equation} \\nudotdotdot\\ = \\frac{n(2n-1){\\dot{\\nu}^3}}{\\nu^2}. \\end{equation} We define the second braking index to be $m_0 \\equiv n(2n-1)$. If equation 1 accurately characterizes the spin-down of the pulsar, then $m_0=m$, where \\begin{equation} m = \\frac{{\\nu}^2 \\nudotdotdot}{\\dot{\\nu}^3}. \\end{equation} A measurement of $m$ provides important insight into the spin-down of young pulsars and $m \\neq m_0$ implies that the characteristic age is an under- or over- estimate of the true age, as well as having implications for a changing magnetic field \\citep{bla94}. Evidently, the measurement of higher-order frequency derivatives will only be possible in the youngest pulsars that have large spin-down rates not dominated by noise processes. The Crab pulsar has a measured value of $\\nudotdotdot = (-6.45\\pm0.02)\\times 10^{-31}\\,$s$^{-4}$ \\citep{lps93}, in agreement with the prediction from the spin-down law, in spite of probable contamination from timing noise and frequent glitches.  \\psr\\ has never been observed to glitch, hence is an excellent candidate to have a measurable $\\nudotdotdot$, as well as a very accurate measurement of $n$. \\citet{kms+94}  measured a value for $\\nudotdotdot = (-1.02 \\pm 0.25) \\times 10^{-31}\\,$s$^{-4}$, based on 11\\,years of radio timing data, implying $m=14.5 \\pm 3.6$, in agreement with the predicted value $m_0=13.26\\pm 0.03$. The uncertainty in $\\nudotdotdot$ and thus $m$ was known to be larger than the formal uncertainty due to contamination from significant timing noise in the data. To account, albeit not rigorously, for this effect, the quoted uncertainty is three times the formal uncertainty. \\par In this paper, we extend the timing analysis performed by \\citet{kms+94} using an additional 10.3\\,yr of radio data as well as 7.6\\,yr of X-ray timing data to increase the precision with which the frequency evolution of \\psr\\ can be measured. This is accomplished by performing a phase-coherent analysis of the the data as well as a partially phase-coherent analysis that is less sensitive to the timing noise superimposed on the deterministic spin-down. In addition, we present a low-resolution spectral analysis of the timing noise and discuss possible physical causes of the noise process. ",
        "conclusions": "\\label{sec:discussion} \\subsection{Timing Analysis} \\label{sec:timing_discussion} The measurement of the third frequency derivative, $\\nudotdotdot = (-1.28 \\pm 0.21) \\times 10^{-31}$s$^{-4}$ is 1.65 standard deviations from the value predicted by the spin-down law. Hence the observed value and the model prediction are consistent. Typically, uncertainties on $\\nudotdotdot$ fall as $t^{-4}$, though the prominence of timing noise in this case will undoubtably lengthen the amount of time required. Based on the bootstrap analysis used to determine uncertainties, another $\\sim 10$\\,yr of observations are required to reduce uncertainties by a factor of $\\sim 2$. Assuming that we have measured a stable value of $\\nudotdotdot$ and that the uncertainties will be reduced by a factor of 2, then the simple spin-down law derived assuming a constant magnetic field may not sufficiently describe the spin-down of \\psr.  \\par The consequences of a measured $\\nudotdotdot$ different from the predicted value are most clear when comparing the measured and predicted values of the second braking index. The measured value of $\\nudotdotdot$ implies a second braking index $m= 18.3 \\pm 2.9$, larger than $m_0=13.26\\pm 0.03$, but not significantly so given the uncertainties. \\citet{bla94} discusses the implications of $m \\neq m_0$, namely that $K$ is a function of time. This discussion is particularly useful because it does not require knowledge of the true braking index, $n_0$, which may be different from the measured value and cannot be predicted from theory. However, since we do not make any assumption about $n_0$, we cannot say whether $K$ is changing at the present time. A measurement of $m \\neq m_0$ does give us insight into the variation of $K$ in the pulsar's past. In particular, a value of $m<m_0$ implies a period of magnetic field growth early in the pulsar's history and a true age larger than the characteristic age, while $m>m_0$ reduces the inferred age of the pulsar.  Given $m-m_0=5.04\\pm2.99$, and using Blandford's prescription for the inferred age of the pulsar, the true age of \\psr\\ is 1000$\\pm$700\\,yr, just over half of the characteristic age, $\\tau_c = 1691$\\,yr. \\par The supernova remnant (SNR) believed to be associated with \\psr\\ was originally estimated to have an age of $6-21$\\,kyr \\citep{shmc83}, indicating that the pulsar's characteristic age was underestimating the age of the system by a factor of $\\sim3.5-12$. However, \\citet{gbm+99} showed that the SNR may be as young as the characteristic age of the pulsar. They show that the SNR and pulsar are interacting and explain the large size and unusual morphology of the remnant by hypothesizing that the remnant is due to an explosion of high kinetic energy or low ejected mass occurring on the edge of a cavity, thus accelerating the expansion of the remnant. At first glance, the measured value of $m$ appears to worsen the possible age discrepancy between the pulsar and the SNR. However, the age of the remnant is clearly not well determined and could be as low as 1000\\,yr.  \\par \\citet{br88} discuss the implications of a constant observed braking index, $n$, where in this case we must make some \\textit{a priori} assumption of the true braking index, typically assumed to be $n_0=3$. They give several reasons why $n<3$ may be observed. The torque may scale with frequency differently, either due to the outflow of plasma removing a significant amount of angular momentum, or a non-dipolar magnetic field.  Another interpretation of $n<3$ is that the magnetic moment is currently varying in time, that is, $K=K(t)$. If this is the case, we can then determine a time scale for torque evolution and a check on the functional form of evolution by finding the first two dimensionless derivatives of $\\dot{\\nu} = -K(t){\\nu}^n$: \\begin{equation} d_1 = n - n_0 = {\\frac{\\dot{K}}{K}}\\,{ \\frac{\\nu}{{\\dot{\\nu}}}}, \\end{equation} and  \\begin{equation} d_2 = m - n_0[2n_0 - 1 + 3(n-n_0)] = \\frac{\\ddot{K}}{K} \\frac{{\\nu}^2}{{\\dot{\\nu}}^2}. \\end{equation} Equation (9) gives a time scale for torque evolution while equation (10) provides a check for an assumed functional form for $K(t)$. For \\psr, $d_1=-0.161 \\pm 0.003$ and $d_2=3.5 \\pm 3$. For exponential growth of $K$, this measurement of $d_1$ implies a time scale for growth of $\\sim 20~000$\\,yr. If the form of $K$ is exponential we would expect that $d_2=(d_1)^2=0.026$, which cannot be ruled out with the present uncertainties for $d_2$. We find that the present measurements are similarily in agreement with a power-law form for $K(t)$. The measurement of $m$ and the large uncertainties, however, are consistent with the simplest explanation of a constant value of $K$ and thus a constant magnetic field. \\par Despite scatter due to timing noise, the braking index is constant over 21.3\\,yr to within 1.5\\%. This result is similar to that of the Crab pulsar's braking index, reported to be constant to within 0.5\\%, where measurements were obtained between glitch events and are typically based on five years of data \\citep{lps93}. This reinforces the need for long-term timing to measure the true value of braking indices to such precision.  \\subsection{Timing noise analysis} \\label{sec:noise_discussion} \\subsubsection{Glitches} Our phase-coherent analysis shows that, due to the absence of any sudden discontinuities in the timing residuals,  \\psr\\ has not glitched in 21.3\\,yr. \\psr\\ is the only young pulsar known that has never glitched over a long period of time.  \\citet{lyn99} showed that there is a correlation not only between $\\dot{P}$ and timing noise, but $\\dot{P}$ and glitch activity, although some of the youngest pulsars (i.e. the Crab) have glitch activity parameters smaller than expected from the correlation. The explanation for this is that the internal temperature of the youngest pulsars is too high for large scale vortex pinning and unpinning. If this is the case, then the fact that \\psr\\ has the lowest glitch rate among young pulsars would suggest that it has the highest internal temperature, in spite its characteristic age being larger than that of the Crab pulsar. On the other hand, so-called `anomalous X-ray pulsars' (AXPs) have been observed to glitch with parameters that are not very different from those seen in radio pulsars \\citep{klc00,kg03,dos+03}.  Given their estimated surface temperatures and luminosities (measured via X-ray spectral observations) AXPs appear to be much hotter than any known radio pulsar. If AXP glitches indeed have the same physical origin as do glitches in rotation-powered pulsars, this argues that the glitch phenomenon is not strongly affected by neutron-star temperature.  \\subsubsection{Spectral Analysis} The spectral index for the timing noise of \\psr, $\\alpha=-4.6\\pm1.0$ is marginally steeper than that measured for other pulsars. The low-resolution nature of the Deeter method, however, provides only a rough estimate of the true spectral index of the noise. The spectral index of the timing noise for the Crab pulsar has been determined several times, the most recent of which found $\\alpha=-3$ as well as a periodic component with period of $568\\pm10$\\,days \\citep{sfw03}, using a modified Fourier analysis (possible in this case due to the dense sampling of the Crab timing data from 1982-1989). The spectral index was then  verified with the Deeter polynomial method, though this analysis does not probe the lowest frequencies (i.e. longest timescales) available. Previous analyses of Crab timing data were consistent with $\\alpha \\sim -4$ \\citep{dee81, cor80}. \\citet{ch80} show that the timing noise of 11 pulsars is consistent with a pure random walk in phase, frequency or frequency derivative. However, \\citet{dmh+95} show in a study of 45 pulsars that only 5 are consistent with a pure random walk, while most pulsars exhibit timing noise either due to a combination of a random walk and discrete jumps, or inconsistent with either hypothesis.  Other pulsars have been shown to have spectral indices of timing noise ranging from $\\alpha \\sim -0.5$ to $-2.4$ \\citep{babd99}. Thus there is a range of possible spectral indices that describe timing noise in different pulsars.  \\par Although there is no physical reason why the spectral index must be an even integer, these are the only cases for which analytical solutions exist. For \\psr, the measured spectral index is consistent with having a component of frequency noise, that is, a random walk in pulse frequency corresponding to $\\alpha= -4$. The measured value of the spectral index indicates that there may also be a component due to a random walk in frequency derivative, characterized by $\\alpha = -6$. Although the measured spectral index alone would seem to indicate a random walk in frequency, the fact that the residuals are not Gaussian distributed after fitting for higher derivatives indicates that this is not the case. A random walk in the $k$th derivative of phase is mathematically equivalent to a white noise process in the $(k+1)$ derivative \\citep{gro75a}. This should correspond to `white' residuals after fitting out the $(k+1)$ derivative, not seen in the timing residuals of \\psr.  \\par One of the problems in quantifying timing noise is the lack of a comprehensive theory of its physical causes. In fact it is possible that there are several causes of timing noise active in one object, thus complicating the interpretation of the measured spectral index. Free precession is an interesting solution to the timing noise problem because it can be detected from the changes in the pulse profile as well as a periodicity in the timing residuals \\citep{sls00}. Unfortunately, the data quality and lack of a long baseline for many pulsars prohibits the discovery of this process. If undetected, precession would contribute to the apparent timing noise of the pulsar and the spectral index of the noise, despite the fact that it is not a random process.  Additionally, other physical causes of noise, (e.g. magnetospheric torques or random vortex pinning and unpinning) will contribute to the measured spectral index.  Another common problem in assessing the cause of timing noise is that quasi-periodicities are apparent in the timing residuals of most pulsars. An artificial quasi-periodicity will change with the length of the data set, while a real quasi-periodicty will not \\citep{klw+99}. To begin to truly understand timing noise in pulsars, it is crucial that new high-resolution spectral methods, capable of discerning discrete features due to periodicities, are developed. New methods must be able to deal with both the spectral leakage and data sampling problems that presently constrain our measurements and thus understanding of timing noise. Another problem in developing high-resolution methods for very long period noise is the extremely long time baselines required to observe many periods of the noise; the only way to solve this is by brute force, that is, by timing pulsars for hundreds of years. A better understanding, both mathematically and physically, of timing noise will allow for a better understanding of the deterministic spin-down of pulsars underlying the noise processes."
    },
    "0410/astro-ph0410682_arXiv.txt": {
        "abstract": "We consider a thin ($\\sim 10^2-10^3$ fm) layer of electrons (the electrosphere) at the quark surface of a bare strange star, taking into account the surface effects at the boundary with the vacuum. The quark surface holds the electron layer by an extremely strong electric field, generated in the electrosphere to prevent the electrons from escaping to infinity by counterbalancing the degeneracy and thermal pressure. Because of the surface tension and depletion of $s$ quarks a very thin (a few fm) charged layer of quarks forms at the surface of the star. The formation of this layer modifies the structure of the electrosphere, by significantly changing the electric field and the density of the electrons, in comparison with the case when the surface effects are ignored. Some consequences of the modification of the electrosphere structure on the properties of strange stars are briefly discussed. ",
        "introduction": "Strange stars made entirely of strange quark matter (SQM) have long been proposed as an alternative to neutron stars (Alcock, Farhi, \\& Olinto 1986; Haensel, Zdunik, \\& Schaeffer 1986; Glendenning 1996; Cheng, Dai, \\& Lu 1998; Weber 1999; Cheng \\& Harko 2000; Alford \\& Reddy 2003; Lugones \\& Horvath 2003 and references therein). The possible existence of strange stars is a direct consequence of the conjecture that SQM, composed of deconfined $u$, $d$, and $s$ quarks, may be the absolute ground state of the strong interaction, i.e., it is absolutely stable with respect to $^{56}$Fe \\cite{B71,W84,FJ84}. If SQM is approximated as a noninteracting gas of quarks, chemical equilibrium with respect to the weak interaction, together with the relatively large mass of the $s$ quarks imply that the $s$ quarks are less abundant than the other quarks. Hence in SQM electrons are required to neutralize the electric charge of the quarks. The electron density at vanishing pressure is of the order of $\\sim 10^{-4}$ of the quark density \\cite{AFO86,KWWG95}.  The electrons, being bounded to SQM by the electromagnetic interaction alone, are able to move freely across the SQM surface, but clearly they cannot move to infinity, because of the bulk electrostatic attraction to the quarks. The distribution of electrons extends several hundred fermis above the SQM surface, and an enormous electric field $E\\simeq 5\\times 10^{17}$ V~cm$^{-1}$ is generated in the surface layer to prevent the electrons from escaping to infinity, counterbalancing the degeneracy and thermal pressure \\cite{AFO86,KWWG95,HX02}.  The thin layer of electrons with a very strong electric field, which is called the \"electrosphere\", drastically affects the observational appearance of strange stars. First, the electrosphere of a hot strange star with a bare SQM surface may be responsible for its thermal emission in a wide range of the surface temperature $T_S$. The point is that the electric field at the elecrosphere is a few ten times higher than the critical field $E_{\\rm cr}\\simeq 1.3\\times 10^{16}$ V~cm$^{-1}$, at which vacuum is unstable to creation of $e^+e^-$ pairs \\cite{S51}. Therefore, a hot strange star with a bare SQM surface may be a powerful source of $e^+e^-$ pairs which are created in the electrosphere and flow away from the star (Usov 1998, 2001a; Aksenov, Milgrom, \\& Usov 2003, 2004). Emission of $e^+e^-$ pairs from the electrosphere dominates in the thermal emission of a bare strange star at $6\\times 10^8~{\\rm K}\\lesssim T_S\\lesssim 5\\times 10^{10}$ K, while below this temperature, $T_S\\lesssim 6\\times 10^8$ K, bremsstrahlung radiation of photons from electron-electron collisions in the electrosphere prevails \\cite{JGPP04}. Besides, the flux of photons generated in the surface layer of SQM via any mechanism (for example, quark-quark bremsstrahlung) may be strongly reduced in the process of propagation through the electrosphere \\cite{CH03}. Second, the surface electric field may be also responsible for existence of a crust of \"normal\" matter (ions and electrons) at the SQM surface of a strange star \\cite{AFO86}. This field is directed outward, and the ions in the inner layer are supported against the gravitational attraction to the underlying strange star by the electric field.  Recently, it was argued that the properties of the electrospheres may be changed essentially because of the surface effects \\cite{U04}. It is the purpose of the present paper to study the structure of the electrospheres by taking into account the surface effects in details. The different QCD phases of SQM are also considered. The inclusion of the surface effects drastically modifies the structure of the electrosphere, by modifying the surface electric fields and the number density of the electrons.  The remainder of the paper is organized as follows. In \\S 2 we formulate the equations that describe the distributions of the electric fields and the density of the electrons is the electrosphere and the boundary conditions. In \\S 3 we obtain the solutions of the equations describing the  properties of the electrosphere and discuss their properties. Finally, in \\S 4 we summarize our results and discuss some astrophysical applications. ",
        "conclusions": "In this paper, we have considered the structure of extremely strong electric fields and the distribution of the density of electrons in the vicinity of the SQM surface of a bare strange star.  It is taken into account that a very thin (a few fm) charged layer forms at the SQM surface because of surface depletion of $s$ quarks. We have shown that the surface charged layer changes significantly both the field structure and the density of electrons as compared with the same calculated in many papers (e.g., Alcock et al. 1986; Kettner et al. 1995; Hu \\& Xu 2002; Cheng \\& Harko 2003), where the surface effects are ignored. These changes are especially large if the surface temperature is not very high ($T_S\\lesssim$ a~few MeV).  For a bare strange star the structure of electric fields and the density of electrons near the stellar surface depend on $m_s$, $\\mu$, and $T_S$ [see equations (\\ref{nezT}), (\\ref{Ez>0}) - (\\ref{z0}), (\\ref{dVdzf}) - (\\ref{V0SQM}) and \\S2]. For SQM at vanishing pressure the value of $\\mu$ is about 300 MeV with accuracy of about 20\\%, and the main uncertainty of the parameters of the electrospheres of bare strange stars is because of uncertainty of $m_s$ that is estimated as 150 MeV within a factor of 2-3 or so.  The density of electrons in the electrosphere of a bare strange star increases because of the surface effects by a factor of few ten (see \\S~3). This may change significantly the thermal emission from the stellar surface. In the Introduction, the main mechanisms of the thermal emission are mentioned. They are creation of $e^+e^-$ pairs in a supercritical electric field \\cite{U98}, quark-quark bremsstrahlung from the surface layer of SQM \\cite{CH03}, and electron-electron bremsstrahlung from the electron layer \\cite{JGPP04}. In supercritical electric fields, $E\\gg E_{\\rm cr}$, the rate of pair production when electrons are created into the empty quantum states is extremely high \\cite{S51}, and all the empty states are occupied by creating electrons practically instantly. Then, the rate of pair production in the electrosphere is determined by the process of thermalization of electrons which favors the empty-state production \\cite{U98}. Using this, we have roughly the following dependence of the strange star luminosity in $e^+e^-$ pairs on $V_0$: \\begin{equation} L_\\pm\\propto\\cases{{\\rm constant} & at $T_S < 0.01\\,eV_0$,\\cr V_0^{1.8}& at $0.01\\,eV_0\\lesssim T_S\\lesssim 0.1\\,eV_0$,\\cr V_0^3 & at $T_S > 0.1\\,eV_0$\\,,\\cr} \\label{Lpm} \\end{equation} where we use $V_0$ which is connected with the density of electrons at the electrosphere [see equation (\\ref{nezT})].  For the most conservative parameters of SQM ($m_s\\simeq 150$ MeV and $\\mu\\simeq 300$ MeV), the value $V_0$ increase $\\sim 3$ times because of the surface effects. Hence, at high temperatures $L_\\pm$ may increase up to a few ten times in comparison with the value calculated by Usov (1998) for the electrosphere model where the surface effects are ignored.  It has been shown by Cheng \\& Harko (2003) that quark-quark bremsstrahlung (and another radiation from SQM) may be strongly suppressed in the process of its propagation through the electrosphere. Most probably, the increase of the density of electrons by a factor of $\\sim 20-30$ because of the surface effects results in almost complete suppression of the outgoing radiation from SQM if the surface temperature is not very high ($T_S < 1~{\\rm MeV}$). In this case practically all radiation of a bare strange star is generated by the electron layer at $z>0$. In contrary, we expect that the electron density increase doesn't affect significantly the electron-electron bremsstrahlung radiation from the electrosphere at rather low temperature ($T_S < 0.1~{\\rm MeV}$) where this radiation prevails. Indeed, in this case the bremsstrahlung radiation is generated mainly at the distance from the SQM surface where the local plasma frequency of electrons is equal to the frequency of radiation (Jaikumar et al. 2004). Since at any temperature the plasma frequency of electrons is connected directly with the density of electrons, the density of electrons in the radiating region is more or less the same irrespective of the structure of the electron layer. Therefore, the electron-electron bremsstrahlung radiation at a given frequency is more or less the same, too. We hope to deal with detail calculations of the thermal emission from a bare strange star for the electrosphere model developed in this paper elsewhere.  Recently, the cooling of young bare strange stars has been studied numerically using the electrosphere model where the surface depletion of $s$ quarks is ignored \\cite{PU02}. It was shown that the thermal luminosity of such a star in photons and $e^+e^-$ pairs may be up to $\\sim 10^{49}- 10^{50}$ ergs~s$^{-1}$ for a few second after the star formation and remains high enough ($\\gtrsim 10^{36}$ ergs~s$^{-1}$) as long as the surface temperature is higher than $\\sim 0.1$ MeV. The increase of the density of electrons in the electrosphere because of the surface effects has to modify the light curves calculated by Page \\& Usov (2002). We expect that at the first stage of the strange star cooling, when the neutrino losses dominate, the thermal radiation from the stellar surface increases in accordance with equation (\\ref{Lpm}), i.e., by a factor of few ten at $T_S > 0.1\\,eV_0$. At the second stage, when the losses in the surface thermal radiation prevails, the thermal luminosity decreases more fast than in the case of the non-modified electrosphere, especially at low temperatures ($T_S < 0.1$ MeV) where the main mechanism of the thermal emission is the electron-electron bremsstrahlung that is ignored by Page \\& Usov (2002).  At the surface of a strange star a massive normal matter crust may be formed by accretion of matter onto the star. From Figure~3 we can see that at rather low temperatures ($T_S\\ll T_*\\sim 10$ MeV) the electrostatic potential of electrons $eV_0$ is more than the electron chemical potential ($\\sim 25$ MeV) at which neutron drip occurs \\cite{BPS71}. Therefore, the maximum density of the crust is limited by neutron drip and is about $4.3\\times 10^{11}$ g~cm$^{-3}$ \\cite{AFO86}. In this case, the maximum mass of the crust is $\\sim 10^{-5}M_\\odot$."
    },
    "0410/astro-ph0410157_arXiv.txt": {
        "abstract": "We re-examine the photometric and spectroscopic evidence available for the star OGLE-TR-56, recently discovered to harbor a giant planet presenting transits and orbiting with a period of 1.21 days. We investigate the possibility that the observational signatures reported might be the result of ``blending\" with the light of an eclipsing binary along the same line of sight. Using techniques developed earlier we perform fits to the light curve under a variety of blend configurations, subject to all observational constraints, and we infer further properties of these possible blends. We then carry out realistic end-to-end simulations based on those properties in order to quantify the spectral line asymmetries and radial velocity variations expected from such scenarios. We confront these calculations with the observations.  The results from these models are clearly inconsistent with the measured radial velocity and bisector span variations, ruling out blends and confirming the planetary nature of the companion. The example of OGLE-TR-56 serves to illustrate the sort of tests that can and should be performed on transiting planet candidates to eliminate the possibility of ``false positives\", and in particular, line-of-sight contamination from eclipsing binaries. ",
        "introduction": "\\label{sec:intro}  The search for extrasolar giant planets has entered an exciting new era, one in which these objects can be discovered indirectly by the tiny photometric signatures they produce as they transit across the disk of their parent star (periodic $\\sim$1\\% drops in brightness). Numerous photometric search programs are underway or under development worldwide \\citep[see][]{Horne:03}. Six stars with transiting planets have been found so far: HD~209458 \\citep{Henry:00, Charbonneau:00}, OGLE-TR-56 \\citep{Konacki:03a}, OGLE-TR-113 \\citep{Bouchy:04, Konacki:04}, OGLE-TR-132 \\citep{Bouchy:04}, TrES-1 \\citep{Alonso:04}, and OGLE-TR-111 \\citep{Pont:04}.\\footnote{Another case, OGLE-TR-3, was reported recently by \\cite{Dreizler:03} and claimed to harbor a transiting planet with a period of 1.19 days, but additional studies showed that it is most likely a false positive \\citep{Konacki:03b}.} The OGLE-TR planets were originally detected as low-amplitude transit events by the most successful of the photometric surveys, OGLE-III \\citep[Optical Gravitational Lensing Experiment;][]{Udalski:02b, Udalski:02c, Udalski:03}, and later confirmed by radial-velocity follow-up.  The planet around HD 209458 was originally discovered by its radial velocity signature ---the small reflex motion of the star in response to the pull from the planet--- and only later was it found to undergo transits.  It is the Doppler technique, in fact, that has been the most successful in discovering extrasolar planets, having produced more than 100 giant planets to date \\citep{Schneider:04}.  Both the Doppler and the transit search methods suffer from the problem of astrophysical ``false positives\", which refers to other phenomena that can produce the same observational signature but do not involve planets.  In the case of the Doppler searches the concern is that similar velocity variations might result from chromospheric activity on the star (e.g., spots), sub-surface convection (granulation), other line asymmetries from undetected stellar companions, or even stellar oscillations \\citep[see, e.g.,][]{Gray:97, GrayH:97, Gray:98}. For the transit searches the main contaminants are (1) eclipsing binaries with grazing geometry, (2) transits of a small star in front of a large star (such as an M dwarf passing in front of an F star, or a solar-type star in front of a giant), and (3) eclipsing binary systems with deep eclipses, blended with the light of another star along the same line of sight that dilutes the minima.  In the latter case (referred to as a ``blend\"), the intruding star may or may not be physically associated with the binary. It is these blend scenarios that are the subject of this paper.  Cases that only involve an eclipsing binary with no contaminants (scenarios [1] and [2] above) are relatively easy to dismiss with radial velocity measurements since the expected amplitudes are typically several tens of~\\kms, as opposed to a few hundred~\\ms\\ for the case of a Jupiter-mass planet around a solar-type star. Blend configurations, on the other hand, can be much more difficult to rule out as the main star may show little or no velocity variation at all, and may therefore appear to have a very low-mass companion in orbit when in reality it does not.  This is particularly the case for very faint candidates ($V = 14$--19) such as those being reported from surveys such as OGLE-III \\citep[e.g.,][]{Udalski:02a}, EXPLORE \\citep{Mallen-Ornelas:03}, MACHO \\citep{Drake:04}, and others.  For these objects the precision of the velocities is much lower than for the typically brighter stars ($V < 11$) that have been targeted in Doppler searches, and the potential for confusion is greater because they also tend to be in very crowded fields. It is therefore important to examine these cases very carefully.  In a recent paper \\citep[][hereafter Paper~I]{Torres:04b} we described a procedure to model the light curves of transit candidates as the combination of three stars forming the blend, and we applied it to the case of OGLE-TR-33, a candidate that turned out to be a false positive. We showed how such modeling tools allow one to predict other observable properties of the blend and can become a useful guide to look for additional signs of a contaminating eclipsing binary. One such sign that can easily be overlooked is asymmetries in the spectral lines, which can lead to spurious velocity measurements.  In the present paper we develop procedures to simulate line asymmetries and their effects on the measured velocities of a transit candidate resulting from typical blend scenarios.  We illustrate these techniques by carrying out a critical examination of the case of OGLE-TR-56 ($V = 16.6$), the first extrasolar transiting planet to emerge from the lists of very faint candidates, and the one with the shortest orbital period so far (1.21~days). We perform blend simulations subject to all observational constraints in an attempt to explain the system as a false positive, but despite all our efforts we are unable to find a plausible configuration involving a contaminating eclipsing binary that does not violate one or another of those constraints by a large margin. This lends further support to the planetary nature of the companion of OGLE-TR-56. ",
        "conclusions": "The fundamental reason for the failure of all of the blend scenarios to reproduce the observations for OGLE-TR-56 is that the variations expected in the radial velocities and in the bisector spans from our simulations are \\emph{of the same order of magnitude}, and typically differ by a factor of two or less, at least for configurations that are consistent with the light curve constraints.  Similar results were reported by \\cite{Santos:02}. In our case the predicted bisector span variations tend to be larger than the velocity variations.  By contrast, the bisector variations actually measured for OGLE-TR-56 are essentially negligible while the velocity variations are highly significant. Thus, no model can reproduce both simultaneously, and our tests show this discrepancy cannot be solved by adjusting the blend parameters.  A further reason blend models fail here is the strong observational constraint on the brightness of star~1. A star with only $\\sim$3\\% of the light of star~3 does not produce enough of an effect to be noticeable, either in the line profiles or in the measured radial velocities.  The strong rotational broadening of star~1, as predicted from our light curve fits, only reduces the effect even further.  The conclusion must be, then, that a blend configuration involving an eclipsing binary along the same line of sight (whether in a physical triple system, or in the background) cannot account for the observational signatures in OGLE-TR-56 (photometric and spectroscopic), and we are left with a planetary companion in orbit as the only viable explanation.  We are unable to devise a more contrived scenario, perhaps involving more stars, that does not go against one or another of the observational constraints.  Dozens of transit candidates are being produced by current deep narrow-field searches that focus on very faint stars, where crowding is potentially a serious issue due to the high stellar densities. Wide-field searches of brighter objects have also entered production mode, and suffer from a similar confusion problem because of the large angular pixel scales. The incidence of blends with a background eclipsing binary in both types of surveys is therefore expected to be rather high.  The same is true in most cases for future transit searches from space such as NASA's {\\it Kepler\\/} mission, which will look for even smaller photometric signals produced by Earth-size planets.  The procedures developed in Paper~I to model blend light curves in detail, along with those described in \\S\\ref{sec:asymmetries} to simulate spectral line asymmetries and velocity variations, are powerful analysis tools that can be of great help for validating these transiting planet candidates. The combination of the two techniques makes it possible to construct self-consistent and realistic blend scenarios and to make predictions that can then be tested quantitatively against the observations. The application to OGLE-TR-56 in this paper has provided crucial evidence of its true nature, and serves as an example of follow-up studies that may be applied to other transit candidates."
    },
    "0410/astro-ph0410247.txt": {
        "abstract": "Measurements of interstellar \\nad\\ $\\lambda \\lambda 5890, 96$ absorption lines in 18 ultraluminous infrared galaxies (ULIGs) have been combined with published \\nad\\ data, in order to reassess the dependence of galactic outflow speeds on starburst luminosity and galactic mass.  The Doppler shifts reveal outflows of relatively cool gas in 15 of 18 ULIGs with an average outflow speed at the line center of  $330 \\pm 100$\\kms. The relation between outflow speed and star formation rate (SFR), defined by the distribution's upper envelope over four orders of magnitude in SFR, demonstrates that winds from more luminous starbursts accelerate interstellar gas to higher speeds as  roughly $v \\propto SFR^{0.35}$. This result is surprising since, in the traditional model for starburst-driven winds, these relatively cool gas clouds are accelerated by the ram pressure of a hot, supernova-heated wind that exhibits weak (if any) \\x temperature variation with increasing galactic mass. The lack of evidence for much hotter winds is partly a sensitivity issue; but the  \\nad\\ velocities in ultraluminous starbursts actually are consistent with acceleration by the tepid wind indicating a hotter component is unlikely to dominate the momentum flux. The \\nad\\ velocities in the dwarf starburst winds do not reach the terminal velocity of a hot wind at the measured temperature of $kT \\sim 0.73$~keV  -- a result which could be interpreted simply as evidence that the hot superbubbles are too confined in dwarf starbursts to generate a free-flowing wind. A dynamically-motivated scenario, however, is that the dwarf starburst winds simply lack enough momentum to accelerate the clouds to the velocity of the hot wind. Among the subsample of starbursts with well-constrained dynamical masses, the terminal outflow velocities are found to always approach the galactic escape velocity. Motivated by a similar scaling relation for {\\it stellar} winds, the galactic Eddington luminosity for dusty starbursts is shown to be within the range measured for ULIGs. If radiation pressure on dust grains, coupled to the cool wind, is indeed important for galactic wind dynamics,  then feedback will be stronger in massive galaxies than previously thought helping shape the high-mass end of the galaxy luminosity function. Regardless of the nature of the acceleration mechanism in ULIGs, the mass flux of cool gas estimated from these data demonstrates that starburst-driven winds transport significant gas during the assembly stage of field elliptical galaxies -- a factor which helps explain the rapid decline in SFR in these systems inferred from elemental abundance ratios. ",
        "introduction": "Galactic-scale gaseous outflows  are ubiquitous in starburst galaxies at all cosmic epochs (Heckman, Armus, \\& Miley 1990; Pettini \\et 2001; Shapley \\et 2003). Powered by supernova explosions, perhaps with additional help from AGN in very luminous galaxies, these galactic winds transport  heavy elements, interstellar dust, and energy into galaxy halos and the intergalactic medium.  The significance for galaxy evolution depends on details of the coupling between the available energy, both radiative and mechanical, and the interstellar gas.  Observations, which describe this interplay empirically, illuminate the dominant physical processes and provide quantitative constraints for galaxy evolution models.  This paper directly addresses feedback at the high mass end of the galaxy luminosity function via measurements of outflow speeds in ultraluminous starbursts. These new observations of Na I and K I (ionization potential of 5.1~eV and 4.3~eV respectively) interstellar absorption reveal relatively cold gas clouds advected into the outflow at the shear interface between the hot wind and the quiescent gas disk. Material evaporated from these clouds could be the source of the mass-loading required to explain the large X-ray surface brightnesses of galactic winds (e.g. Suchkov \\et 1994; Suchkov \\et 1996; Strickland \\et 2002; Martin, Kobulnicky, \\& Heckman 2002). Unlike emission lines, absorption lines leave no ambiguity between infall and outflow and are sensitive to extended, low density gas. The results reported here complete a larger effort to describe the low-ionization absorption kinematics over four orders of magnitude in star formation rate (Heckman \\et 2000, luminous infrared galaxies (LIGs);  Schwartz \\& Martin 2004, dwarf starbursts). The observations have little overlap with the Rupke \\et (2002) sample of ULIGs due to different selection criteria (see \\S 2).  Feedback is expected to have a large impact on dwarf galaxies owing to their shallow gravitational potential.  It can explain the absence of young galaxies with very low mass, $V_{rot} \\sles\\ 10 - 20$\\kms, after reionization (Barkana \\& Loeb 1999), the slope of the galaxy luminosity function (Dekel \\& Silk 1986; Dekel \\& Woo 2003; Benson \\et 2003), and the mass-metallicity relation among galaxies (Larson 1974). Ablation of the gas in dwarf galaxies by winds from nearby galaxies may help flatten the faint-end of the luminosity function (Scannapieco, Thacker, \\& Davis 2001).   Observations strongly support the idea that a larger fraction of the metals are removed from dwarf galaxies than massive galaxies. Metallicity measurements from X-ray spectra indicate the hot wind carries the heavy elements produced by the starburst population (Martin \\et 2002).  Temperature estimates for the hot gas, i.e. $kT_x$, vary little with galactic mass indicating a critical galaxy mass, $V_c \\sim 130$\\kms, below which most of the hot wind escapes (Martin 1999). Retention of a larger fraction of the metals in more massive galaxies would produce an increasing effective yield with galaxy mass up to this scale where the stellar yield is recovered (Larson 1974; Dekel \\& Silk 1986; Vader 1987).  Although the X-ray data have some significant shortcomings -- including a small number of observed objects and a temperature biased towards the densest regions of the hot plasma, the mass-metallicity relation has been measured locally for a complete sample of galaxies. The metallicities increase steeply with mass over the range from $10^{8.5} - 10^{10.5}$\\msun $h_{70}^{-2}$ but flatten at $10^{10.5}$\\msun $h_{70}^{-2}$ (Tremonti \\et 2004). The mass scale of the flattening appears to be consistent with the inferred mass scale of metal retention (Garnett 2002).  Despite the efficiency of metal ejection, it is not clear whether the total mass loss rates which are of the order of the SFR are sufficient to explain the measured faint-end slope of the galaxy luminosity function, which is flatter than  that of the halo mass function (Somerville \\& Primack 1999).   This paper finds an increase in the velocities of cold clouds with galactic mass suggesting feedback may play a more prominent role in massive galaxies than previously thought, an idea which could be important for several outstanding problems. First, wind velocities are quite important in determining how much of the IGM is polluted (Aguirre  2001). Fast winds from quasars may well deposit more energy into the IGM than starburst winds (Scannepieco \\& Oh 2004), but they seem unlikely to dominate the metal transport. Second, the galaxy luminosity function cuts off more steeply than the theoretical halo mass function at bright luminosities, and a feedback mechanism physically distinct  from that in dwarf galaxies appears to be required to explain both ends of the luminosity function. The mass scale separating these forms of feedback shows up not only in the mass-metallicity relation and \\x temperatures but also characterizes the galaxy mass scale $\\sim 3 \\times 10^{10}$\\msun\\ where (1) surface-brightness stops rising with luminosity and mass and (2) less massive galaxies have stellar populations weighted towards younger ages (Blanton \\et 2003; Kauffmann \\et 2003a,b). Any feedback processes invoked to explain these relations must do so without disrupting the Tully-Fisher and Faber-Jackson relations between velocity, i.e. the depth of the gravitational potential, and stellar mass, $V \\propto M_*^{1/4}$ (e.g. Bernardi \\et 2003), which show no obvious change at this scale (Zwaan \\et 1995; Sprayberry \\et 1995; Dale \\et 1999).  Sample selection and the observational strategy are described in \\S 2. The kinematics and column densities of the nuclear ULIG spectra are presented in \\S 3. Section~4 develops scaling relations between outflow properties and galaxy parameters, including the dynamical state of the merger. The results are shown to challenge the standard model of supernova-driven winds in \\S 5, and the role of radiative acceleration is discussed. \\S 6 summarizes the main results. The second paper in this series (Paper II) describes the surprisingly large spatial extent of \\nad\\ absorption, compares kinematics of \\nad\\ and \\Ha lines, and discusses the outflowing mass flux in detail. ",
        "conclusions": "Doppler shifts measured from interstellar absorption lines in starburst spectra reveal cool gasesous outflows.  The fraction of starbursts presenting outflows increases with starburst luminosity reflecting either the brief duty cycle of the ultraluminous phase or the large solid angle subtended by these winds. The cool outflow does not entirely cover the optical nucleus indicating clouds or shells are a better description than a continuous fluid. The covering fraction approaches unity in the galaxies with the strongest winds (i.e. largest equivalent width), and these merging systems are usually found to be near perigalactic passage. More widely separated systems and single-nuclei systems generally have lower covering factors and line strengths.  These results lead to the conclusions that (1) much of the mass loading occurs very near perigalactic passage and (2) the mass loading is gradually shut down as the nuclei finish merging.  The mass entrainment rate of cold clouds is quite large, possibly comparable  to the SFR, and is discussed in more detail in Paper II.   Galaxies with higher SFRs accelerate the absorbing clouds to higher velocities. Projection effects introduce a large range in measured velocities at a given SFR, but the maximum outflow velocity rises roughly as $SFR^{1/3}$. This trend challenges the hypothesis that the clouds are accelerated by the hot winds, whose emission-measure weighted temperatures appear to vary little with circular velocity.  The simplest explanation is that the clouds in the ultraluminous starbursts reach terminal velocities comparable to the velocity of the hot winds, and that the low \\nad\\ wind velocities in dwarf starbursts reflect a shortage of momentum in dwarf starburst winds.  However, in order for momentum-driven winds to explain the intriguing similarity between the  terminal cloud velocity and the galactic escape velocity, the rate of momentum injection must be tightly correlated with the depth of the gravitational potential.  Radiatively-driven winds naturally provide this connection via a galactic equivalent to the stellar Eddington luminosity, while rather ad hoc recipes between the SFR and dark halo mass must be cooked up for supernova-driven winds.  A highly simplified description of how dusty clouds would be accelerated by the starburst continuum luminosity demonstrates that ULIG luminosities are sufficient to accelerate significant columns of cool gas to the observed \\nad\\ velocities.   These results draw attention to the importance of feedback in very massive galaxies. Preliminary estimates of the mass loss rates suggest a significant fraction of the ISM participates in the cool outflows. This result is surprising in the context of energy-driven winds because a large fraction of the supernova energy is expected to be radiated away in the dense molecular ISM of an ULIG.  Much as supernova feedback is most effective in dwarf galaxies where hot gas can escape from the halo, radiative feedback is naturally strongest in the densest, most luminous galaxies.  Hence, radiative feedback may help explain the differences between the shape of the halo mass function and the galaxy luminosity function at the high-mass end. The dynamical properties of the stars in ULIGs do appear to be consistent with those of field ellipticals (Genzel \\et 2001), so it is also appealing to associate the winds seen in ULIGs with the event that shuts the star formation down quickly enough to create high ratios of alpha elements to iron in elliptical galaxies.  Filling in \\fig~\\ref{fig:v14_sfr} and \\fig~\\ref{fig:vterm_model} with more observations will help clarify some of these isssues. Also, since radiatively-driven winds do not require a hot phase,  X-ray observations of these ULIGs may distinguish acceleration mechanisms. Spitzer and GALEX observations will be instrumental for constraining the amount of extraplanar dust in starbursts and its composition. A few of the most luminous objects present velocities too high to explain with starburst winds, so it will be possible to explore the nature of outflows in warm ULIGs during the SB/AGN transition. More detailed dynamical modeling of the survivability and fate of the cool clouds as well as a more realistic treatment of  radiation pressure will allow a more detailed comparison to the data (e.g. line widths, spatial gradients, etc.) and thereby a more stringent test of the ideas presented here."
    },
    "0410/astro-ph0410631_arXiv.txt": {
        "abstract": "The real-space power spectrum of $L_*$ galaxies measured from the 2dF Galaxy Redshift Survey (2dFGRS) is presented. Markov-Chain Monte-Carlo (MCMC) sampling was used to fit radial and angular modes resulting from a Spherical Harmonics decomposition of the 2dFGRS overdensity field (described in Percival et~al. 2004) with 16 real-space power spectrum values and linear redshift-space distortion parameter $\\beta(L_*,0)$. The recovered marginalised band-powers are compared to previous estimates of galaxy power spectra. Additionally, we provide a simple model for the 17 dimensional likelihood hyper-surface in order to allow the likelihood to be quickly estimated given a set of model band-powers and $\\beta(L_*,0)$. The likelihood surface is not well approximated by a multi-variate Gaussian distribution with model-independent covariances. Instead, a model is presented in which the distribution of each band-power has a Gaussian distribution in a combination of the band-power and its logarithm. The relative contribution of each component was determined by fitting the MCMC output. Using these distributions, we demonstrate how the likelihood of a given cosmological model can be quickly and accurately estimated, and use a simple set of models to compare estimated likelihoods with likelihoods calculated using the full Spherical Harmonics procedure. All of the data are made publically available (from {\\tt http://www.roe.ac.uk/{\\tt\\char'176}wjp/}), enabling the Spherical Harmonics decomposition of the 2dFGRS of Percival et al. (2004) to be easily used as a cosmological constraint. ",
        "introduction": "In addition to a wealth of information about the physics of galaxy formation, galaxy surveys also provide a statistical measure of the large-scale primordial density perturbations. If, as is theoretically expected, the density field on large-scales is well-approximated by a random-phase Gaussian distribution, then the power spectrum provides a statistically complete description of the field. Given a galaxy redshift survey, the most direct route to an approximation to the power spectrum is to decompose the overdensity distribution into Fourier modes (e.g. Feldman, Kaiser \\& Peacock 1994). However, systematic observed distortions in the radial direction caused by the coherent peculiar velocities of galaxies are difficult to incorporate into a model of a Fourier decomposition.  Consequently, methods to decompose the galaxy overdensity field into 3-dimensional bases consisting of radial and angular functionals have been developed by a number of groups (Heavens \\& Taylor 1995; Fisher et~al. 1994;1995; Hamilton, Tegmark \\& Padmanabhan 2000). By splitting the basis into radial and angular components, radial redshift-space distortions can be readily included in the analysis method. Recently, the 2dFGRS and SDSS galaxy redshift surveys have both been analysed in this way (Percival et~al. 2004:P04, Tegmark et ~al. 2004) resulting in tight constraints on both the power spectrum shape and the redshift-space distortions. Other analyses of 2-point statistics measured from these surveys include Percival et~al. (2001) and Cole et~al. (2004) who decomposed the 2dFGRS into Fourier modes and Pope et~al. (2004) who undertook a Karhunen-Lo\\`{e}ve eigenmode analysis of the SDSS based on the correlation function of counts-in-cells.  In this paper we build upon the analysis of P04 who decomposed the 2dFGRS overdensity field into an orthonormal basis of spherical harmonics and spherical Bessel functions. Throughout we refer to such a basis as a Spherical Harmonics basis. In P04, cosmological models were tested against combinations of the radial and angular modes calculated for the decomposed density field. The procedure for calculating the likelihood of each model required inversion of a large covariance matrix ($\\sim1800\\times1800$) and was computationally expensive. It would obviously be advantageous if the same Likelihood surface could be recovered without having to invert a large matrix for each cosmological model.  The Markov-chain Monte-Carlo (MCMC) technique provides a method for mapping a likelihood surface, providing both marginalised parameters and information about the form of the surface. In this paper, we use the MCMC technique to map the true likelihood surface as a function of $16$ power spectrum band-powers and $\\beta(L_*,0)$. Here $\\beta(L_*,0)\\equiv\\Omega_m^{0.6}/b(L_*,0)$ is a measure of the increased fluctuation amplitude for present day $L_*$ galaxies caused by the linear movement of matter onto density peaks and out from voids (Kaiser 1987). $\\Omega_m$ is the present day matter density and $b(L_*,0)$ is the present day bias of $L_*$ galaxies. The resulting sampling of the likelihood surface provides an estimate of the real-space galaxy power spectrum through the marginalised power spectrum band-powers (presented in Section~\\ref{sec:real_pk}).  In addition, we fit models to the distribution of each band-power, marginalised over the other parameters (Section~\\ref{sec:like_fit}). The distributions are assumed to be Gaussian in a combination of the band-power and its logarithm. Simply assuming a Gaussian distribution in each parameter and model-independent covariances could potentially bias the likelihood estimates. This is directly analogous to the problem of compressing CMB maps into angular band-powers, and the use of those band-powers to estimate the true likelihood surface (e.g. Bond Jaffe \\& Knox 2000). The model of the likelihood surface can be used to quickly estimate the likelihood of any theoretical power spectrum and $\\beta(L_*,0)$. In Section~\\ref{sec:simple_models} the likelihood surface estimated for a simple set of cosmological models from these distributions is shown to be in excellent agreement with the original likelihood surface. ",
        "conclusions": "A new method using the MCMC technique has been developed to recover the real-space power spectrum from a decomposition of a galaxy overdensity field and has been applied to the 2dFGRS Spherical Harmonics analysis presented in P04. The recovered 2dFGRS $L_*$ power spectrum was presented as $16$ marginalised power spectrum band-powers. The MCMC technique was also used to obtain information about the shape of the likelihood surface. Using this information, the likelihood surface was fitted with a simple model, providing a quick method for estimating the likelihood of different cosmological models.  The advantage of the MCMC method over previous deconvolution routines lies in providing a map of the likelihood surface. Maximum likelihood techniques have previously been used to recover the real-space power spectrum in a series of band-powers (Ballinger, Heavens \\& Taylor 1995), using a standard routine to iterate to the likelihood maxima. However, such a step-wise maximum likelihood method only partially explores the likelihood surface giving poorly determined conditional errors (Ballinger et~al. 1995; Tadros et~al. 1999). The MCMC technique provides a significant advance over such methods, as it enables a more complete exploration of the likelihood surface to be quickly undertaken.  An alternative to the step-wise approach to determining ML band-powers, is to use quadratic band-power estimators (Hamilton 1997a;b) favoured by Tegmark et~al. (2002; 2004). The band-powers that we model have no window functions, and therefore must correspond to the linear combinations of the quadratic estimators given by the inverse of the Fisher information matrix (see, for example, section 3.5 of Tegmark et~al. 2002). In this approach, the resulting band-powers are assumed to have a Gaussian likelihood with model-independent covariances, which, as we argued in Section~\\ref{sec:like_true}, for an all-sky noiseless survey must lead to a different form for the likelihood compared with the true likelihood given by Eq.~\\ref{eq:like_allsky}. This is consistent with the MCMC analysis presented in this paper which suggests that the likelihood distribution around the band-powers is not well described by a multi-variate Gaussian distribution. Either the covariance matrix must change its form, or equivalently the shape of the function assumed for the likelihood surface must change. The MCMC technique is therefore advantageous over the quadratic band-power method as it provides the information required to allow the likelihood surface to be accurately modelled beyond the likelihood peak.  Another factor that makes the MCMC routine particularly well suited to this application is the nature of the likelihood surface. No evidence was found for multiple likelihood maxima, which would slow the MCMC mapping. Additionally, more simple techniques would fail because the surface does not follow an obvious form -- for instance it is close to, but not exactly a multi-variate Gaussian so a simple fit of a multi-variate Gaussian would be inappropriate. Fitting to the likelihood surface has been shown to provide an accurate estimate of the combined likelihoods for a number of cosmological models covering a range of power spectrum shapes. This gives us confidence that it can be extended to cover a wider range of cosmological models.  From the MCMC chains, we have estimated marginalised band-powers, and have compared these with measurements of the real-space power spectrum from the SDSS, and with other measurements from the 2dFGRS. Excellent agreement is observed and the basic shape of the power spectrum now seems to be well constrained. We have not directly compared with the real-space power spectrum of the IRAS 1.2Jy galaxies (Ballinger et~al. 1995) or with the real-space power spectrum of the PSCz galaxies estimated by Tadros et~al. (1999), in order not to further clutter Fig.~\\ref{fig:real_pk_cmpr}. However, it is worth noting that the general shape of the recovered power spectra are similar. We have compared with the deprojected APM data of Efstathiou \\& Moody (2001) and Padilla \\& Baugh (2003), although the 2dFGRS and SDSS data represent a significant improvement over these data.  In addition to the band-power measurements we simultaneously fit the linear redshift-space distortions. A best-fit value of $\\beta(L_*,0)=0.55$ was recovered with 1-$\\sigma$ confidence region given by $0.46<\\beta(L_*,0)<0.64$. This compares very well with the value of $\\beta(L_*,0)=0.58\\pm0.08$ recovered in P04 marginalising over the basic set of cosmological models considered. Note that from the approximation to the likelihood we recover the same parameter value if we consider the range of models chosen, and their associated priors on the shape of the power spectrum.  The marginalised band-power and $\\beta(L_*,0)$ measurements were presented in Table~\\ref{tab:pk}. These should not be confused with the parameters of the fit to the likelihood surface that were presented in Table~\\ref{tab:like_fit} and associated covariance matrix. All of these data are publically available from {\\tt http://www.roe.ac.uk/{\\tt\\char'176}wjp/}, where simple segments of code demonstrating how the likelihood can be quickly calculated are also provided. It should therefore be possible to easily combine the cosmological constraints provided by the Spherical Harmonics analysis of P04 with other cosmological probes.  Obviously, these constraints should not be combined with other analyses of the 2dFGRS such as provided by Percival et~al. (2001) or Cole et~al. (2004) as these results are strongly inter-dependent. The analysis presented by Cole et~al. (2004) uses more galaxies and therefore provides tighter constraints on the power spectrum shape than the Spherical Harmonics analysis considered here and should be used in preference to this work if the power spectrum shape is the important consideration. The advantage of the Spherical Harmonics analysis is that, although it provides weaker information on $P(k)$ shape, it provides additional correlated information about $\\Omega_m^{0.6}\\sigma_8$ from fitting the redshift-space distortions.  An example application where using Spherical Harmonics rather than Fourier likelihoods may be advantageous is to provide constraints on the neutrino mass density. Here, the amplitude of the low-redshift matter power spectrum provides a significant constraint, and has previously been determined by converting a redshift-space power spectrum amplitude to $\\sigma_8$ using $\\beta$ and galaxy bias $b$ calculated from other analyses (e.g. Spergel et al. 2003). The problem with such an approach is that $\\beta$ and $b$ are often measured on different scales, using different analysis techniques and model dependencies. The Spherical Harmonics method instead sets a constraint on $\\sigma_8$ (through $\\Omega_m^{0.6}\\sigma_8$) in a self-consistent way. Combining the Spherical Harmonics likelihoods with CMB likelihoods should therefore provide a more robust determination of the neutrino mass density."
    },
    "0410/astro-ph0410407_arXiv.txt": {
        "abstract": "In the context of the standard model of particle physics, there is a definite upper limit to the density of stable compact stars. However, if there is a deeper layer of constituents, below that of quarks and leptons, stability may be re-established far beyond this limiting density and a new class of compact stars could exist. These objects would cause gravitational lensing of white dwarfs and gamma-ray bursts, which might be observable as a diffraction pattern in the spectrum. Such observations could provide means for obtaining new clues about the fundamental particles and the origin of cold dark matter. ",
        "introduction": "The different types of compact objects traditionally considered in astrophysics are white dwarfs, neutron stars (including quark and hybrid stars), and black holes. The first two classes are supported by Fermi degeneracy pressure from their constituent particles. For white dwarfs, electrons provide the pressure counterbalancing gravity. In neutron stars, the neutrons play this role. For black holes, the degeneracy pressure is overcome by gravity and the object collapses to a singularity, or at least to the Planck scale ($\\rho \\sim 10^{93}$~g/cm$^3$). For a recent review of neutron stars, hybrid stars, and quark stars, see, {\\it e.g.},~\\cite{Weber:04} and references therein.  The distinct classes of degenerate compact stars originate directly from the properties of gravity, as was made clear by a theorem of Wheeler and collaborators in the mid 1960s~\\cite{Harrison:65}. This theorem states that for the solutions to the stellar structure equations, whether Newtonian or relativistic, there is a change in stability of one radial mode of vibration whenever the mass reaches a local maximum or minimum as a function of the central density. The theorem assures that distinct classes of stars, such as white dwarfs and neutron stars, are separated in central density by a region in which there are no stable configurations.  In the standard model of quarks and leptons (SM), the theory of the strong interaction between quarks and gluons predicts that with increasing energy and density, the coupling between quarks asymptotically fades away~\\cite{Gross:73,Politzer:73}. As a consequence of this so-called asymptotic freedom, matter is expected to behave as a gas of free fermions at sufficiently high densities. This puts a definite upper limit to the density of stable compact stars, since the solutions to the stellar equations end up in a never-ending sequence of unstable configurations, with increasing central density. Thus, in the light of the standard model, the densest stars likely to exist are neutron stars, quark stars, or the slightly more dense hybrid stars~\\cite{Gerlach:68,Glendenning:00,Schertler:00}. However, if there is a deeper layer of constituents, below that of quarks and leptons, the possibility of a new class of compact stars opens up~\\cite{Hansson:04}.  Though being a quantitatively successful theory, the SM consists of a large number of exogenous {\\it ad hoc} rules and parameters, which were introduced solely to fit the experimental data. The SM provides no explanation for the deeper meaning of these rules. At a closer look, however, the SM seems to be full of hints to its deeper background. By considering these rules from a historical point of view, a ``simple\" and appealing explanation is {\\it compositeness}~\\cite{Fredriksson:03}, {\\it i.e.}, that the quarks, leptons, and gauge bosons are composite particles, built out of more elementary {\\it preons}~\\cite{dSouza:92}. Preons provide natural explanations for the particle families of the SM and phenomena such as neutrino oscillations, mixing of the weak gauge bosons, and quarks of different flavour.  Over the last decades, many papers have been written about preons, but so far there are no direct evidence for (or against!) their existence. In the late 1970s, a number of consistency conditions were formulated by 't~Hooft~\\cite{tHooft:79}. In the same work, a vector-like non-Abelian $SU(3)$ gauge group was considered, but no solution to the consistency conditions was found. Later, it was shown that with another choice for the gauge group and the flavour structure of preons, {\\it e.g.}, three different preon flavours, the consistency conditions are satisfied~\\cite{Barbieri:80}. For a more detailed discussion of preon models, see~\\cite{Fredriksson:03,dSouza:92,PreonTrinity:02} and references therein.  Not all clues favour preon models, but the existence of preons is still an open question, and as a consequence so is the question whether a new class of compact stars exists or not. This paper is based on the ideas and results presented in~\\cite{Hansson:04}. Assuming that quarks and leptons are composite particles, built out of more elementary preons, I will: \\newcounter{Lcount} \\begin{list}{\\Roman{Lcount}.}{\\usecounter{Lcount}} \\item{Give an estimate for the mass and radius of stars composed of preons.} \\item{Show that for some particular equations of state, stable solutions to the general relativistic stellar equations do exist, with densities far beyond the maximum density in stars composed of quarks and leptons.} \\item{Briefly discuss some potential astrophysical consequences and how these objects could be observed. Herein lies the potential importance of this qualitative speculation, since these objects are candidates for cold dark matter and could be found as, {\\it e.g.}, gravitational ``femtolenses\".} \\end{list} ",
        "conclusions": "If there is a deeper layer of fermionic constituents (preons), below that of quarks and leptons, a new class of stable compact stars could exist. By fitting a simple equation of state for fermions to the minimum energy density of an electron, the maximum mass for stars composed of preons can be estimated to $\\sim 10^{2}$~Earth masses and the maximum radius to $\\sim 1$~m. The minimum central density is of the order of $\\sim 10^{23}$~g/cm$^3$. Preon stars with a maximum mass of $\\sim 10^{-1}$~Earth masses and radius $\\sim 1$~mm could have been formed by the primordial density fluctuations in the early universe. By utilizing gamma-ray bursts, or white dwarfs in the large magellanic cloud as light sources, an intermediate preon star would produce a diffraction signal in the spectrum, which might be observable. An observation of an object as dense as a preon star would be a direct evidence for a new state of matter, which is not composed of quarks and leptons. Due to the need for observational clues in the cold dark matter sector, this could prove compositeness plausible, without much dedicated effort. This approach might complement direct tests of preon models~\\cite{PreonTrinity:02} performed at particle accelerators."
    },
    "0410/astro-ph0410294_arXiv.txt": {
        "abstract": "A new nonlinear electromagnetic mode in a magnetized plasma is predicted. Its existence depends on the interaction of an intense circularly polarized electromagnetic wave with a plasma, where quantum electrodynamical photon--photon scattering is taken into account. This scattering gives rise to a new coupling between the matter and the radiation.  Specifically, we consider an electron--positron plasma, and show that the propagation of the new mode is admitted. It could be of significance in pulsar magnetospheres, and result in energy transport between the pulsar poles. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410588_arXiv.txt": {
        "abstract": "We present FUSE observations of the \\ion{He}{2} Ly$\\alpha$ forest in the redshift range $2.3 < z < 2.7$ towards HS~1700+6416. Between October 2002 and February 2003, the brightness of the QSO increased by a factor 2. Therefore, with an exposure time of 203\\,ks during orbital night, the quality of the resulting spectrum is comparable to the HE~2347-4342 data. This second line of sight with a resolved \\ion{He}{2} Lyman alpha forest reveals a similar variation of several orders of magnitude of the column density ratio $\\eta = N($\\ion{He}{2}$)/N($\\ion{H}{1}$)$ and confirms the results of previous studies. The well-known metal line spectrum of HS~1700+6416 permits to examine the influence of metal line absorption on the \\ion{He}{2} column densities. ",
        "introduction": "Resolving the \\ion{He}{2} Ly$\\alpha$ forest in HE~2347-4342 with FUSE \\citep{krissetal01} has led to a detailed picture of the ionization if the IGM. The cosmic UV background which governs the \\ion{He}{2}/\\ion{H}{1} column density ratio was found to fluctuate by at least two orders of magnitude, with the highest values observed in the low density part (voids) of the Ly$\\alpha$ forest \\citep{shulletal04}. We present here FUSE observations of HS~1700+6416 ($z_{\\mathrm{em}} = 2.73$), a second line of sight of comparable quality. ",
        "conclusions": ""
    },
    "0410/astro-ph0410280_arXiv.txt": {
        "abstract": "We perform a blind multi-component analysis of the WMAP 1 year foreground cleaned maps using SMICA (Spectral Matching Independent Component Analysis). We provide a new estimate of the CMB power spectrum as well as the amplitude of the CMB anisotropies across frequency channels. We show that the CMB anisotropies are compatible with temperature fluctuations as expected from the standard paradigm. The analysis also allows us to identify and separate a weak residual galactic emission present significantly in the Q-band outside of the Kp2 mask limits, and mainly concentrated at low galactic latitudes.  We produce a map of this residual component by Wiener filtering using estimated parameters. The level of contamination of CMB data by this component is compatible with the WMAP team estimation of foreground residual contamination. In addition, the multi-component analysis allows us to estimate jointly the power spectrum of unresolved point source emission. ",
        "introduction": "The Cosmic Microwave Background (CMB) is one of the most powerful probe of modern cosmology. The shape of the spatial power spectrum of the small temperature fluctuations depend on the cosmological parameters describing, in the frame of the standard model, the matter content, the geometry, and the evolution of the Universe \\citep{jungman96}. Since the first detection of CMB anisotropies by the COBE satellite in 1992 \\citep{DMR}, several ground--based and balloon--borne experiments have provided an accurate estimate of the power spectrum on a large range of angular scales.  The recent WMAP mission \\citep{WMAP1yr}, after one year of data acquisition, provided measurements of the power spectrum with unprecedented accuracy.  Measuring the CMB power spectrum is a difficult task. It requires a good characterization of noise contribution in the observations, as well as the subtraction of foreground astrophysical emission present at millimeter wavelengths \\citep{fb-rg99}.  Both noise and foreground emissions can significantly bias CMB power spectrum estimates if not well accounted for.  The WMAP data is the best presently available observation of the sky in the 20-90~GHz range.  A careful separation of CMB from foregrounds, however, is required in order to make the best out of these observations.  In particular, isolating CMB from foregrounds is of extreme importance for measuring accurately the angular power spectrum $C(\\ell)$ of CMB anisotropies and for cosmological interpretation.   The approach of the WMAP collaboration consists in cleaning observations from foreground contaminations for performing a CMB power spectrum estimation on the cleaned maps. More specifically, galactic foregrounds are subtracted using templates obtained from external data sets.  In addition, strong known extra-galactic point sources, as well as the region of the galactic plane, are masked prior to the analysis. The CMB spectrum is estimated using a weighted average of the cross-power spectra of cleaned observation maps in the Q, V and W channels \\citep{WMAPSp}.  This avoids biases due to detector noise, assuming that the noise of different detectors is uncorrelated.  This estimated spectrum is corrected for residual point source contamination by subtracting an estimate of the contribution of point sources in the cross power spectra. The level of residual galactic contaminations is estimated by cross-correlating the maps with foreground templates.  A multi-component approach of CMB spectral estimation has been proposed by \\cite{SM-mnras}. Additional publications on the subject give variants and details \\citep{cardoso02,patanchon03}.  The method, called SMICA (Spectral Matching Independent Component Analysis), is based on matching the cross- and auto-power spectra of the observed maps to a parametric model described by the power spectra of all the astrophysical components, their relative amplitudes in the different channels, and the noise power spectra. It is a very flexible approach: depending on available knowledge, most of the parameters may be either estimated from the data or kept at fixed values.  In particular, emission laws for all or some of the components can be estimated.  There are several benefits to using SMICA in CMB analysis.  First, SMICA is equivalent to a maximum likelihood estimation of the CMB power spectrum if the observed maps are a linear mixture of Gaussian stationary components and noise. Therefore, if there are no foregrounds but only Gaussian CMB and noise, SMICA is expected to outperform quadratic estimates.  Secondly, SMICA allows to estimate jointly the spatial power spectra \\emph{and} the amplitude of components in each channel (which are related to emission laws of components as well as calibration coefficients). In particular, the application on WMAP data allows us to check the following strong prediction of the standard model: CMB anisotropy should have a spectral emission law given, to first order, by a derivative (with respect to the temperature) of the blackbody law. \\cite{fixsen03} has measured the CMB anisotropy contribution in the COBE FIRAS data by looking for the WMAP anisotropy template in FIRAS. The author shows the compatibility of CMB anisotropy at large angular scales ($> 5^\\circ$) with temperature fluctuations in the Wien part of the spectrum. Our blind estimation of the component emission laws provides us with a unique tool for extending this result to frequencies covered by WMAP and to all measured angular scales (including the first acoustic peak) with a remarkable precision as shown below.  Finally, if there are foreground contributions in the observations, SMICA is designed to detect them and allows their separation.  It permits indeed to jointly estimate power spectra of multiple components in the data, {\\emph{to assess the significance of components}}, and eventually to separate the effect of all the emissions.  \\medskip  In this paper, we investigate the existence of residual foreground emission in foreground--cleaned WMAP data.  In particular, residual galactic emission resulting from an error in the subtraction of galactic templates can be present in the published maps. Such residuals may exist if external templates of galactic emissions, extrapolated using a physical model to WMAP frequencies, are not quite representative of the actual galactic components. Multi-component analysis with SMICA allows to check for the existence of such a residual, and to quantify the level of contamination if any, without prior information on its power spectrum or on its amplitude across channels.  In addition, the WMAP collaboration has shown that unresolved point source emissions have a non-negligible contribution to the data \\citep{WMAPSp,komatsu}. This result has been independently confirmed by \\cite{SZ-WMAP}. As will be shown, SMICA also yields a coarse estimate the level of point source emission in WMAP data.  The paper is organized as follow: section \\ref{sec:data} introduces the data used for the analysis, the simulations performed in order to check our results and the model of the WMAP observations for SMICA. The method is described in section \\ref{sec:method}. Section \\ref{sec:results} presents the different results. Finally, conclusions are provided in section \\ref{sec:conclusion} ",
        "conclusions": "\\label{sec:conclusion}  We have performed a blind multi-component analysis of the first year WMAP data.  We used the eight foreground-cleaned high frequency maps (Q1 to W4) provided by the WMAP collaboration. Our analysis uses SMICA, a maximum likelihood spectral matching method.  Three significant astrophysical components \\begin{itemize} \\item the CMB anisotropies \\item a residual galactic emission \\item unresolved point source emissions \\end{itemize} are jointly characterized in the foreground cleaned maps after masking the strongest point sources and the galactic plane using Kp2 mask. No significant thermal SZ emission is found with the present study, in agreement with the WMAP team \\citep{wmap-foreg} and \\cite{SZ-WMAP}.  Blind analysis allows us to estimate the power spectrum of CMB anisotropies as well as their amplitude across frequency channels confirming their cosmological nature. Our power spectrum estimate is in excellent agreement with the WMAP team estimate. We show that the measured emission law of CMB anisotropies at WMAP frequencies is compatible with the derivative of a blackbody, as expected for temperature fluctuations. The statistical errors on parameters related to the amplitude of the anisotropies are about 0.3\\%, and are smaller than calibration errors (0.5\\%) on dipole modulation. Conversely, if CMB anisotropies are assumed to be pure temperature fluctuations, then the estimation of CMB amplitude across channels provides a relative calibration across frequency bands at a better precision than dipole calibration.  The second estimated component, corresponding to a weak residual galactic emission, is mainly concentrated in Q-band maps outside of the Kp2 mask.  We believe that this component results mainly from spatial variations of the difference between Haslam map (used as a synchrotron template for the subtraction) and synchrotron emission at WMAP frequencies.  This component is weak compared to CMB anisotropies.  The estimated power spectrum is about $\\ell(\\ell+1)C(\\ell)/2\\pi \\approx 10-12\\mu K^2$ for $\\ell < 100$ in Q band and less than $2\\mu K^2$ in V and W band. Those estimates are compatible with WMAP team estimates of foreground contamination. Finally, much of the power from this residual large scale component disappears for galactic latitude higher than 40 degrees.  The third component corresponds to residual point sources emission. By fixing the parameters related to the emission law of point sources (we assume $(\\nu/\\nu_0)^{-2}$) and to the power spectrum (flat power spectrum), we estimate the amplitude of unresolved point source power spectrum. We find at high galactic latitude: $C(\\ell) = (11.3\\pm3.7) \\times 10^{-3}\\mu K^2$ at 41~GHz compatible with the WMAP team estimation.  We also provide an estimate of the power spectrum of point sources at high galactic latitude.  The goodness of fit of the 3-component model is excellent, except for a very small discrepancy around $\\ell\\approx 20$ which is likely to be attributed to a systematic in W1.  This inconsistency has no impact on the CMB power spectrum estimate (because errors are dominated by cosmic variance), but could affect the estimates of weak components in data."
    },
    "0410/astro-ph0410139_arXiv.txt": {
        "abstract": " ",
        "introduction": "Some of the predictions of fundamental theories of physics can only be observed on energy scales as high as the Planck scale. The availability of such high energies in the early universe and the huge separation between conventional accelerator experiments and the Planck scale has led many to turn from accelerator-based experiments to cosmological observations in order to test such theories. Inflationary cosmology \\cite{1} is one of the paradigms that may serve this purpose. There, it is assumed that quantum fluctuations of the inflaton are stretched by inflationary expansion to cosmological scales. About 60 e-folds of inflationary expansion are necessary to solve many of the puzzles of big bang cosmology but in most inflationary models the expansion is much larger. In models proposed by Linde \\cite{2} the universe has expanded by a factor of $10^{10^{12}}$. An implication of these models is, therefore, that our observable universe was of \\textit{sub-Planckian} size at the beginning of the (last) inflationary period. This suggests that inflation could act as a magnifying glass for probing the short distance structure of space-time.  A similar question had arisen concerning a possible sensitivity of black hole radiation to transplanckian physics. There, it has been found that Hawking radiation is largely immune to transplanckian effects, see e.g. \\cite% {brout-bh}. In the case of inflation, however, it has been found that the inflationary predictions for the cosmic microwave background (CMB) do possess a small and possibly even observable sensitivity to modifications of quantum field theory in the ultraviolet. To this end, various examples of ultraviolet-modified dispersion relations, some motivated by solid state analogs, have been tested for their effects on inflation, see \\cite{3,4,5,6}% . In particular, and this will be our interest here,  the ultraviolet cutoff described by a lower bound in the formal uncertainty in position, $\\Delta x_{min}$, has also been investigated for its implications in inflation, see % \\cite{7,20,10,11}.  To model the small scale structure of space through a finite minimum position uncertainty $\\Delta x_{min}$ is of interest because the corresponding modified uncertainty principle has been motivated to arise from quite general quantum gravity arguments as well as from string theory, see e.g. \\cite{8,9,akucr,KMM}. In fact, any theory with this type of ultraviolet cutoff can be written, equivalently, as a continuum theory and as a lattice theory, see \\cite{akl}. While in the continuum formulation the theory displays unbroken external symmetries, the theory's ultraviolet regularity is displayed in its lattice formulation.  Indeed, it has been found that inflationary predictions for the CMB are sensitive to the natural ultraviolet cutoff if the cutoff is modelled through a finite minimum uncertainty in positions, $\\Delta x_{min}$. The magnitude by which the cutoff affects the predicted scalar and tensor spectra in the CMB was found to depend crucially on the initial conditions when a mode's evolution begins, which is when its proper wave length is the minimum length. These initial conditions determine how close the modes' state is to the adiabatic vacuum during the period of adiabatic evolution before the mode crosses the Hubble horizon. If the modes are in the adiabatic vacuum during the phase of adiabatic evolution then the effects of Planck scale physics on inflationary predictions should be no bigger than of the order of $\\sigma ^{2}$, see \\cite{20}, where: \\begin{equation} \\sigma =\\frac{\\Delta x_{min}}{L_{Hubble}} \\end{equation}% Here, $L_{Hubble}$ is the Hubble length during inflation. Thus, we have approximately $\\sigma \\approx 10^{-3}$ if the cutoff length, $\\Delta x_{min}$% , is at the string scale and $\\sigma \\approx 10^{-5}$ if the cutoff length is at the Planck scale. In principle, however, the cutoff can lead to arbitrarily large effects, namely if the modes' state during the adiabatic phase differs strongly from the adiabatic vacuum (see also \\cite{pad}). In this case, the modulus of the mode functions oscillates at horizon crossing and these oscillations translate into characteristic oscillations in the CMB spectra. This possibility is restricted, however, by the need to keep the back-reaction small \\cite{tanaka}. Interestingly, Easther \\textit{et.al. }% \\cite{10,11} found that this constraint still allows nontrivial vacua with effects as large as of order $\\sigma $. Effects of this magnitude might reach the threshold of observability.  So far, initial conditions have been proposed based on analyticity arguments % \\cite{20} and based on similarity to the Bunch Davies vacuum \\cite{10,11}. A further suggestion is to minimize the field uncertainties \\cite% {danielsson,greene3}. Still, however, the crucial question how to determine initial conditions for the new comoving modes that are continually being created during an expansion has not been conclusively answered. The problem is of course equivalent to identifying the vacuum state.  Here, we address this problem by reconsidering how the vacuum state is usually identified within inflationary QFT without a minimum length. Namely, the usual strategy is to make use of the fact that the action can be rewritten so as to resemble the familiar action of a field on Minkowski space with time-dependent mass term. When quantizing, one then chooses the vacuum as one does for Minkowski space theories. We will find that this method is no longer reliable when there is a minimum length. The reason is that the reformulation of the action requires the neglect of a boundary term that ceases to be a boundary term once the minimum length is introduced. We find that the differences are small but noticeable both in the initial conditions and in the evolution equations. {This shows that in any approach to introducing a minimum length into inflation this will have to be taken into account: reformulations of an action that appear to be harmless due to neglect of a boundary term can lead to an unintended modification of the theory.}  To see how this phenomenon can arise, let us recall that the particular model of a natural ultraviolet cutoff that we are considering is described by quantum mechanical uncertainty relations with correction terms in the ultraviolet, of the form \\begin{equation} \\Delta x \\Delta p \\ge \\frac{\\hbar}{2}\\left(1+\\beta~(\\Delta p)^2 + ...\\right) \\label{1a} \\end{equation} where $\\beta >0$ is a positive constant. In the simplest case, such an uncertainty relation arises from the modified commutation relation: \\begin{equation} \\lbrack \\mathbf{X},\\mathbf{P}]=i\\hbar (\\mathbf{1}+\\beta ~\\mathbf{P}^{2}) \\label{1} \\end{equation} It is not difficult to show that the uncertainty relation then implies a finite lower bound to the position uncertainty $\\Delta x$: \\begin{equation} \\Delta x_{min} = \\hbar \\sqrt{\\beta} \\end{equation} By choosing $\\beta$ appropriately we obtain a cutoff at the string or at the Planck scale. This type of ultraviolet cutoff was introduced into quantum field theory in \\cite{akf} and then into inflationary cosmology in \\cite{7}.  It is clear that similar to quantization, which changes the commutativity properties and therefore comes with a well-known ordering ambiguity, the introduction of a minimum length through an equation such as Eq. (\\ref{1}) changes the commutativity properties and therefore comes with an ordering ambiguity. In principle, of course, ordering ambiguities can be of arbitrary magnitude. For example, a classical system is unchanged by adding terms of the form $(xp-px)f(x,p)$ to its Hamiltonian $H$. When promoting  $x$ and $p$ to operators the resulting terms become proportional to $\\hbar $ and could be arbitrarily large and significant to the evolution. In the case we consider here, normally vanishing terms of the form $(xp-px-i\\hbar )g(x,p)$ similarly become nonzero when $\\beta \\neq 0$. Here, the new Hamiltonian is determined only up to terms that vanish when setting the minimum length to zero. Those terms can be arbitrarily large and, in principle, only experiments could decide which choice is correct. This is to be expected in any approach to introducing some form of a natural minimum length.  Of course, in the case of quantization it has proven to be a very reliable strategy to adopt the minimalist approach to resolving the ordering ambiguity: write the Hamiltonian in its most simple and symmetric form and leave it unchanged when introducing $\\hbar$, i.e. do not introduce terms by hand. The same minimalist approach has tacitly been applied in the literature when the minimum length uncertainty relation has been used in inflationary quantum field theory. We will now review this procedure, thereby uncovering potential pitfalls with implications for the determination of the vacuum. ",
        "conclusions": "While the framework of quantum field theory is well-tested down to distances of about $10^{-18}m$, it is generally expected that there are corrections due to quantum gravity when approaching the Planck length of about $10^{-35}m $ which may well constitute a fundamental smallest length in nature.  If, therefore, there exists a finite minimum wavelength then, during inflation, comoving modes are continually being created. Initially, a new comoving mode will evolve under the influence of Planck scale effects but it is clear that at late times a comoving mode's equation of motion will reduce to the usual low-energy mode equation, namely when the mode's proper wavelength becomes much larger than the minimum length. Thus, as was pointed out in \\cite{sta}, effects of the Planck scale can propagate into the observable low energy realm essentially only by selecting a solution of the mode equation which at late times differs from the usually assumed solution for the usual mode equation.  This suggests a simple technique for exploring possible effects that Planck scale physics could have on inflationary predictions for the CMB. Assume that standard quantum field theory holds \\textit{unchanged }down to the minimum wavelength where modes are being created. Then, consider a variety of possible initial conditions for the newly created modes by applying candidate criteria for identifying the vacuum state. It is clear that in the time translation invariant de Sitter case all effects reduce to merely an overall re-normalization of the flat spectrum (if each mode's initial condition is chosen by applying the same criterion). When the Hubble parameter varies, however, then modes oscillate a variable number of times before crossing the horizon. Thus, generically, a mode's amplitude will be alternatingly large and small when crossing the horizon. This can lead to potentially observable characteristic oscillations in the spectrum \\cite% {greene3}. In this approach, quantum field theory is implicitly assumed to hold unchanged down to the Planck scale and therefore Planck scale physics is modelled so as to affect the predictions of inflation, for any given evolution of the scale factor $a(\\tau )$, merely through the initial conditions. In any realistic model, of course, the quantum field theoretic mode equations will be modified when approaching the Planck scale. This too will have an effect on the number of oscillations that a mode undergoes before horizon crossing and it will therefore contribute to the predicted oscillations of the CMB spectra.  Here, we considered a concrete model for how quantum field theory is modified when approaching the Planck length, namely by introducing the minimum length uncertainty principle. The equations of motion then indeed became modified at scales close to the Planck scale. As was shown in \\cite% {20,10,11}, the inflationary predictions for the CMB are to some extent affected, possibly leading to observable oscillations in the fluctuation spectra. However, while the equations of motion are known, the details of the predictions still significantly depend on precisely which initial condition is chosen, i.e. on the identification of the vacuum.  As yet, it is not fully understood in any model how Planck scale physics determines the initial conditions of modes as they are being created, i.e. when their proper wave length is the minimum length. Within our model of spacetime, in which there is a minimum length uncertainty relation, the problem of determining the initial conditions for new comoving modes is further complicated by the fact that the mode equation possesses an irregular singular point at the initial time, see \\cite{7,20,10,11}. So far, in the literature, a mathematical argument based on analyticity \\cite{20} and a physical argument based on similarity to the Bunch Davies vacuum \\cite% {10,11} have been discussed and the implications for the CMB have been investigated. Nevertheless, the crucial problem of determining the initial state of modes when they emerge from the Planck scale in an expansion is still essentially unsolved.  Therefore, we here reconsidered the conventional approach to fixing the vacuum: introduce new variables in terms of which the action resembles that of a Minkowski space theory with variable mass - a theory for which the correct vacuum is known. Interestingly, we found that introducing the minimum length into this reformulated action does not yield the same theory - the action and the equations of motion differ slightly. This raised the question as to the extent to which any predictions are affected by this ambiguity.  Let us, therefore, recall the mechanism that leads to the prediction of characteristic oscillations in the CMB spectra due to the minimum length uncertainty relation. The modulus of the mode solution oscillates before horizon crossing if it deviates from the adiabatic solution. In the case of non-de Sitter inflation, modes then cross the Hubble horizon alternatingly with small and large amplitudes and this translates into the prediction of oscillations in the CMB spectra. The oscillations are the more pronounced the more the solutions deviate from the usual adiabatic vacuum in the phase preceding horizon crossing. For the minimum length uncertainty relation an effect of order $\\sigma $ appears possible \\cite{10,11}. This depends, of course, on the choice of initial condition.  On the basis of the current experimental approaches, it is clear that, at best, only effects on the CMB spectrum that are of order $\\sigma $ might become observable over the foreground. Effects of order $\\sigma ^{2}$ must be expected to compete with numerous other effects such as those due to backreaction and the nonlinearity of gravity. We found that in the mode equation the term $\\mu $ dominates the ambiguous term $d$ by a factor $% \\sigma ^{2}$.  On the one hand, this means that the ambiguous term's effect on the amplitude of the minimum-length-induced characteristic oscillations in the CMB spectra is very small. This is because any generic criterion for choosing the modes' initial condition should be affected only to the order of $\\sigma ^{2}$. We note, however, that one might of course conceive of special initial conditions that are arbitrarily different for the two cases. For any generic criterion for initial conditions, the modes' deviation from the adiabatic solution is, therefore correspondingly little affected by whether one starts from the actions of type $S^{(1)}$ or those of the type $% S^{(2)}$. On the other hand, this leaves open the possibility that the phases of the predicted characteristic oscillations in the CMB spectra are affected by the ambiguous term $d$. Indeed, the term $d$ implies a small frequency shift for the mode in the period before horizon crossing. We found, however, that the shift in the oscillation frequency accumulates only to a small fraction ($\\sigma \\approx 10^{-5}$) of a single oscillation (whose amplitude is itself at most of order $\\sigma $) over the entire adiabatic phase.  To summarize, choosing either the actions of type $S^{(1)}$ or the actions of type $S^{(2)}$ generically leads to practically indistinguishable predictions for the effect of the minimum length uncertainty relation on the CMB. Nevertheless, the two actions $S^{(1)}$ or $S^{(2)}$ are actually different and the much-sought-after precise physical criterion for identifying the newly created modes's initial condition may be sensitive to the correct choice of equation of motion. We draw the lesson, therefore, that great care must be taken when introducing a further fundamental constant such as a minimum length into quantum field theory: even if, as usual, ordering ambiguities are resolved by adopting the minimalist approach, previously harmless reformulations of the theory up to a boundary term can lead to an actual change in the theory.  In the particular case at hand one may try to circumvent the problem by starting from the most plausible actions, i.e. those of type $S^{(1)}$, introducing the minimum length and only then reformulating the theory such as to resemble a Minkowski space theory with variable mass. Even this approach faces difficulties due to the introduction of a minimum length scale, however. To see this, note that, usually, modes can be followed arbitrarily far back in the past to when their proper wave length was arbitrarily short and therefore essentially unaffected by curvature, yielding an arbitrarily close match with Minkowski space. With a minimum length, however, the initial conditions must be chosen when the proper wavelength of the mode is the cutoff length. At that time, unavoidably, the match with the Minkowski case is imprecise. In addition, a further redefinition of the field variable followed by a suitable redefinition of the time variable could preserve the appearance of a free field theory on flat space with a time-dependent mass. This would modify, however, the generator of time evolution, i.e. the Hamiltonian. This can affect Hamiltonian-based criteria for picking the vacuum state; a similar point was made in \\cite{12}. An example of a criterion for choosing the vacuum that does not rely on the Hamiltonian, but on field uncertainties instead, was given in \\cite{danielsson, greene3}. It should be interesting to apply this criterion in the case of the minimum length uncertainty relation."
    },
    "0410/astro-ph0410625_arXiv.txt": {
        "abstract": "{An investigation of the stability of the transfer function of the European Southern Observatory's Very Large Telescope Interferometer has been carried out through observations of Fomalhaut, which was observed over a range in hour angle from 21:50--05:24 on 20 October 2002.  No significant variation in the transfer function was found for the zenith angle range 5\\degr--70\\degr. The projected baseline varied between 139.7~m and 49.8~m during the observations and, as an integral part of the determination of the transfer function, a new accurate limb-darkened angular diameter for Fomalhaut of 2.109 $\\pm$ 0.013 mas has been established. This has led to improved values for the emergent flux = (3.43 $\\pm$ 0.10)$\\times10^{8}$~Wm$^{-2}$, effective temperature = 8819 $\\pm$ 67~K and radius = (1.213 $\\pm$ 0.011)$\\times10^{9}$~m (R/R$_{\\sun}$ = 1.744 $\\pm$ 0.016). The luminosity has been found to be (6.34 $\\pm$ 0.20)$\\times10^{27}$~W (L/L$_{\\sun}$ = 16.5 $\\pm$ 0.5). ",
        "introduction": "\\label{sec:intro} A major problem in ground-based optical/infrared long-baseline interferometry is establishing the true interference fringe visibility from measurements made through the Earth's turbulent atmosphere.  The observed visibility is always less than the true visibility due to residual seeing and instrumental effects no matter how well the interferometer is designed, compensated and aligned.  The technique generally adopted to correct for the loss in visibility is to interleave observations of the object being studied with observations of a calibrator - a star for which the true visibility at the baseline and wavelength being employed is known.  A calibrator is thus a star that is either essentially unresolved or whose angular diameter is known with sufficient accuracy to predict the true visibility with the desired accuracy. Division of the observed visibility for the programme object by the observed visibility for the calibrator yields the true visibility for the programme object, providing the instrumental and atmospheric conditions do not change significantly between the measurements.  Optical/infrared interferometers generally measure the square of the fringe visibility as defined by Michelson and, for convenience, we will follow Davis et al. (\\cite{susi2}) and use correlation $C$ to represent the square of the fringe visibility $V$. We define the transfer function $T_\\mathrm{F}$ of an interferometer as the ratio of the observed correlation $C_\\mathrm{obs}$ over the expected correlation $C_\\mathrm{exp}$. The transfer function is thus given by  \\begin{equation} T_\\mathrm{F} = \\frac{V_\\mathrm{obs}^{2}}{V_\\mathrm{exp}^{2}} = \\frac{C_\\mathrm{obs}}{C_\\mathrm{exp}}. \\label{eqn:transfer} \\end{equation}  The true correlation for a programme object is then given by dividing its observed correlation by the transfer function for its associated calibrator.  This procedure works well when programme object and calibrator are close in the sky and the observations are made close in time. However, problems have been experienced at blue wavelengths ($\\sim$440 nm) with the Sydney University Stellar Interferometer (SUSI) (Davis et al. \\cite{susi1}) when the program object and calibrator are separated by more than $\\approx5\\degr$ on the sky, or when the seeing conditions are changing on the timescale of the individual measurements.  As part of the commissioning programme of the European Southern Observatory's Very Large Telescope Interferometer (VLTI) (Glindemann et al. \\cite{vlti}), a large number of calibrator observations have been accumulated (Percheron et al. \\cite{calibvlti}) and are being pursued to assess, among other things, the stability and repeatability of the transfer function in the spectral K band at a wavelength of 2.2~$\\mu$m.  The correlation measurements are made with the VINCI instrument (Kervella et al. \\cite{vinci}), which is based on the FLUOR instrument (Coud\\'{e} du Foresto et al. \\cite{fluor}), and it employs an optical fibre beam-combiner and photometric monitoring of the beams.  The spatial and temporal scales of the turbulence-induced wavefront distortions scale as $\\lambda^{6/5}$ and, with the spatial filtering provided by VINCI, the VLTI is expected to be significantly less affected by atmospheric effects than instruments working at shorter wavelengths and with conventional beam-combiners.  One step proposed for the investigative programme was to observe a single star over a large range in hour angle, and hence in zenith angle $Z$, to investigate whether the transfer function had a dependence on $Z$. Time was allocated for these programmes when the VLTI was operating with its two test siderostats (Derie et al. \\cite{sids}) separated by 140 m. The limiting K magnitude of the system for an unresolved star was $\\sim+3.0$ and, for a partially resolved star, the limiting magnitude was brighter by $2.5\\log C_\\mathrm{exp}$. With this sensitivity limit a shorter baseline would have provided a choice of unresolved stars for the experiment. However, given the circumstances, it was decided to make observations at the longer baseline since they were of intrinsic interest and, in addition, would provide experience for planning more detailed experiments with a shorter baseline in the future. The limiting magnitude placed severe restrictions on the choice of stars, particularly as an unresolved star with declination close to the latitude of the VLTI location on Paranal ($\\approx-25\\degr$) would ideally be chosen for following over a large range in hour angle. At the scheduled time for the observations no suitable unresolved candidate stars were available.  It was decided to follow Fomalhaut ($\\alpha$ PsA, HR 8728, HD216956, spectral type A3V), even though it would be significantly resolved at 140~m based on the Narrabri Stellar Intensity Interferometer limb-darkened angular diameter of $(2.10\\pm0.14)$~mas (Hanbury Brown et al. \\cite{hbdanda}). ",
        "conclusions": "\\label{sec:summary} We have analyzed a set of VLTI interferometric measurements obtained on the bright star Fomalhaut using the VLTI. These measurements were recorded with the main goal of studying the accuracy and stability of the VLTI equipped with the VINCI instrument. Unlike standard interferometric practice, they were not interspersed with independent calibrator observations.  For the night concerned, while it is not possible to completely rule out a dependence on zenith angle, we have shown that the transfer function was remarkably stable even though the observations were taken over large ranges in zenith angle, projected baseline, and observed correlation values.  The observational data are an excellent fit to the transform for a uniform disk and our analysis shows that, if there is a dependence on zenith angle, it is at a level that does not significantly affect the value of the fitted uniform disk angular diameter.  A bonus, and an integral result of the analysis of the observations, was the accurate determination of the angular diameter for Fomalhaut.  This has allowed fundamental quantities for this important star to be determined with significantly improved accuracy.  The angular diameter has been determined without recourse to observations of a calibrator.  This has been possible as a result of observations having been made over a wide range in projected baseline length combined with the stability of the transfer function. It is emphasised that the interleaving of observations of calibrators with those of programme stars will continue to be the normal mode of operation, since it would not be realistic to abandon it on the basis of one night's observations.  However, the implications of the results presented here are that the VLTI in its current configuration, using the VINCI optical fiber beam combiner, gives remarkably stable transfer functions in the short term with a single baseline, and that the angular separation of target and calibrators is not as crucial a matter as found at shorter wavelengths with conventional beam-combiners.  It augurs well for the future of the VLTI and for optical/infrared interferometry in general."
    },
    "0410/astro-ph0410469_arXiv.txt": {
        "abstract": "We examined offsets between H{\\sc ii} regions and molecular clouds belonging to spiral arms of a late type spiral galaxy NGC 4254 (M99). We used a high resolution $^{12}$CO ($J$=1-0) image obtained by Nobeyama Millimeter Array (NMA) and an H$\\alpha$ image. We derived angular offsets ($\\theta$) in the galactic disk, and found that these offsets show a linear dependence on the angular rotation velocity of gas ($\\Omega_{\\rm G}$). This linear relation can be expressed by an equation: $\\theta =(\\Omega_{\\rm G}-\\Omega_{\\rm P})\\cdot t_{\\rm H\\alpha}$, where $\\Omega_{\\rm P}$ and $t_{\\rm H\\alpha}$ are constant. Here, $\\Omega_{\\rm P}$ corresponds to the pattern speed of spiral arms and $t_{\\rm H\\alpha}$ is interpreted as the timescale between the peak compression of the molecular gas in spiral arms and the peak of massive star formation. We may thus determine $\\Omega_{\\rm P}$ and $t_{\\rm H\\alpha}$ simultaneously by fitting a line to our $\\theta - \\Omega_{\\rm G}$ plot, if we assumed they are constant. From our plot, we obtained $t_{\\rm H\\alpha} = ( 4.8 \\pm 1.2 )\\times 10^6$ yr and $\\Omega_{\\rm P} = 26^{+10}_{-6} ~{\\rm km~s^{-1}~kpc^{-1}}$, which are consistent with previous studies. We suggest that this $\\theta - \\Omega_{\\rm G}$ plot can be a new tool to determine the pattern speed and the typical timescale needed for star formations. ",
        "introduction": "\\subsection{Pattern Speed} Pattern speed ($\\Omega_{\\rm P}$) is defined as an angular rotation velocity of spiral pattern. According to the spiral density wave theory \\citep{LS64}, $\\Omega_{\\rm P}$ determines the location and even the existence of resonances, which would greatly affect the morphology and kinematics of a galaxy. The results of numerical simulations also show the large dependence of spiral structure on $\\Omega_{\\rm P}$ \\citep{Semp95}. Despite its importance on the study of galaxies, it cannot be determined directly from observations, since the pattern structure is not a material but a density wave propagating through the disk.  Several methods have been proposed to date for the determination of pattern speed. \\citet{TW84} presented the method using the continuity equation for the surface brightness of galaxies. \\citet{Canz93} showed that the residual velocity field, obtained by subtracting the axisymmetric component from observed velocity field, has a difference in pattern between the inside and outside of the corotation. \\citet{CB90} pointed out that where the arm to interarm ratio of star formation efficiency has a sharp dip could be thought as the corotation. These methods work well for some galaxies, but there are still inconsistencies or large uncertainties.  \\subsection{CO and H$\\alpha$} In this letter, we examined offsets between H$\\alpha$ and CO arms of a spiral galaxy. We suggest that we can use them as a new method for the determination of pattern speed and typical timescale for star formation.  The H$\\alpha$ emission is the tracer of H{\\sc ii} region, which surrounds newly born massive stars, while the CO emission is the tracer of molecular gas. According to the galactic shock wave theory \\citep{Fuji68, Robe69}, most molecular clouds are formed and then condensed by the compression owing to the spiral shock wave and thus belong to the spiral arm. Since galactic scale star formations take place in these molecular arms, some molecular gases are used to form stars, and the others are dissociated by UV emission from massive stars born in that star formation. In consequence, molecular clouds disappear soon after their formation and their position appears to be fixed to the spiral pattern, though individual molecular clouds move at the speed of gas. Therefore, the position of spiral pattern and H$\\alpha$ arm can be measured by using that of a molecular arm and an ensemble of H{\\sc ii} regions respectively. In other words, offset between the arms of H$\\alpha$ and CO represents the difference between the speed of gas and spiral pattern, as well as the time needed for star formation from molecular clouds.  This offset has been found in many spiral galaxies (eg. \\cite{RK90}). We can easily see it in a $B$-band image as a displacement of an arm's peak and a dust lane, since the light of $B$-band is mainly emitted by the young OB star and the distribution of dust lane and molecular gas are almost the same. However, this offset has not been studied quantitatively yet, since the resolution of the CO image was too large to examine detail structures of spiral arms. ",
        "conclusions": "We examined the offset between the arm of H$\\alpha$ and CO of SAc galaxy NGC 4254, derived ten offset values and found the linear relation between the offset angle and the angular rotation velocity. This linearity implies that a rigid spiral pattern exists and that the physical processes of star formation can be regarded as not to extremely change in the disk. We emphasize that this is the first work to show this relationship.  With this relation, we derived the pattern speed $\\Omega_{\\rm P} =26^{+10}_{-6} ~{\\rm km~s^{-1}~kpc^{-1}}$, and the time delay for star formation $t_{\\rm H\\alpha} = ( 4.8 \\pm 1.2 )\\times 10^6$ yr. We again emphasize that we can derive both $\\Omega_{\\rm P}$ and $t_{\\rm H\\alpha}$ simultaneously. Since the obtained values are consistent with previous studies, the way of our analysis can be a new method for the determination of the pattern speed and the typical timescale needed for star formations. Thus, this method can investigate not only the kinematics, but also the star formation mechanisms quantitatively. Moreover, this method can be applied to many spiral galaxies, since offsets between H$\\alpha$ and CO arms have been found in many spirals.  \\bigskip We are very grateful to Rebecca A. Koopmann, Jeffrey D. P. Kenney, and Judith Young for kindly providing us with their $R$ and H$\\alpha$ image of NGC 4254. H. N. is financially supported by a Research Fellowship from the Japan Society for the Promotion of Science for Young Scientists."
    },
    "0410/astro-ph0410143_arXiv.txt": {
        "abstract": "The outbursts of the dwarf nova WZ Sagittae occur on timescales which are much longer than expected from the standard thermal viscous disc instability model, which correctly reproduces the outburst behaviour of many other dwarf novae. In addition short `echo' outbursts have sometimes been observed after the main outburst. The increased timescales could be explained by a magnetic propeller, anchored to the primary star, which clears the inner disc. We argue that the propeller, combined with the resonances of the binary orbit, may also be the cause of the echo outbursts. We present numerical models which reproduce the echo outbursts. ",
        "introduction": "WZ Sagittae is a member of the SU Ursae Majoris class of dwarf novae. These are short period systems which exhibit long superoutbursts in addition to ordinary outbursts. During these superoutbursts periodic modulations of the optical light curve, known as superhumps are observed. These features have a frequency which is slightly longer than that of the binary orbit, and which is usually attributed to a precessing disc. Superoutbursts in the SU UMa systems typically occur with recurrence times of weeks or months.  Stars from the subclass to which WZ Sge belongs are sometimes known as TOADs or Tremendous Outburst Amplitude Dwarf novae. These systems have very long recurrence times and correspondingly bright superoutbursts. WZ Sge itself showed a 33 year recurrence time, until this pattern was broken by the 2001 outburst which occurred after only 23 years. WZ Sge, the brightest of all dwarf novae, is also notable for its extreme mass ratio. According to \\citet{ski02} $M_{2}=0.045  \\pm 0.003 \\ M_{\\sun}$ and $M_{1} = 0.74 \\pm 0.07 \\ M_{\\sun}$. It has a binary period of $P_{\\rm{orb}} = 81.6 \\ {\\rm{min}}$ \\citep{krz62} and a primary spin period of $P_{\\rm{spin}} = 27.87 \\ {\\rm{s}}$. In this outburst \\citep[e.g][]{pat02} superhumps were observed, but prior to these a modulation at the frequency of the binary orbit was observed by \\citet{pat02} as discussed by e.g. \\citet{kun04}. In addition, the main outburst was followed by twelve echo outbursts. Echoes have also been observed in the similar system EG Cancri \\citep{osa01}.  Numerical models incorporating the thermal viscous disc instability model (DIM) due to \\citet{hos79}, have reproduced the outburst behaviour of many dwarf noave. These models have not however, without using extremely low viscosities, been able to explain the long recurrence time of WZ Sge type systems. \\citet{osa89} extended the DIM to binaries with small mass ratios when \\citet{whi88} showed that discs in such systems could become tidally unstable. Osaki proposed that, during the normal outburst cycle, material is allowed to build up in the outer disc. When the disc becomes sufficiently large that it may access the 3:1 resonance of the binary orbit, it will become eccentric, enhancing dissipation and the loss of angular momentum, leading to a superoutburst. Simulations \\citep[e.g.][]{tru01} supported this theory.  The extended recurrence times of the WZ Sge type systems may be explained by the use of a magnetic propeller to evacuate the inner regions of the disc. This mechanism, similar to that proposed by \\citet{mat04a} for the FU Orionis stars, is detailed in the forthcoming paper by \\citet{wyn04}. The magnetic propeller is a mechanism by which the inner regions of an accretion disc can be evacuated. The co-rotation radius $R_{\\rm{co}}$ is the radius at which a Keplerian orbit has the same angular frequency as that of the star's spin. If the primary magnetic field rotates as a solid body then this radius can be said to divide the magnetic accretion and magnetic propeller regimes. Particles at radii smaller than the co-rotation radius lose angular momentum as a result of their interactions with the stellar magnetic field, and fall rapidly onto the star. This magnetically enhanced accretion occurs because the particles are moving more rapidly than the magnetic field. Elements of plasma which exist outside co-rotation are, conversely, accelerated by the magnetic field, and gain angular momentum at the star's expense. This is the magnetic propeller regime. An example of a strong magnetic propeller is that of AE Aquarii \\citep{wyn97}, in which accretion is completely prevented by the magnetic field. What is proposed here is a far weaker propeller, driven by a rapidly spinning, moderately magnetic white dwarf. Simulations using weaker fields have been shown to warp accretion discs \\citep{osu04}. \\citet{las99} also used a magnetically truncated disc in their model of WZ Sge.  The cause of the echo outbursts is as yet uncertain. \\citet{osa01} suggested that the mass which accumulates near the 3:1 resonance in the disc is released gradually after the main outburst, leading to echo outbursts. In this paper we model a full outburst of WZ Sge using smoothed particle hydrodynamics (SPH). The rapidly spinning white dwarf creates a magnetic propeller which, in addition to evacuating the inner disc, aids the creation of a reservoir of material in the outer disc. Echo outbursts arise naturally from this configuration. ",
        "conclusions": "In this paper echo outbursts are reproduced in an SPH model of WZ Sge. These echoes are similar to those observed in the light curve of the system However, it is not yet clear whether this mechanism is likely to be the one which occurs in reality. In order to establish this there are a number of refinements which will be made to the model. Firstly the viscous alphas used in the model will be brought closer to their accepted values. It is not desirable to rely on time scaling arguments in this system since neither magnetic nor resonant effects scale with alpha. The $\\sim 30,000$ particles used are not sufficient to fully resolve such a disc in three dimensions, so higher resolution calculations will be made. In addition a mechanism for the termination of the echoes will be sought."
    },
    "0410/astro-ph0410696_arXiv.txt": {
        "abstract": "AA\\,Dor (LB\\,3459) is an eclipsing, close, single-lined, post common-envelope binary (PCEB) consisting of an sdOB primary star and an unseen secondary with an extraordinary small mass -- formally a brown dwarf. The brown dwarf may have been a former planet which survived a common envelope phase and has even gained mass.  A recent determination of the components' masses from results of state-of-the-art NLTE spectral analysis and subsequent comparison to evolutionary tracks shows a discrepancy between masses derived from radial-velocity and the eclipse curves. Phase-resolved high-resolution and high-SN spectroscopy was carried out with FUSE in order to investigate on this problem.  We present preliminary results of an ongoing NLTE spectral analysis of FUSE spectra of the primary. ",
        "introduction": "A recent spectral analysis of the primary of AA\\,Dor (Rauch 2000, 2004) has shown a discrepancy between the resulting stellar parameters from spectral analysis and those derived from the radial-velocity and the eclipse curves. Recently, Hilditch et al\\@. (2003) presented new photometry data and an improved photometric model of AA\\,Dor which verified earlier results within smaller error ranges. Thus, one reason for this discrepancy might be uncertainties in the spectral analysis, i.e\\@. mainly in the surface gravity determination.  Phase-resolved high-resolution UV spectra are necessary in order to reduce the error limits and to identify individual iron and nickel lines in the spectrum. Ten LWRS observations with short exposure times (200\\,s) should be performed, in order to minimize the effects of orbital smearing. Four have already been executed on Aug 29, 2003. Five were carried out on Jun 22, 2004. The preliminary analysis which is presented here makes use of the 2003 data, only.   \\begin{figure}[ht] \\plotone{rauch1_fig1.eps} \\caption{Comparison of the four co-added FUSE spectra from Aug 2003 (bottom) of the sdOB primary of AA\\,Dor with the synthetic spectrum of our final model (top, Rauch 2000). Both are smoothed with a Gaussian of 0.2\\,\\AA\\ FWHM for clarity. H, C, N, and O lines considered in the model are indicated. All measured Kurucz lines (POS, Kurucz 1996) are considered. The synthetic spectrum is convolved with a rotational profile with 35 km/sec. } \\end{figure}  The plane-parallel, static models (H+He+C+N+O+Mg+Si+IronGroup included), used for the spectral analysis of AA\\,Dor, are calculated with TMAP, the T\\\"ubingen NLTE Model Atmosphere Package (Werner et al\\@. 2003, Rauch \\& Deetjen 2003).   \\begin{figure}[htb] \\plotone{rauch1_fig2.eps} \\caption{Total abundance of the iron-group elements. We reduced the resolution of the FUSE spectrum to 7\\,\\AA\\ (like the low-resolution IUE spectrum SWP\\,04887 for wavelength greater than 1186\\,\\AA). The overall slope of the flux is matched well by a solar abundance of the iron-group elements. } \\end{figure}  \\clearpage  \\begin{figure}[ht] \\plotone{rauch1_fig3.eps} \\caption{Rotational velocity of the primary AA\\,Dor. The FUSE spectrum around C\\,{\\sc iii} 1175\\,\\AA\\ is compared with a synthetic spectrum which is convolved with a rotational profile with $v_\\mathrm{rot}=35\\,\\mathrm{km/sec}$ (bottom). The top panel shows three synthetic spectra calculated with $v_\\mathrm{rot}=20, 35, 50\\,\\mathrm{km/sec}$ (dashed, thick, thin line).} \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410375_arXiv.txt": {
        "abstract": "We present neutral hydrogen observations of the nearby edge-on spiral galaxy NGC\\,891 which show extended extra-planar emission up to distances of 15 kpc from the plane. 3D modeling of the galaxy shows that this emission comes from halo gas rotating more slowly than the gas in the disk. We derive the rotation curves of the gas above the plane and find a gradient in rotation velocity of $-$15 \\kms\\ kpc$^{-1}$. We also present preliminary results of a galactic fountain model applied to NGC\\,891. ",
        "introduction": "NGC\\,891 is one of the best known and studied nearby edge-on spiral galaxies. It is at the distance of 9.5 Mpc, is classified as a Sb/SBb, and it is often referred to as a galaxy very similar to the Milky Way \\citep{vdk81}. Because of its very high inclination \\citep[i$\\geq$88.6\\ci, ][]{rup91} it is very suitable for the study of the distribution and kinematics of the gas above the plane.  NGC\\,891 has been the subject of numerous studies at different wavelengths that have led to the detection of various halo components: an extended radio halo \\citep*{all78, hum91}, an extended layer of diffuse ionised gas (DIG) \\citep[e.g.][]{det90} and diffuse extra-planar hot gas \\citep{bre94}. Also ``cold'' ISM components have been detected in the halo such as \\hi\\ \\citep*{swa97}, dust \\citep{how99} and CO \\citep{gar92}.  Here we concentrate on the neutral gas and present results from recent third generation \\hi\\ observations obtained with the Westerbork Synthesis Radio Telescope (WSRT). NGC\\,891 was first studied in \\hi\\ in the late seventies with the WSRT and the presence of neutral gas seen in projection out of the plane was reported \\citep{san79}. Subsequently, a new study with higher sensitivity showed that the extra-planar emission was very extended, up to 5 kpc from the plane. 3D modeling indicated that such emission was produced by a thick layer of neutral gas rotating more slowly than the gas in the disk \\citep{swa97}.  Since then, several other studies have confirmed the presence of neutral gas in the halos of spiral galaxies. It has been detected in edge-on or nearly edge-on systems (e.g. UGC\\,7321, \\citet{mat03} and NGC\\,253, \\citet{boo04}), as well as in galaxies viewed at different inclination angles \\citep[e.g.\\ NGC\\,2403,][]{fra01}. Indications of vertical gradients in rotation velocity have been found in several galaxies also in the ionised gas (e.g.\\ NGC\\,891, \\citet{pil94}, \\citet{ran97}; NGC\\,5055, \\citet{ran00}).  In the first part of the paper we present the new \\hi\\ observations of NGC\\,891 together with a 3D modeling of the \\hi\\ layer. In the second part (Section 3) we study the kinematics of the extra-planar gas and, in the third one (Section 4), we present results from a dynamical model of the extra-planar gas. ",
        "conclusions": "New \\hi\\ observations of the edge-on spiral galaxy NGC\\,891 show the presence of extraplanar neutral gas up to distances of 15 kpc from the plane. We have used these data to derive, for the first time, a 2D rotation velocity field (in R and z) of an edge-on spiral galaxy. We find that the extra-planar gas is corotating with the gas in the disk up to about 1.3 kpc; beyond that it rotates more slowly by about 15 \\kms\\ kpc$^{-1}$. Dynamical models of NGC\\,891 show that a pure ballistic galactic fountain can qualitatively reproduce the data. However, problems with missing low angular momentum material suggest that other mechanisms such as accretion may also play a role."
    },
    "0410/astro-ph0410719_arXiv.txt": {
        "abstract": "The Cl 1604 supercluster at $z\\sim 0.9$ is known to contain at least four distinct member clusters, separated in both projection and redshift.  In this paper we present deep, multicolor wide-field imaging of a region spanning $\\sim45'$ on a side, corresponding to $21~h^{-1}_{70}$ Mpc (physical) at the supercluster redshift.  We select galaxies whose colors correspond to those of spectroscopically confirmed cluster members in the $i'$ versus $r'-i'$ color-magnitude diagram. Using an adaptive kernel, we generate a map of the projected red galaxy density and identify numerous new candidate clusters that are likely supercluster members. Assuming that all of the density peaks are associated with the supercluster, its transverse size is $\\sim10~h^{-1}_{70}$ Mpc, which is still significantly smaller than the nearly $93~h^{-1}_{70}$ Mpc depth in redshift space. ",
        "introduction": "The Cl 1604 supercluster at $z\\sim0.9$ is one of the best studied high-redshift large scale structures. It was initially detected as two separate clusters in the plate-based survey of \\citet{gho86}, with redshifts and preliminary velocity dispersions measured by \\citet{pos98,pos01}. The proximity of the two clusters Cl 1604+4304 and Cl 1604+4321 in both radial velocity (4300~{\\kms}) and position on the sky ($17'$) prompted \\citet{lub00} to perform deep multi-band imaging with COSMIC \\citep{kel98} on the Palomar 5-m telescope covering the area between the two clusters. The imaging revealed additional overdensities of red galaxies whose colors were consistent with spectroscopically-confirmed, early-type galaxies in the two original clusters \\citep[see Figs. 2 and 4 of][]{lub00}, suggesting that the clusters were part of a high-redshift supercluster. A total of four distinct red galaxy density peaks were identified in a contiguous area of $10\\farcm4 \\times 18\\farcm2$. \\citet{gal04} conducted a large spectrographic survey using LRIS and DEIMOS \\citep{fab03} on the Keck 10-m telescopes to confirm the new cluster candidates and provide accurate velocity dispersions for all four supercluster components. They detected 230 total cluster members, confirming that all four components were in fact supercluster members. These four clusters were nevertheless well-separated both in projection and in redshift, and have velocity dispersions ranging from 480 to 980 km s$^{-1}$, corresponding to Abell richness classes of $R\\sim0-2$.  The large radial depth of the supercluster ($\\sim93~h^{-1}_{70}$ Mpc) in comparison to the small area imaged with COSMIC ($4.8\\times8.5~h^{-1}_{70}$ Mpc) suggested that this structure could extend to a significantly larger area on the sky. To test this hypothesis and search for new supercluster components, we undertook a wide area, multi-band imaging survey using the Large Format Camera \\citep[LFC;][]{sim00} on the Palomar 5-m telescope.  In this paper, we present the results of this survey, including several new candidate supercluster members.  We assume a $\\Lambda$CDM cosmology with $\\Omega_m=0.3, \\Lambda=0.7,$ and ${\\rm H_0}=70~h_{70}~{\\rm km~s^{-1}~Mpc^{-1}}$. ",
        "conclusions": "If the red galaxy overdensities detected in this imaging survey are truly members of the Cl 1604 supercluster at $z\\sim0.9$, we have identified by far the largest known structure at redshifts approaching unity. The largest projected separation (AG) is approximately $9.6~h^{-1}_{70}$ Mpc, still much smaller than the apparent $93~h^{-1}_{70}$ Mpc depth. It appears that we have detected the southeast limit of the supercluster; no significant overdensities are visible within $\\sim10~h^{-1}_{70}$ Mpc directly to the east of A. We cannot rule out the structure continuing to the north/northeast or southwest. Most of the highest density peaks, including the four known clusters, are surrounded by smaller, lower density enhancements, suggesting a filamentary nature that is typical of large scale structure in simulations \\citep{evr02}. The large range of densities in this supercluster makes it an ideal laboratory to study the dependence of galaxy properties and their evolution on their physical environment.  We note that the apparent 8:1 axial ratio of this supercluster is comparable to or even less than that found in simulations. \\citet{sha04} provide an excellent discussion of shape measurements for filaments and voids in numerical simulations with $\\Lambda$CDM cosmologies, demonstrating that superclusters typically have non-trivial, highly filamentary topologies with order-of-magnitude or greater differences in the lengths of their dimensions (see their Figures 14 and 15). Furthermore, any complex filament will be easier to detect if it is viewed in projection along its longest axis. If the angular separation between the sufficiently dense regions of a supercluster becomes larger than the contiguously imaged area in a survey, then it will only be detected as a single cluster.  Our results demonstrate the efficiency of deep, moderately large area imaging in the vicinity of known clusters for detecting further structures. Hubble volume simulations show that similar superclusters are not uncommon at $z\\sim1$ \\citep{evr02}. Therefore, we have undertaken a wide-field imaging survey of a large sample of $z > 0.5$ clusters to search for other large scale structures at these cosmologically significant redshifts. Using color cuts corresponding to early-type galaxies at a variety of redshifts we can also detect clusters at distances other than that of the target, in a manner similar to the cut-and-enhance technique of \\citet{got02} and the Red Sequence Cluster Survey \\citep{gla00}."
    },
    "0410/astro-ph0410233_arXiv.txt": {
        "abstract": "We present theoretical predictions for the spectral, temporal and intensity signatures of the electromagnetic radiation emitted during the process of the gravitational collapse of a stellar core to a black hole, during which electromagnetic field strengths rise over the critical value for $e^+e^-$ pair creation. The last phases of this gravitational collapse are studied, leading to the formation of a black hole with a subcritical electromagnetic field, likely with zero charge, and an outgoing pulse of initially optically thick $e^+e^-$-photon plasma. Such a pulse reaches transparency at Lorentz gamma factors of $10^2$--$10^4$. We find a clear signature in the outgoing electromagnetic signal, drifting from a soft to a hard spectrum, on very precise time-scales and with a very specific intensity modulation. The relevance of these theoretical results for the understanding of short gamma-ray bursts is outlined. ",
        "introduction": "The discovery in 1997 of the afterglows of Gamma-Ray Bursts (GRBs) \\cite{C97} has evidenced the cosmological nature of these sources. By the analysis of the first and second BATSE catalogs\\footnote{see http://cossc.gsfc.nasa.gov/batse/} Tavani in 1998 \\cite{T98} confirmed the existence of two families of GRBs: the so-called ``long-bursts'' with a soft spectrum and duration $\\Delta t > 2.5$sec and the ``short-bursts'' with harder spectrum and duration $\\Delta t < 2.5$sec. In 2001 the theory was advanced \\cite{lett2} that both short-bursts and long-bursts originate from the same underlying physical process due to the vacuum polarization of electromagnetic overcritical gravitational collapse leading to the creation of $e^+-e^-$ pairs at the expenses of the extractable energy of a black hole \\cite{CR71}. The difference between the short-bursts and long-bursts in this theory is mainly due to the amount of baryonic matter encountered by the $e^+e^-$ pairs in their relativistic expansion. A support of such a theory was given by Schmidt \\cite{S01} showing that short-bursts and long-bursts have the same isotropic-equivalent characteristic peak luminosity.  In recent work we have systematically developed the theoretical background of a process of gravitational collapse of matter involving an electromagnetic field with field strength higher than the critical value for $e^+e^-$ pair creation \\cite{CRV02,RV02,RV03,RVX03a,RVX03b,RVX03c}. The goal has been to clarify the physical nature of the process of extracting the mass energy of a black hole by the creation of $e^+e^-$ matter pairs \\cite{CR71} and to analyze the electromagnetic radiation emission process during the transient dynamical phases of the gravitational collapse leading to the final formation of the black hole.  In this letter we conclude this analysis by making precise predictions for the spectra, the energy fluxes and characteristic time-scales of the radiation for short-bursts. If the precise luminosity variation and spectral hardening of the radiation we have predicted will be confirmed by observations of short-bursts, these systems will play a major role as standard candles in cosmology.  These considerations will also be relevant for the analysis of the long-bursts when the baryonic matter contribution will be taken into account. ",
        "conclusions": "The above results were obtained considering $e^+e^-$ plasma without any baryonic contamination and are therefore directly relevant for short-bursts \\cite{lett2}. The characteristic spectra, time variabilities and luminosities of the electromagnetic signals from collapsing overcritical stellar cores, here derived from first principles, agrees very closely with the observations of short-bursts \\cite{P...99}. New space missions must be planned, with temporal resolution down to fractions of $\\mu$s and higher collecting area and spectral resolution than at present, in order to verify the detailed agreement between our model and the observations. It is now clear that if our theoretical predictions will be confirmed, we would have a very powerful tool for cosmological observations: the independent information about luminosity, time-scale and spectrum can uniquely determine the mass, the electromagnetic structure and the distance from the observer of the collapsing core, see e.g.\\ Fig.~\\ref{ff5} and Ref.~\\cite{RFVX04}. In that case short-bursts may become the best example of standard candles in cosmology \\cite{R03}. The introduction we are currently analysing is the introduction of baryonic matter in the optically thick phase of the expansion of the $e^+e^-$ plasma which can affect the structure of the Proper-GRB (P-GRB) \\cite{lett1} as well as the structure of the long-bursts \\cite{Brasile}.  An interesting proposal was advanced in 2002 \\cite{it02} that the $e^+e^-$ plasma may have a fundamental role as well in the physical process generating jets in the extragalactic radio sources. The concept of dyadosphere originally introduced in Reissner-Nordstr\\\"{o}m black hole in order to create the $e^+e^-$ plasma relevant for GRBs can also be generalized to the process of vacuum polarization originating in a Kerr-Newman black hole due to magneto-hydrodynamical process of energy extraction (see e.g. \\cite{p01} and references therein). The concept therefore introduced in this letter becomes relevant for both the extraction of rotational and electromagnetic energy from the most general black hole \\cite{CR71}.  After the submission of this letter we have become aware that Ghirlanda et al. \\cite{g03} have given evidence for the existence of an exponential cut off at high energies in the spectra of short bursts. We are currently comparing and contrasting these observational results with the predicted cut off in Fig.~\\ref{ff4} which results from the existence of the separatrix introduced in \\cite{RVX03a}. The observational confirmation of the results presented in Fig.~\\ref{ff4} would lead for the first time to the identification of a process of gravitational collapse and its general relativistic self-closure as seen from an asymptotic observer."
    },
    "0410/nucl-th0410040_arXiv.txt": {
        "abstract": "We study the hadron-quark phase transition in the interior of neutron stars (NS). For the hadronic sector, we use a microscopic equation of state (EOS) involving nucleons and hyperons derived within the Brueckner-Hartree-Fock approach. For the quark sector, we employ the MIT bag model, as well as the Nambu--Jona-Lasinio (NJL) and the Color Dielectric (CD) models, and find that the NS maximum masses lie in the interval between 1.5 and 1.8 solar masses. ",
        "introduction": "The appearence of quark matter in the interior of massive neutron stars is one of the main issues in the physics of these compact objects \\cite{pulsar}. Calculations of NS structure, based on a microscopic nucleonic equation of state, indicate that for the heaviest NS, close to the maximum mass (about two solar masses), the central particle density reaches values larger than $1/\\rm fm^{3}$. In this density range, it can be expected that the nucleons start to loose their identity, and quark degrees of freedom are excited at a macroscopic level.  The value of the maximum mass of NS is probably one of the physical quantities that are most sensitive to the presence of quark matter in the core. Unfortunately, the quark matter EOS is poorly known at zero temperature and at the high baryonic density appropriate for NS. One has, therefore, to rely on models of quark matter, which contain a high degree of uncertainty. In this paper we use a definite nucleonic EOS, which has been developed on the basis of the Brueckner many-body theory, and three different models for the quark EOS, respectively the MIT bag model, the Nambu--Jona-Lasinio and the Color Dielectric models. We compare the predictions of different models, and estimate the uncertainty of the results for the NS structure and mass. ",
        "conclusions": ""
    },
    "0410/astro-ph0410268_arXiv.txt": {
        "abstract": "An impressive amount of different astrophysical data converges towards the picture of a spatially flat universe undergoing a today phase of accelerated expansion. The nature of the dark energy dominating the energy content of the universe is still unknown and a lot of different scenarios are viable candidates to explain cosmic acceleration. Most of the methods employed to test these cosmological models are essentially based on distance measurements to a particular class of objects. A different method, based on the lookback time to galaxy clusters and the age of the universe, is used here. In particular, we constrain the characterizing parameters of  three classes of dark energy cosmological models to see whether they are in agreement with this kind of data,  based on time measurements rather than distance observations. ",
        "introduction": "The increasing bulk of data that have been accumulated in the last few years have paved the way to the emergence of a new standard cosmological model usually referred to as the {\\it concordance model}. The Hubble diagram of Type Ia Supernovae (hereafter SNeIa), measured by both the Supernova Cosmology Project \\cite{SCP} and the High\\,-\\,z Team \\cite{HZT} up to redshift $z \\sim 1$, first indicated that the universe is undergoing a phase of accelerated expansion. On the other hand, balloon born experiments, such as BOOMERanG \\cite{Boomerang} and MAXIMA \\cite{Maxima}, determined the location of the first and second peak in the anisotropy spectrum of  cosmic microwave background radiation (CMBR) strongly indicating that the geometry of the universe is spatially flat. If combined with constraints coming from galaxy clusters on the matter density parameter $\\Omega_M$, these data indicate that the universe is dominated by a non-clustered fluid with negative pressure generically dubbed {\\it dark energy} which is able, eventually, both to close the universe and drive its accelerated expansion. This picture has been further strengthened by the more precise measurements of the CMBR spectrum, due to the WMAP experiment \\cite{WMAP}, and by the extension of the SNeIa Hubble diagram to $z > 1$ \\cite{Riess04}.  After the discovery of this scenario, an overwhelming flood of papers, presenting a great variety of models which try to explain this phenomenon has appeared; in any case,  the simplest explanation is claiming for the well known cosmological constant $\\Lambda$ \\cite{LCDMrev}. Although the best fit to most of the available astrophysical data \\cite{WMAP}, the $\\Lambda$CDM model failed in explaining why the inferred value of $\\Lambda$ is so tiny (120 orders of magnitude lower) compared to the typical vacuum energy values predicted by particle physics and why its energy density is comparable to the matter density right now (the so called {\\it coincidence problem}). As a tentative solution, many authors have replaced the cosmological constant with a scalar field rolling down its potential and giving rise to the model now referred to as {\\it quintessence} \\cite{QuintFirst}. Even if successful in fitting the data, the quintessence approach to dark energy is still plagued by the coincidence problem since the dark energy and matter densities evolve differently and reach comparable values for a very limited portion of the universe evolution that is taken to coincide with the present era. Moreover, it is not clear where this scalar field originates from thus leaving a great uncertainty on the choice of the scalar field potential (see \\cite{QuintRev} for comprehensive reviews).  The subtle and elusive nature of  dark energy has led many authors to look for completely different scenarios able to give a quintessential behavior without the need of exotic components.  To this aim, it is worth stressing that the acceleration of the universe only claims for a negative pressure dominant component, but does not tell anything about the nature and the number of cosmic fluids filling the universe. This consideration suggests that it could be possible to explain the accelerated expansion by introducing a single cosmic fluid with an  equation of state causing it to act like dark matter at high densities and dark energy at low densities. An attractive feature of these models, usually referred to as {\\it Unified Dark Energy} (UDE) or {\\it Unified Dark Matter} (UDM) models, is that such an approach naturally solves, al least phenomenologically, the coincidence problem. Some interesting examples are the generalized Chaplygin gas \\cite{Chaplygin}, the tachyon field \\cite{tachyon} and the condensate cosmology \\cite{Bassett}.  A different class of UDE models has been recently proposed by some of us \\cite{Hobbit} where a single fluid is considered whose energy density scales with the redshift in such a way that the radiation dominated era, the matter domination era and the accelerating phase can be naturally achieved. It is worth noting that such a model is extremely versatile since it can be interpreted both in the framework of UDE models and as a two-fluid scenario with dark matter and scalar field dark energy.  Actually, there is still a different way to face the problem of cosmic acceleration. As  stressed in Lue et al. \\cite{LSS03}, it is possible that the observed acceleration is not the manifestation of another ingredient in the cosmic pie, but rather the first signal of a breakdown of our understanding of the laws of gravitation. From this point of view, it is thus tempting to modify the Friedmann equations to see whether it is possible to fit the astrophysical data with a model comprising only the standard matter. Interesting examples of this kind are the Cardassian expansion \\cite{Cardassian} and the DGP gravity \\cite{DGP}.  Moving in this same framework, it is possible to find alternative schemes where a quintessential behavior is obtained by simply taking into account the effective contribution to cosmology of some usually neglected fundamental physics \\cite{capozcurv,curvature,review}. For instance, a cosmological constant term may be recovered as a consequence of a non\\,-\\,vanishing torsion field thus leading to a model which is consistent with both SNeIa Hubble diagram and Sunyaev\\,-\\,Zel'dovich data coming from clusters of galaxies \\cite{torsion}. SNeIa data could also be efficiently fitted including higher-order curvature invariants in the gravity Lagrangian \\cite{curvfit}. It is worth noting that these alternative schemes provide naturally a cosmological component with negative pressure whose origin is simply related to the geometry of the universe thus overcoming the problems linked to the physical significance of the scalar field.  It is evident, from this short overview, the high number of cosmological models which are viable candidates to explain the observed accelerated expansion. This abundance of models is somehow a result of having only a limited number of cosmological tests to discriminate among rival theories so that any new observable which could be added to the usual ones is welcome. To this aim, it is useful to remember that both the SNeIa Hubble diagram and the recently proposed angular size\\,-\\,redshift relation of compact radio sources \\cite{AngTest} are distance based methods to probe cosmological models. From this point of view, it is  interesting to look for  tests based on time-dependent observables. To this aim, one can take into account the {\\it lookback time} to distant objects since this quantity can discriminate among different cosmological models. The lookback time is observationally estimated as the difference between the present day age of the universe and the age of a given object at redshift $z$. Such an estimate is possible if the object is a galaxy observed in more than one photometric band since its color is determined by its age as a consequence of stellar evolution. It is thus possible to get an estimate of the galaxy age by measuring its magnitude in different bands and then using stellar evolutionary codes to choose the model that reproduces the observed colors at best. A quite similar approach was pursued by Lima \\& Alcaniz \\cite{LA00} who used the age (rather than the lookback time) of old high redshift galaxies to constrain the dark energy equation of state (see also \\cite{Jimenez}). The same method was then applied also to braneworld models \\cite{AJD02} and the Chaplygin gas \\cite{AJD03}.  It is worth noting, however, that the estimate of the age of a single galaxy may be affected by systematic errors which are difficult to control. Actually, this problem can be overcome by considering a sample of galaxies belonging to the same cluster. In this way, by averaging the estimates of all galaxies, one obtains an estimate of the cluster age and reduces the systematic errors. Such a method was first proposed by Dalal et al. \\cite{DAJM01} and then used by Ferreras et al. \\cite{FMT03} to test a class of models where a scalar field is coupled with the matter term giving rise to a particular quintessence scheme. We improve here this analysis using a different cluster sample and testing three different approaches to the dark energy problem. Moreover, we add a further constraint to better test the dark energy models and assume that the age of the universe for each model is in agreement with recent estimates. Note that this is not equivalent to the lookback time as we will discuss below.  The layout of the paper is the  following. In Sect.\\,II, we briefly present the three classes of cosmological models considered giving also the main quantities which we  need for lookback time test. The method  is developed in Sect.\\,III and the used data are presented in Sect.\\,IV. A detailed discussion of the results is the subject of Sect.\\,V. We summarize and conclude in Sect.\\,VI. ",
        "conclusions": "The impressive amount of data  indicating a spatially flat universe in accelerated expansion has posed the problem of dark energy and stimulated the search for cosmological models able to explain such unexpected behavior. Many rival theories have been proposed to solve the puzzle of the nature of  dark energy ranging from a rolling scalar field to a unified picture where a single exotic fluid accounts for the whole dark sector (dark matter and dark energy). Moreover, modifications of the gravity Lagrangian has also been advocated. Although deeply different in their underlying physics, all these scenarios share the common feature of well reproducing the available astrophysical data. It is worth stressing, however, that the most widely used cosmological tests and in particular the SNeIa Hubble diagram and the angular size\\,-\\,redhisft relation) are essentially based on distance measurements to high redshift objects and are thus affected by similar systematic errors. It is hence particularly interesting to look for  methods which are related to the estimates of different quantities. Being affected by other kinds of observational problems, such  methods could be considered as   cross checks for the results obtained by the usual tests and they should represent complementary probes for cosmological models.  The technique we have devised here is a first  step in this direction. We have used the present age of the universe and the lookback time to galaxy clusters to build a sort of {\\it time diagram} of the universe in order to reconstruct its age evolution. Relying on stellar evolutionary codes, the estimate of the lookback time is related to a different astrophysics than the distance based methods and it is thus free from any problem connected with the evolution of standard candles (such as the SNeIa absolute magnitude and the intrinsic linear size of radio sources). Actually, this technique could be affected by its own systematics (such as, for instance, the degeneracy between age and metallicity), but these may be better controlled. Moreover, comparing the results thus obtained with those coming from distance based methods allows to strengthen the conclusions suggested by both techniques.  Motivated by these considerations,  we have first applied our method to the $\\Lambda$CDM model in order to estimate the present day values of the matter density parameter $\\Omega_M$ and of the dimensionless Hubble constant $h$ obtaining (at the $1 \\sigma$ CL)\\,:  \\begin{displaymath} \\Omega_M = 0.22_{-0.10}^{+0.13} \\ \\ \\ \\ ,\\ \\ \\ \\ h = 0.71_{-0.08}^{+0.07} \\ . \\end{displaymath} These values are in agreement with the previous estimates in literature. This can be considered an independent confirmation not only of the viability of the $\\Lambda$CDM model, but also of our method. It is worth noting that the $\\Lambda$CDM scenario receives further support from this test and henceforth the cosmological constant $\\Lambda$ still remains the best candidate to explain the dark energy puzzle from an observational point of view.  Nonetheless, the $\\Lambda$CDM model is severely affected by conceptual problems so that it is worthwhile to look fo alternative approaches. This has stimulated a plenty of models where the cosmic acceleration is due to a dominant scalar field rolling down its potential. However, such a scheme, on the one hand, still does not solve the coincidence problem and, on the other hand, is plagued by the unidentified nature of the scalar field itself and the ignorance of its self interaction potential. These considerations have opened the way to different models that are able to give an accelerated expansion without the need of scalar fields. Two of these approaches have been tested in this paper.  In \\cite{Hobbit}, it has been proposed a phenomenological unified model where a single fluid with a given energy density assigned by few parameters is able to fit both the SNeIa Hubble diagram and the angular size\\,-\\,redshift relation for ultracompact radio sources without the need of any cosmological constant or scalar field. This parametric density model has been tested here with lookback time method and the following estimates (at the $1 \\sigma$ CL) for its characterizing parameters has been obtained\\,:  \\begin{displaymath} q_0 = -0.68_{-0.13}^{+0.21} \\ \\ \\ \\ , \\ \\ \\ \\ h = 0.71 {\\pm} 0.07 \\ , \\end{displaymath} in agreement with the results from the test performed in \\cite{Hobbit}. It is worth stressing that our method turns out to be more efficient than the usual fit to the SNeIa Hubble diagram for the parametric density model. It is also interesting to note that both $q_0$ and $t_0$ predicted by this model are in almost perfect agreement with those computed for the $\\Lambda$CDM model with the parameters $(\\Omega_M, h)$ as discussed above. This is not surprising since Eq.(\\ref{eq: rhor}) shows that, nowadays, the energy density of the parametric density model is very similar to that of the $\\Lambda$CDM model so that these two scenarios share most of the observable properties referred to today quantities. However, this does by no way mean that the two models are the same. Actually, it is the underlying philosophy that is completely different having now a single fluid rather than a cosmological constant dominating over the matter term. Both models predict similar  values ofthe today observed quantities  simply because they are tied to reproduce the same data and not because they are two different manifestations of the same underlying physics.  Another possible approach to the cosmic acceleration is to consider this feature as the first signal of the breakdown of the Einstein General Relativity at some characteristic scale. In this picture, the universe is still dominated by  standard matter, but the Friedmann equations have to be modified as consequence of a different gravity Lagrangian. This philosophy inspired  curvature quintessence scenarios where an effective dark energy is related to the properties of the function $f(R)$ which replaces the Ricci scalar $R$ in the gravity Lagrangian \\cite{capozcurv,curvature,review}. The coupling between matter and curvature for a theory with $f(R) = f_0 R^n$ gives rise to fourth order nonlinear differential equations for the evolution of the scale factor that is not analytically solvable. This difficulty disappears if one considers a toy model with $\\Omega_M = 0$ so that power law solutions, $a \\sim t^{\\alpha}$, are possible provided that $\\alpha$ is linked to $n$ by Eq.(\\ref{eq: alphancurv}). Such a model is particularly attractive from a theoretical viewpoint (since it allows to give a purely geometric interpretation of the dark energy) and has also been shown to fit well SNeIa Hubble diagram \\cite{curvfit}. We have shown here that, while it is able to pass successfully the loockback time test implemented in this paper, only weak constraints can be imposed on the parameter $n$ with the best fit value lying in a region corresponding to decelerating rather than accelerating models. We are thus tempted to conclude that this scenario could be rejected with some confidence. Actually, this conclusion is also suggested by the recent result by Riess et al. \\cite{Riess04} which, using the SNeIa Hubble diagram up to $z = 1.55$, have convincingly detected a change in the sign of the deceleration parameter $q$. The power law solutions predict, on contrary, a constant $q$ so that they are ruled out by the result in Riess et al. \\cite{Riess04}. It is important to stress, however, that this is not a definitive exclusion of the curvature quintessence scenario. Actually, we have only considered power law solutions in absence of matter. It is indeed conceivable that including a matter term completely change the evolution of the scale factor possibly giving a past deceleration followed by the present acceleration thus leading the curvature quintessence model in agreement with what is suggested by the extended SNeIa Hubble diagram (see also \\cite{curvature} and references therein for a further discussion of curvature models).  Having tested three different models, it is worth asking what is the better one. Unfortunately, this is not possible on the basis of the test results only. Comparing the $\\chi^2$ values, one could naively conclude that the parametric density model is the best one since it gives the lowest $\\chi^2$. However, while the difference among the $\\chi^2$ values is significant between the parametric density model and the curvature quintessence scenario (almost an order of magnitude), it is too small (0.07 vs. 0.09) to conclude that the $\\Lambda$CDM model is disfavored. We have thus to conclude that this test alone is not able to discriminate among these two dark energy candidates. On the other hand, Sandvik et al. \\cite{Sandvik} recently claimed that UDE models are not viable because the growth of density perturbations will lead to matter power spectrum in disagreement with what is observed. This should be an evidence against the parametric density model. However, it is worth noting that Sandvik et al. explicitely consider the generalized Chaplygin gas model which is characterized by a negative squared sound speed that seems to be the main cause of the anomalous growth of perturbations. For the parametric density model, the sound speed is always positive definite so that it is likely that the argument of Sandvik et al. should be at least reconsidered.  Finally, we would briefly comment on the possibility to ameliorate our method by increasing the maximum redshift probed. To this aim, galaxy clusters do not appear as good candidates since it is quite difficult to detect a significative number of member galaxies up to redshifts larger than $\\sim 1.3$. However, high redshift galaxies may be taken into account provided that they are detected in as more photometric bands as possible. This latter requirement is fundamental since it allows not only to better estimate the age of the galaxy, but also to infer constraints on its formation redshift $z_F$ that cannot be assumed to be the same for all the galaxies (as it has been possible for clusters). The $i$ drop out technique allows to discover galaxies up to redshift $z \\sim 6$ \\cite{Dick04} and thus let us hope to measure lookback time up to such high redshift. With its ability of both furnishing multicolor photometry of high redshift galaxies (and thus better estimates of their color ages and formation redshift) and increasing the number of SNeIa, the GOODS survey \\cite{GOODS} seems to be the most promising source of cosmological constraints in the near future."
    },
    "0410/astro-ph0410742_arXiv.txt": {
        "abstract": " ",
        "introduction": "The study of the light speed anisotropy with respect to the dipole of the Cosmic Microwave Background  (CMB) radiation as suggested in \\cite{GM}, is the modern analog of the Michelson-Morley experiment.  CMB, besides being a unique cosmological messenger, also is determining the hierarchy of inertial frames  and their relative motions, and is defining an \"absolute\" inertial frame of rest, i.e. the one where the  dipole and quadrupole anisotropies vanish.  The dipole anisotropy of the temperature $T$ of CMB without doubts is of Doppler nature, \\begin{eqnarray} \\frac{\\delta T(\\theta)}{T}=(v/c)\\, cos \\theta + (v^2/2c^2)\\, cos2\\theta + O(v^3/c^3) \\end{eqnarray} (the first term in the right hand side is the dipole term) and is indicating the Earth's motion with velocity $$v/c = 0.000122 \\pm  0.00006; v = 365 \\pm 18\\, km\\, s^{-1},$$ with respect to the above mentioned CMB frame.  WMAP satellite's 1-year data define the amplitude of the dipole $3.346 \\pm 0.017$ mK and the coordinate of the apex of the motion  (in galactic coordinates) \\cite{Ben03} \\begin{equation} l=263.85^{\\circ} \\pm 0.1^{\\circ}, b=48.25^{\\circ}\\pm 0.04^{\\circ}. \\end{equation} The errors are due to the calibration of the WMAP and will be improved in due course. This coordinate is in agreement with estimations based on the hierarchy of motions involving the Galaxy,  Local group and Virgo supercluster \\cite{RG}.  The probing of {\\it the anisotropy of the speed of light with respect to the direction of CMB dipole,} therefore, is a profound aim.  An experiment was proposed in \\cite{GM} to check such a light speed anisotropy using the effect of  inverse Compton scattering of photons on monochromatic electron beams. Originally the idea  of using  the Compton scattering of laser photons on accelerated electrons has been proposed in \\cite{HTM},.  Estimations showed the feasibility for reaching high accuracies at available parameters of laser beams, as well as for highly monochromatic electrons produced in existing accelerators, such as SLAC, ESRF (Grenoble), TJNAF. Besides of methodical difference with other experiments such a test would reach accuracies higher than the existing limits.  Majority of performed measurements of the light speed isotropy were dealing with   a closed path propagation of light (see \\cite{W} and references therein). Such round-trip propagation are insensitive to the first order but are sensitive only to the  second order of the velocity of the reference frame of the device with  respect to a hypothetical universal rest frame.  Mossbauer-rotor experiments yield a one-way limit $\\Delta{c/c}<2\\times 10^{-10}$, using fast beam  laser spectroscopy \\cite{R}.  The latter using the light emitted by the atomic beam  yield a limit $\\Delta{c/c}<3\\times10^{-9}$ for the anisotropy of the one-way  velocity of the light. Similar limit was obtained by Vessot et al. \\cite{V} for the difference in speeds of the uplink and the downlink signals used in the NASA GP-A rocket experiment to test the gravitational redshift effect. One-way measurement of the velocity  of light  has  been  performed using also NASA's Deep Space Network  \\cite{KRISH}: the obtained limits  yield $\\Delta c/c <3.5\\times 10^{-7}$  and $\\Delta c/c <2\\times10^{-8}$ for  linear and quadratic dependencies, respectively. Another class of experiments dealt not with angular but frequency dependence of the speed of  light. For the visible  light and $\\gamma$-rays the limit $(1.8\\pm6)\\times 10^{-6}$ was obtained, while the difference with the velocity of 11~GeV energy electrons $v_e$ was  $(c-v_e)/c=(-1.3\\pm2.7)\\times10^{-6}$ (see \\cite{MAC}). The data on gamma ray bursters, sources  located on cosmological distances, were used by Schaefer \\cite{Brad} to obtain strong limits  for the variation of the speed of light on frequency, namely $\\Delta c (\\nu) /c < 6.3 \\times 10^{-21}$ was obtained  for 30-200 keV photons of the flare GRB 930229.  Below we present the first results of the test of the isotropy of the  one-way speed of light by means of the Compton effect \\cite{GM} using the  data obtained at GRAAL. The results obtained are not only methodically different from those of the above mentioned experiments but also provide stronger constraints  on the light speed anisotropy in CMB frame. ",
        "conclusions": "From the stability limit shown in Figure 6, we established the following constraint for the light speed anisotropy: $ \\Delta c/c \\leq (1 \\pm 0.5) \\times  10^{-12}. $ The systematic errors of the experiment are indeed very difficult to estimate, thus in view of the large and unexplained fluctuations of the fitted value for the CE, we adopt a conservative limit \\begin{eqnarray} \\Delta c/c < 3 \\times  10^{-12}. \\end{eqnarray}  However, even this conservative limit is about 3 orders of magnitude better than the one reachable from measurements involving the space probes (including the Cassini probe), since it will correspond to several meter accuracy for a planetary orbit; also it deals with one-way effect, while the space measurements are two-way ones (Nordtvedt, personal communication, 2004).  It is important to note that, the electrons are circulating inside the machine at  a rate of $10^6$ revolutions per second,  and therefore, an anisotropy effect cannot be compensated by the stabilization procedure of the machine as it could happen in a linear accelerator. In  other words, the angle between the electron speed  inside the ring and the direction of the CMB dipole is changing so fast, that the machine cannot compensate for a change in energy induced by an anisotropy.  Further studies of the light speed anisotropy with respect to the apex of CMB dipole, possibly in dedicated experiments with higher $\\gamma$ and lower $d\\gamma$, seem crucial and can reduce our limit or give evidence for possible anisotropy effects.  We thank K.Nordtvedt for valuable discussions and the ESRF teams for their permanent help, in particular, L.Hardy, J.L.Revol, D.Martin and the members of the alignment group."
    },
    "0410/astro-ph0410097_arXiv.txt": {
        "abstract": "Quadratic methods with heuristic weighting (e.g.\\ pseudo-$C_l$ or correlation function methods) represent an efficient way to estimate power spectra of the cosmic microwave background (CMB) anisotropies and their polarization. We construct the sample covariance properties of such estimators for CMB polarization, and develop semi-analytic techniques to approximate the pseudo-$C_l$ sample covariance matrices at high Legendre multipoles, taking account of the geometric effects of mode coupling and the mixing between the electric ($E$) and magnetic ($B$) polarization that arise for observations covering only part of the sky. The $E$-$B$ mixing ultimately limits the applicability of heuristically-weighted quadratic methods to searches for the gravitational-wave signal in the large-angle $B$-mode polarization, even for methods that can recover (exactly) unbiased estimates of the $B$-mode power. We show that for surveys covering one or two per cent of the sky, the contribution of $E$-mode power to the covariance of the recovered $B$-mode power spectrum typically limits the tensor-to-scalar ratio that can be probed with such methods to $\\sim 0.05$. ",
        "introduction": "\\label{sec:intro}  With the recent results from the Degree Angular Scale Interferometer (DASI; \\citealt{kovac02,leitch04}), the Cosmic Anisotropy Mapper (CAPMAP; \\citealt{barkats04}) and the Cosmic Background Imager (CBI; \\citealt{readhead04}), and the imminent two-year results from the \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP; \\citealt{kogut03}), the characterization of the polarization of the cosmic microwave background (CMB) polarization is gathering pace. The ultimate goal of such experiments is to constrain parameters describing the cosmological model and its fluctuations, and hence test models for the generation of these fluctuations, such as inflation. In Gaussian theories, all of the cosmological information encoded in the CMB anisotropies and their polarization is contained in their power spectra. For this reason, estimating CMB power spectra is a useful and efficient form of data compression in analysis pipelines~\\citep*{bond00}. However, estimating the spectra is not enough: it is also essential to have an accurate covariance matrix, and preferably an accurate model for the shape of the likelihood function (see e.g.~\\citealt{bond00,efstathiou04}). Modelling the covariance matrix for a particularly efficient class of polarization power spectra estimators -- heuristically-weighted quadratic estimators, or pseudo-$C_l$s -- is the subject of this paper.  The polarization field can be decomposed uniquely on the full sky into a gradient part, with electric ($E$) parity, and a curl part with magnetic ($B$) parity~\\citep*{kamionkowski97,seljak97}. The cosmological importance of this decomposition is that linear density perturbations cannot produce magnetic polarization, and so, in standard scenarios and at linear order, large-angle $B$-mode polarization is produced solely by gravitational waves. Second-order effects, most notably weak gravitational lensing by large-scale structure, can generate $B$ modes from primary $E$ modes~\\citep{zaldarriaga98}. The lens-induced $B$ modes have a blue spectrum, and should dominate any primary signal from gravitational waves on sub-degree scales. On larger scales, they may still dominate if the gravitational wave amplitude is sufficiently low. Either way, the variance of the $B$ modes is expected to be much smaller than that of the $E$ modes. Given the different parity properties of the $E$ and $B$ fields, CMB polarization in a statistically-isotropic and parity-respecting universe is characterized by two power spectra, $C_l^E$ and $C_l^B$, and also the cross-power $C_l^{TE}$ between $E$ and the temperature anisotropies.  A number of different methods have been developed to estimate these CMB polarization power spectra. The optimal method, brute-force maximum likelihood (e.g.\\ \\citealt*{gorski94,bond98}), and its quadratic approximation~\\citep{tegmark02}, should produce the most accurate (minimum-variance) estimates. However, these methods require the inversion of $N_{\\rmn{pix}}\\times N_{\\rmn{pix}}$ matrices, where $N_{\\rmn{pix}}$ is the number of pixels in the map, which is prohibitively slow for the analysis of mega-pixel maps. Given these limitations, a number of fast alternative methods have been suggested~\\citep{hansen02b,kogut03,chon04} building on earlier work on the analysis of the temperature anisotropies~\\citep*{szapudi01,wandelt01b,hivon02,hansen02}. These fast methods are closely related, and all involve compressing heuristically-weighted maps of the Stokes parameters into a set of pseudo-$C_l$s making use of spherical transforms -- a step that requires only $O(N_{\\rmn{pix}}^{3/2})$ operations with fast Fourier transform methods on iso-latitude pixelisations such as the widely-used HEALPix~\\citep{gorski04}. The methods differ in the details of how they transform the pseudo-$C_l$s to power spectrum estimates. In the absence of instrument noise, the mean of the pseudo-$C_l$s is linear in the true power spectra, so, for observations over a large enough part of the sky, unbiased power spectrum estimates can be obtained by applying the inverse operation to the pseudo-$C_l$s~\\citep{kogut03}. If the observed part of the sky is sufficiently small, direct inversion is not possible and some form of regularisation is required. \\citet{hansen02b} advocate using likelihood methods, while~\\citet{chon04} make use of correlation functions as an intermediate data product. A notable feature of the latter approach is that the inversion is constructed so that electric and magnetic polarization power is not mixed in the mean. Pseudo-$C_l$ methods have now been applied to a number of real data sets, both for temperature~\\citep{netterfield02,hinshaw03} and polarization analyses~\\citep{kogut03,ponthieu05}.  Despite the attention that pseudo-$C_l$ techniques have received, relatively little work has been done on trying to \\emph{understand} their error properties. This paper aims to fill that gap by providing the basis for a reliable, semi-analytic error analysis technique for CMB polarization power spectrum estimates obtained with pseudo-$C_l$ methods. This extends earlier analytic work that discusses temperature anisotropies only~\\citep{hinshaw03,efstathiou04}, and the recent thesis work by one of us~\\citep{chon_thesis} where both temperature and polarization are considered. A significant complication that arises for polarized data is the mixing of $E$ and $B$ modes for observations covering only part of the sky~\\citep*{lewis02,bunn03}. We develop approximations to the pseudo-$C_l$ sample covariances, taking this mixing into account at leading order for maps with smoothly-apodized edges. Our approximations are accurate in the regime where the mixing is perturbative, i.e.\\ for $l \\gg l_{\\rmn{max}}$, where $l_{\\rmn{max}}$ is the effective band-limit of the function used to weight the Stokes parameters. It is difficult to construct good approximations for the covariance on large scales, $l \\ll l_{\\rmn{max}}$. This is partly because of the increased importance of $E$-$B$ mixing there, but also because the polarization power spectra do not go smoothly to zero on large scales in the presence of reionization. (The behaviour of the temperature-anisotropy power spectrum on large scales produces a similar effect for the temperature pseudo-$C_l$ covariance.) Following~\\citet{efstathiou04}, the low-$l$ sector of the covariance matrices can be constructed directly with exact methods, and used to overwrite the inaccurate approximation. Being able to evaluate accurate covariances efficiently is important in any application where multiple evaluations are required. Examples include parameter estimation, where the covariance generally depends on position in parameter space, and optimising the design and analysis of surveys, where one might wish to assess a large number of weighting schemes.  Throughout, we ignore the effects of instrument noise so we deal only with the sample covariance of the $\\tilde{C}_l$. This is to highlight the geometric effects of heuristic weighting that are peculiar to polarization analysis, particularly $E$-$B$ mixing. An obvious extension of this work that will be required for application to real data is to include the contribution from instrument noise. For uncorrelated noise this should be straightforward, but in a more realistic set-up, one will probably have to resort to Monte-Carlo simulations for the noise contribution.  As noted above, one of the future goals of CMB polarization observations is to detect (or constrain) the stochastic gravitational wave background via its signature in large-angle $B$-mode polarization~\\citep{seljak97,kamionkowski97}. Magnetic polarization allows much smaller amplitude gravitational wave backgrounds to be detected than with temperature or electric polarization alone, since, in the latter case, the gravitational-wave signal is swamped by the cosmic variance of the signal from the dominant density perturbations. Methods to separate out exactly pure magnetic modes from observations with incomplete sky coverage exist at the map level~\\citep{lewis02,bunn03,lewis03}. Such methods have the desirable property that quadratic power spectrum estimates formed from the pure-$B$ modes have no cosmic variance if the $B$-mode power is zero. Were it not for gravitational lensing, separating $B$ modes at the level of the map would thus allow the detection of arbitrarily low amplitude gravitational wave backgrounds in the absence of instrument noise. In practice, imperfections in removing the lensing signal limit the amplitude to tensor-to-scalar ratios\\footnote{Our definition of the tensor-to-scalar ratio follows~\\citet{martin00}, i.e.\\ it is the ratio of the primordial gravitational wave and curvature power spectra.} $r \\sim 3\\times 10^{-5}$ (\\citealt*{kesden02,knox02}; see~\\citealt{seljak03} for refinements from more optimal lensing-reconstruction methods). Heuristically-weighted pseudo-$C_l$ methods do not share the property that the cosmic variance of the estimated $C_l^B$ is immune to $E$-mode power, even if we take care to construct unbiased estimates so that e.g.\\ the mean of our $C_l^B$ estimator decouples from $C_l^E$~\\citep{chon04}. We investigate this issue further here, exploring the limits of detectability for $r$ that arise solely from cosmic variance of the $E$-mode polarization leaking into the variance of quadratic estimates of $C_l^B$.  This paper is organised as follows. Section~\\ref{sec:properties} reviews the properties of the polarization pseudo-$C_l$s and establishes our notation. Following a brief discussion of the exact pseudo-$C_l$ covariance matrices in Section~\\ref{sec:exact-cov}, we develop analytic approximations to them that can be rapidly computed in Section~\\ref{sec:approx-cov}. We begin with a toy model, based on a Gaussian weight function applied to a small survey area in the flat-sky limit. (Some of the more involved manipulations are relegated to Appendix~\\ref{app:flat}.) This instructive example illustrates the key issues that are specific to the analysis of polarization, i.e.\\ the geometric mixing of $E$ and $B$ modes. We then move on to develop approximations for arbitrary (smooth) weight functions, and test these against the exact covariances in cases where we can easily compute the latter (in particular, for azimuthally-symmetric weight functions). Simple rules-of-thumb for band-power variances are also derived in Section~\\ref{subsec:rot} as limiting cases of the more general expressions; some of the manipulations of the $3j$ symbols required for these derivations are relegated to Appendix~\\ref{app:useful_sums}. Finally, in Section~\\ref{sec:implications} we consider the implications of our results for the detection of gravitational waves with $B$-mode polarization estimated with pseudo-$C_l$ methods.  Throughout, we illustrate our results with the best-fitting power-law concordance $\\Lambda$CDM model to the \\emph{Wilkinson Microwave Anisotropy Probe} (\\emph{WMAP}) and 2dF Galaxy Redshift Survey data (\\citealt{spergel03}; Table 7), but with the tensor-to-scalar ratio $r=0.15$ corresponding to $\\phi^2$ inflation. We include the effects of gravitational lensing in the $B$-mode power spectrum, but ignore effects due to its non-Gaussianity. Theoretical power spectra, pseudo-$C_l$s and estimated power spectra in all figures include the smoothing effect of a Gaussian beam of ten-arcmin FWHM. ",
        "conclusions": "\\label{sec:conclusion}  Accurate calculation of the covariance matrix is an important part of power spectrum estimation during the analysis of CMB data. A `correct' estimator (i.e.\\ unbiased), but with poorly quantified errors, is useless for subsequent parameter estimation. In this paper we have constructed the contribution from sample variance to the covariance matrices of heuristically-weighted quadratic estimates of the CMB polarization power spectra. Exact calculation of these matrices for general weighting is numerically feasible at large scales ($l$ less than a hundred or so), but the $O(l^6)$ scaling is prohibitive on small scales. To resolve this problem we developed analytic approximations to the polarization covariance matrices, including the effects of mode coupling and, in a perturbative manner, the mixing of $E$ and $B$ modes that inevitably arises with heuristic weighting. Our approximations, which generalise earlier work on the temperature anisotropies~\\citep{hinshaw03,efstathiou04}, can be computed efficiently to high $l$. Accurate covariance matrices can be obtained on all scales by replacing the low-$l$ sectors of our approximations with a numerical calculation.  We tested the approximations against numerical results for two toy-model surveys for which azimuthal symmetry allows efficient evaluation of the exact covariances. For the more challenging survey, a $15\\degr$-radius survey with cosine-squared apodization of the outer $5\\degr$, the most significant errors were on the variance of $\\tilde{C}_l^B$. These errors arise from our local treatment of the $E$-mode power spectrum when it appears in a convolution with the broad kernel that describes the mixing of $E$ and $B$ modes in equation~(\\ref{eq:5-38}). For the variance of $\\tilde{C}_l^E$ (and the temperature anisotropies) the approximation is more accurate because the dominant convolution involves a kernel that is intrinsically more narrow. It will probably be very difficult to improve the error on the variance of $\\tilde{C}_l^B$ in an analytic treatment. Fortunately, convolution errors tend to cancel in the approximation of the \\emph{correlation} matrices. We observed maximum errors $\\la 0.05$ in the correlations, except on the largest scales, suggesting that, in a practical application, it should be possible to construct accurate covariance matrices by combining the approximate correlations with variances calibrated off simulations. By binning our approximate covariances, we also derived useful rules of thumb for the variances of band-power estimates of $C_l^E$ and $C_l^B$. These are particularly simple to evaluate, involving only integrals of products of the weight function and its derivatives.  A further aim of this paper was to quantify the conditions under which pseudo-$C_l$ methods can be applied to gravitational-wave searches via $B$-mode polarization. Measurements of the $B$-mode power spectrum from gravitational waves with future surveys with sub-$4\\,\\mu\\rmn{K}$-arcmin thermal noise are potentially limited by the sample variance of lens-induced $B$ modes. We showed that for surveys covering one or two per cent of the sky, the additional variance due to leaked $E$ modes dominates the lensing noise. Adopting pseudo-$C_l$ methods would therefore by highly sub-optimal: the error on the tensor-to-scalar ratio would be $\\sim 25$ times higher than if $B$ modes had been isolated before forming quadratic estimates. Ultimately, this additional variance would limit the tensor-to-scalar ratio that is detectable with pseudo-$C_l$ methods to $\\sim 0.05$. Map-level $E$-$B$ separation, or more optimal weighting such as in maximum-likelihood power spectrum estimation, will therefore be essential to ensure the maximum returns from future $B$-mode surveys."
    },
    "0410/math-ph0410004_arXiv.txt": {
        "abstract": "Maxwell's multipoles are a natural geometric characterisation of real functions on the sphere (with fixed $\\ell$). The correlations between multipoles for gaussian random functions are calculated, by mapping the spherical functions to random polynomials. In the limit of high $\\ell,$ the 2-point function tends to a form previously derived by Hannay in the analogous problem for the Majorana sphere. The application to the cosmic microwave background (CMB) is discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410604_arXiv.txt": {
        "abstract": "We present high resolution ($0.''12 \\times 0.''079$) observations of the ultracompact \\hii\\ region G10.6-0.4 in 23~GHz radio continuum and the \\nhthree\\ line. Our data show that the infall in the molecular material is largely spherical, and does not flatten into a molecular disk at radii as small as 0.03~pc. The spherical infall in the molecular gas matches in location and velocity the infall seen in the ionized gas. We use a non-detection to place a stringent upper limit on the mass of an expanding molecular shell associated with pressure driven expansion of the \\hii\\ region. These data support a scenario in which the molecular accretion flow passes through an ionization front and becomes an ionized accretion flow onto one or more main sequence stars, not the classical pressure-driven expansion scenario. In the continuum emission we see evidence for externally ionized clumps of molecular gas, and cavities evacuated by an outflow from the central source. ",
        "introduction": "\\label{sec:intro}  The ultracompact (UC) \\hii\\ region G10.6-0.4, at a distance of 6.0~kpc \\citep{dow80}, is the brightest member of a complex of variously compact \\hii\\ regions \\citep{woo89b} in an area of active star formation. The associated IRAS point source, IRAS 18075-1956, has a luminosity of $9.2 \\times 10^5$~\\lsun\\ \\citep{cas86} and has colors that meet the criteria of \\citet{woo89a} for an \\uchii\\ region. G10.6 is known to be embedded in a hot molecular core (HMC) \\citep{bra83,ho86,plu92}. The core is thought to contain at least 1200~\\msun\\ of gas within a radius of 0.2~pc, based on an analysis of a variety of dust continuum measurements \\citep{mue02}, and at least 3300~\\msun\\ of gas within 1.1~pc based on \\ceio\\ and \\cseo\\ measurements \\citep{hof00}. Previous studies of the inversion lines of \\ammonia\\ have determined that the molecular core is rotating and collapsing inward toward the \\uchii\\ region \\citep{ho86,ket87a,ket88,ket90}. In these studies, using the \\nhone\\ and \\nhthree\\ lines, rotation is seen at size scales from 1~pc down to 0.08~pc, and infall is detected in the form of red-shifted absorption seen against the continuum source. \\methanol\\ and \\htwoo\\ masers are seen distributed linearly in plane of the rotation \\citep{wal98,hof96}, while OH masers seem to be distributed along the axis of rotation \\citep{arg00}. In \\ceioone, \\citet{ho94} see $10^3$~\\msun\\ of dense (n$\\sim10^6$\\percc), rotating gas in a flattened (0.3$\\times$0.1~pc) disk-like structure. At the highest resolution achieved in earlier work, infall and rotation in the molecular gas were seen simultaneously in absorption, showing that the molecular gas was spiraling inward on size scales comparable to the size of the \\uchii\\ region. Our new observations resolve the \\uchii\\ region and show clearly that the molecular gas is infalling mostly spherically toward the \\uchii\\ region, with only slow rotation and little flattening in the plane of rotation.  Recent observations of the ionized gas within the \\uchii\\ region showed that the ionized gas is also spiraling inward toward the stars at the center of the \\uchii\\ region \\citep{ket02a}. Subsequent theoretical work showed that in small \\hii\\ regions, the gravitational effect of the central star(s) can overcome the thermal pressure of the ionized gas causing the molecular accretion flow to pass through the \\hii\\ region boundary and continue inward as an ionized accretion flow \\citep{ket02b}. In this model, the \\hii\\ region boundary exists as a standing R-type ionization front within a continuous accretion flow. These results differ from classical treatments of the pressure driven expansion of \\hii\\ regions, which predict outward motion of the ionized gas as soon as the \\hii\\ region is formed \\citep{str39,spitzer}. In the classical model for pressure driven expansion, the \\hii\\ region boundary, after a very short phase as a moving R-type front, will develop a characteristic double front structure composed of an isothermal shock followed by a moving D-type ionization front. In this model, as the \\hii\\ region expands, most of the displaced molecular material remains between the shock and the ionization front as a dense outward moving shell, which snow-plows ahead of the \\hii\\ region. In the alternative model of \\citet{ket02b}, however, the accretion flow passes through a standing R-type ionization front at the \\hii\\ region boundary and continues toward the star(s) as an ionized flow. In that case there will be no swept-up, dense molecular layer at the boundary, and all the molecular gas will be moving inward. Ionized accretion flows represent a fundamentally different mode of accretion than that seen in low-mass star formation. In the ionized accretion flow scenario a main sequence star accretes gas which was once part of a molecular accretion flow, but has been photo-ionized as it approaches the star. G10.6 contains the first example of evidence for a massive accretion flow continuing through an ionization front. Our new observations constitute a stringent non-detection of any outward moving molecular layer thus contradicting the pressure driven expansion model (see Section \\ref{sec:nature}). ",
        "conclusions": ""
    },
    "0410/physics0410255_arXiv.txt": {
        "abstract": "{ We present the timing and spectral properties of the transient Be/X-ray binary pulsar 3A~0535+262 during quiescence using three observations with the narrow field imaging instruments (NFI) of BeppoSAX. Assuming a distance of 2 kpc for this system, the 2--10 keV X-ray luminosities measured from the three observations are in the range of 1.5--4.0 $\\times$ 10$^{33}$ erg s$^{-1}$, indicating a very low rate of accretion. We report the detection of pulsations at a very low luminosity of 2 $\\times$ 10$^{33}$ erg s$^{-1}$ during one of the three observations, though at this accretion rate the system is expected to be in the centrifugally inhibited regime. The X-ray spectra for the unpulsed observations are best modeled as power law type while a combined model of power law and black-body is required to fit the pulsed spectrum.} ",
        "introduction": "3A~0535+262 is a close binary system which contains a O9.7 IIIe Be star (HDE 245770, Giangrande et al. 1980) and a slowly rotating, strongly magnetized neutron star. The X-ray source was discovered by the Ariel V satellite (Rosenberg et al. 1975) when it was in a `giant outburst' mode with X-ray intensity of the order of $\\sim$ 2 Crab in the 3--7 keV range and showed pulsations at 104 s. The system has been observed to undergo several transient X-ray outbursts classified as ``giant'' (2--10 keV flux $>$ 1 crab) and ``normal'' (2--10 keV flux $< 1$ crab) outbursts. After 1975 two more giant outbursts were observed, one in 1980 detected with the Hakucho satellite (Nagase et al. 1982) and another one in 1994 by the BATSE detectors onboard CGRO (Finger et al. 1996). A large orbital period of 111.0$\\pm$0.4 days was discovered from the long term light curves obtained with the Vela 5B satellite (Priedhorsky and Terrell 1983). Interestingly, the system shows significant periodic optical variability at a different period of $\\sim$ 103 days (Larionov, Lyuty \\& Zaitseva 2001). During the 1994 outburst, quasi-periodic oscillations (QPO) were observed with frequency correlated to the X-ray flux and spin-up torque (Finger et al. 1996, Ghosh 1996). The orbital parameters were also determined from observations made with BATSE during this outburst. The magnetic field strength of the neutron star in 3A~0535+262 is $\\sim$ 1 $\\times$ 10$^{13}$ Gauss as measured from the cyclotron resonance scattering feature with the OSSE experiment onboard CGRO during the 1994 outburst (Grove et al. 1995).  With EXOSAT observations made in 1985-1986 in between the outbursts, a moderate and pulsating X-ray flux was detected from 3A~0535+262 at various orbital phases (Motch et al. 1991). The X-ray flux in the 1--20 keV energy band was measured to be in the range of 1.5--3.0 $\\times$ 10$^{-10}$ erg cm$^2$ s$^{-1}$, which corresponds to a luminosity of 0.7--1.4 $\\times$10$^{35}$ erg s$^{-1}$ for a distance of 2 kpc (Steele et al. 1998). The detection of X-ray pulsations in two of the EXOSAT observations showed that the mass accretion rate was large enough not to cause centrifugal inhibition and that material from the outflow of the Be star companion accretes through a disk onto the neutron star. In more recent observations of 3A~0535+262 made with the Proportional Counter Array of the Rossi X-ray Timing Explorer (RXTE-PCA), the source was detected at a much lower luminosity of 2.0--4.5 $\\times$ 10$^{33}$ erg s$^{-1}$ at which the onset of centrifugal inhibition of accretion or propeller effect is expected. The X-ray luminosity that was still detected in the propeller regime was ascribed to material leaking through the magnetosphere or thermal emission from the heated core of the neutron star (Negueruela et al. 2000). However, as the RXTE-PCA detectors are non-imaging instruments and therefore have large internal and X-ray sky background, the flux and spectral measurements at low intensity are not very accurate.  In this paper, we present X-ray timing and spectral properties of 3A~0535+262 in a quiescent state from three observations made with the Beppo-SAX narrow field imaging spectrometer instruments. Detection of pulsations and measurement of the X-ray spectrum in quiescent state may resolve the outstanding issue regarding the different accretion regimes in which the system would reside. In the subsequent sections we describe the observations (\\S2), results from timing and spectral analysis (\\S3) and a discussion on the propeller effect in this pulsar and the X-ray spectral properties in the pulsating and non-pulsating states (\\S4) based on the results obtained here. ",
        "conclusions": "We have investigated the timing and spectral properties of the X-ray pulsar 3A~0535+262 in quiescent state using the narrow field instruments of the BeppoSAX observatory. We have detected pulsations in observation C in parts of which the 2--10 keV X-ray luminosity was as low as 2 $\\times$ 10$^{33}$ erg s$^{-1}$. For all the three observations of 3A~0535+262 the accretion is expected to be centrifugally inhibited though the detection of pulsations at these low flux levels in one of them indicates that some matter may have leaked through onto the neutron star surface. When the emission is non-pulsating, the X-ray spectrum is power law type. The origin of an additional black-body component in the pulsed observation is not very clear though it seems to be essential in fitting the spectra."
    },
    "0410/astro-ph0410432_arXiv.txt": {
        "abstract": "{This paper presents the X-ray spectroscopy of an X-ray selected sample of 25 radio-loud (RL) AGNs extracted from the XMM-{\\it Newton} Bright Serendipitous Survey (XBSS). The main goal of the work is to assess and study the origin of the X-ray spectral differences usually observed between radio-loud and radio-quiet (RQ) AGNs. To this end, a comparison sample of 53 RQ AGNs has been also extracted from the same XBSS sample and studied together with the sample of RL AGNs. Since there are many claims in the literature that RL AGNs have, on average, a flatter spectral index when compared to the RQ AGNs, we have focused the analysis on the distribution of the X-ray spectral indices of the power-law component that models the large majority of the spectra in both samples. We find that the mean X-ray energy spectral index is very similar in the 2 samples and close to $\\alpha_X\\sim1$. However, the intrinsic distribution of the spectral indices is significantly broader in the sample of RL AGNs. In order to investigate the origin of this difference, we have divided the RL AGNs into blazars (i.e. BL Lac objects and FSRQs) and ``non-blazars'' (i.e. radiogalaxies and SSRQs), on the basis of the available optical and radio information. Although the number of sources is small, we find strong evidence that the broad distribution observed in the RL AGN sample is mainly due to the presence of the blazars. Furthermore, within the blazar class we have found a link between the X-ray spectral index and the value of the radio-to-X-ray spectral index ($\\alpha_{RX}$) suggesting that the observed X-ray emission is directly connected to the emission of the relativistic jet. This trend is not observed among the ``non-blazars'' RL AGNs. This favours the hypothesis that, in these latter sources, the X-ray emission is not significantly influenced by the jet emission and it has probably an origin similar to the RQ AGNs. Overall, the results presented here indicate that the observed distribution of the X-ray spectral indices in a given sample of RL AGNs is strongly dependent on the amount of relativistic beaming present in the selected sources, i.e. on the relative fraction of blazars and ``non-blazars''. ",
        "introduction": "The origin of the differences between radio-loud (RL) and radio-quiet (RQ) AGNs is one of the fundamental questions in the AGN phenomenology. The broad and possibly bimodal distribution of the radio-to-optical flux ratios observed within the class of AGNs (e.g \\citealt{kellermann89}; \\citealt{dellaceca94}; \\citealt{ivezic02}; but see also \\citealt{white00}; \\citealt{cirasuolo03} for evidences against the bimodality) is not yet fully understood and it may reflect important differences between the physical processes at work in the inner regions of the 2 classes of objects. It has been proposed that the ``radio-loudness\" of the AGNs could be directly connected to the mass of the central black-hole (\\citealt{franceschini98}; \\citealt{laor00}) or to its spin (e.g. \\citealt{meier99}). Alternatively, a different accretion rate (e.g. \\citealt{woo02}; \\citealt{ho02}) or a different structure of the accretion disk (e.g. see \\citealt{ballantyne02} and references therein) may be at the origin of the observed phenomenology.\\newline In this context, X-ray observations can play a key role  since they bring information on the innermost region of the AGN structure. In particular, a direct comparison between the X-ray properties of RL and RQ AGNs can shed light on the origin of the ``radio-loudness\".  If the main properties of the X-ray spectrum of RQ AGNs are quite well established, RL AGNs are, in general, less studied. Early observations carried out with {\\it Einstein} have suggested that RL AGNs have, on average, a flatter spectrum than RQ AGNs (e.g. \\citealt{wilkes87}; \\citealt{shastri93}). This difference is more evident if only RL AGNs with a flat radio spectrum (i.e. the Flat Spectrum Radio Quasars, FSRQs) are considered, while it becomes  marginal when only the AGNs with a steep radio spectrum (i.e. the Steep Spectrum Radio Quasars, SSRQs) are compared to the RQ AGNs (\\citealt{canizares89}). This trend has been subsequently confirmed with ASCA and BeppoSAX data: on the one hand, the X-ray spectrum of a sample of RL AGNs (mostly composed by FSRQs) studied by \\citet{reeves97} and \\citet{reeves00} turned out to be significantly flatter, on average, when compared to RQ AGNs; on the other hand, this difference in the X-ray spectral index is not clearly present in the samples studied by \\citet{sambruna-er-m99} and \\citet{hasenkopf02}, mostly composed by radiogalaxies or lobe-dominated quasars.  Apart from the spectral index, other differences in the X-ray spectra of RL and RQ AGNs include a weaker ``reflection'' hump and Fe I emission line in RL AGNs (e.g. \\citealt{hasenkopf02}; \\citealt{eracleous00}; \\citealt{sambruna-er-m99}; \\citealt{wozniak98}). These results have been interpreted as related to a different structure of the accretion disk in RL AGNs (e.g. \\citealt{ballantyne02}). Moreover, \\citet{sambruna-er-m99} have pointed out that among RL AGNs, the presence of a thermal component (from hot plasma related to the host galaxy and/or to a cluster of galaxies) is very common.  Altogether, the data collected so far seem to indicate that there are significant differences in the X-ray spectrum of RL and RQ AGNs although the details of these differences depend on the radio spectral index of the source. In particular, the presence of a significant contribution to the X-ray emission coming from the relativistic jets can play an important role in some RL AGNs (like the FSRQs).  Chandra and XMM-{\\it Newton} can give a fundamental contribution to this topic: different components of the X-ray emission in RL AGNs can be resolved spatially, with Chandra, (e.g. \\citealt{gambill03} or \\citealt{schwartz04} for a review) or spectrally, with XMM-{\\it Newton} (e.g. \\citealt{brocksopp04},  \\citealt{ferrero03}, \\citealt{page04}). In this paper, we present a systematic X-ray spectral analysis of an X-ray selected sample of 25 RL AGNs extracted from the XMM-{\\it Newton Bright Serendipitous Survey} (XBSS, \\citealt{dellaceca04}). In this sample we have distinguished between the AGNs with flat and steep radio spectrum in order to establish the importance of the jets in the X-ray emission of RL objects. In order to assess how different the selected RL AGNs are with respect to the RQ AGNs, we have also selected from the same XBSS a ``comparison'' sample of 53 objects, mostly composed by RQ AGNs. Since the two samples have a common origin, the only difference being the radio emission, we expect that any observed difference could be connected to the different mechanisms at work in the 2 classes of AGNs.\\newline Throughout this paper H$_{0}$=65 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}$=0.7 and $\\Omega_{m}$=0.3 are assumed. ",
        "conclusions": "We have presented the analysis of a sample of RL AGNs extracted from the XBSS survey. The NVSS data  have been used to select 28 sources of the sample that are also radio emitting at 1.4~GHz with a flux larger than 2.5~mJy. Out of these 28 radio/X-ray matches, 25 are classified as radio-loud AGNs on the basis of the radio-loudness parameter (R$>$10). The presence of an AGN is detected in the optical and/or in the X-ray spectrum for 23 out of the 25 sources. In the remaining 2 sources, optically classified as ``galaxy'' and ``cluster'' respectively, and for which the X-ray spectrum is modelled by a thermal emission, the presence of an AGN (a radiogalaxy) is strongly suggested by the radio properties (radio-to-optical flux, the radio power or morphology). Therefore, the sample studied here is fully composed by AGN.  Taking into account the flux limit of the NVSS (2.5 mJy/beam), we have estimated that the RL AGNs represent 13\\% of the X-ray selected AGNs at a flux limit of $\\sim$7$\\times$10$^{-14}$ erg s$^{-1}$ cm$^{-2}$. We have then selected from the same XBSS a ``comparison'' sample of 53 sources not detected in the NVSS survey, i.e. with a radio flux density below 2.5 mJy/beam. We have estimated that the large majority of the sources in this sample is represented by RQ AGNs and that at most 2 objects could be RL objects.  The X-ray spectra of the sources in both samples have been extracted, studied and compared.  The main results can be summarized as follows:  \\begin{itemize}  \\item The spectra of the majority of the sources in both samples ($\\geq$90\\%) are well described by a single power-law. The mean spectral index is similar in both samples ($\\Gamma\\sim$2).  \\item  Although the mean value of $\\Gamma$ is similar, the intrinsic spread of the values of  $\\Gamma$ (i.e. the one computed taking into account the uncertainties) is significantly different in the two samples, with the RL AGNs having a broader distribution ($\\sigma$=0.34) when compared to the RQ AGNs ($\\sigma$=0.18).  \\item When blazars (BL Lacs and FSRQs) and ``non-blazars'' (radiogalaxies and SSRQs) in the RL AGN sample are analysed separately, the larger spread of the spectral index is observed only in the blazars sample while the ``non-blazars'' show a distribution fully consistent with the one observed in the RQ AGNs. In particular, BL Lacs and FSRQ are populating respectively the steep and the flat tails of the  distribution of $\\Gamma$.  \\item The values of the X-ray spectral index of blazars are correlated with the type of SED, i.e. with the radio-to-X-ray spectral index ($\\alpha_{RX}$). Such a correlation is not observed in the RL AGNs classified as ``non-blazars''. This result suggests that the X-ray emission in the blazar class is directly linked, as expected, to the jet emission while this is probably not the case for the ``non-blazars'', i.e. those RL AGNs whose relativistic jet is not oriented towards the observer. In these objects the relativistic beaming is not important and the observed X-ray emission is likely to have the same origin as in the RQ AGNs.  \\end{itemize}  In conclusion, the analysis presented here has shown that the distribution of the X-ray spectral indices observed in a sample of RL AGNs strongly depends on the importance of relativistic beaming in the selected sources, i.e. on the type of RL AGNs included in the sample under analysis. First, the inclusion of a significant fraction of ``oriented'' RL AGNs, i.e. of blazars, significantly broadens the distribution. Second, the type of blazars included in the analysis changes the observed distribution of X-ray spectral indices: if the sample selection favours objects with low values of $\\alpha_{RX}$ (like in a typical X-ray selected survey), the distribution will be more populated towards steep X-ray indices. Samples that include more sources with large values of $\\alpha_{RX}$ (like in a typical radio survey) are likely to include more blazars with a flat X-ray spectral index."
    },
    "0410/astro-ph0410481_arXiv.txt": {
        "abstract": "We examine the disc--jet connection in stellar mass and supermassive black holes by investigating the properties of their compact emission in the hard X-ray and radio bands.  We compile a sample of $\\sim$100 active galactic nuclei with measured mass, 5 GHz core emission, and 2-10 keV luminosity, together with 8 galactic black holes with a total of $\\sim$ 50 simultaneous observations in the radio and X-ray bands. Using this sample, we study the correlations between the radio ($L_{\\rm R}$) and the X-ray ($L_{\\rm X}$) luminosity and the black hole mass ($M$).  We find that the radio luminosity is correlated with {\\em both} $M$ and $L_{\\rm X}$, at a highly significant level. We show how this result can be used to extend the standard unification by orientation scheme to encompass unification by mass and accretion rate. ",
        "introduction": "Some galaxies are known to emit radiation with extremely high luminosities in the $\\gamma$-ray, X-ray, UV and radio continuum from a very concentrated volume in the nuclear region. Such active cores are the so-called Active Galactic Nuclei (AGN) and their radiation is believed to be produced by accretion onto a supermassive black hole.  The intrinsically complex nature of such systems and the differences in the terminology among different scientific communities (radio, optical, X-ray astronomers) has led to an extremely complicated nomenclature for the AGN zoo.  As the wealth of observations piled up, and with them the number of different AGN types, the opposite enterprise of finding unification schemes has progressively gained support. The basic idea behind the standard unification scheme is that AGN are asymmetric and anisotropic systems. This is natural, as all rotating systems necessarily single out a preferential axis in space and break the full spherical symmetry of non-rotating bodies. Therefore, the orientation of the AGN rotation axis with respect to our line of sight becomes another important parameter that can cause apparent observational differences in two sources that are intrinsically identical ({\\it unification by orientation}).  According to the current widely accepted paradigm (Urry and Padovani 1995), there are two principal causes of anisotropic radiation: obscuration and relativistic beaming.  The former is usually associated with a torus of gas and dust obscuring the optical, UV and (sometimes) soft X-ray radiation along some line of sights, the latter with outflows of energetic particles (jets) along the symmetry axis: the high velocity plasma in the jet beams radiation relativistically in the forward direction.  The powerful source of radiation and the launching site of the collimated jet lies in the center of the system.  Observationally, jet morphologies and spectral properties of both radio and X-ray cores are remarkably similar in the case of black holes of stellar mass (Galactic black holes, hereafter GBH) and of their supermassive counterparts in the nuclei of galaxies (hereafter SMBH).  If jets are launched in the innermost parts of the accretion flows, as commonly assumed, then these similarities suggest that it should be possible to extend the unification scheme further and to understand the physics of both black hole accretion and jet production by studying all those systems as a single class ({\\it unification by mass}).  Furthermore, the recent discovery that SMBH lie at the center of the majority (possibly all) galactic nuclei, even in apparently inactive ones (Kormendy and Richstone 1995; Magorrian et al. 1998), naturally leads to the idea that a grand unification is possible by taking into account the differences in fueling rates among different objects ({\\it unification by accretion rate}).  In the following, we will show how to proceed quantitatively towards a grand unification of active black holes by studying the multivariate correlation among masses, radio luminosities and hard X-ray luminosities of objects traditionally classified in the more diverse ways. The underlying theoretical assumption is that both accretion and jet production are fundamental manifestation of black hole activity, and are somehow physically connected (Begelman et al. 1984; Rawlings and Saunders 1991; Falcke and Biermann 1995; Heinz and Sunyaev 2003). ",
        "conclusions": ""
    },
    "0410/astro-ph0410162_arXiv.txt": {
        "abstract": "We present X--ray observations of nine Weak-Line Radio Galaxies (WLRGs), optically classified as confirmed or possible Low Ionization Nuclear Emission--line Regions (LINERs).  The data were taken from the \\chandra, \\xmm, and \\sax\\ archives.  The \\chandra\\ images typically show complex X-ray morphologies, with hard (2--10 keV) point sources embedded in diffuse soft (0.3--2.0 keV) emission in all cases except 1246--41 (NGC 4696), where only diffuse emission is detected on the scale of the cluster, and 0334--01 (3C 15), where only a point source is detected.   The nuclear X--ray spectra are well fitted at hard energies by an absorbed power law, with a wide range of photon indices, $\\Gamma=1.5-2.7$. Excess absorption over the Galactic value is detected in 6/9 sources, with column densities N$_H \\approx 10^{21-22}$ \\nh.  A thermal component is required at softer energies, in agreement with the results of the spatial analysis. We find that there is no correlation between the core X--ray luminosity and the radio core dominance parameter, suggesting that the bulk of the core X-ray emission is not beamed, but rather is isotropic and thus likely related to the accretion flow. In an attempt to constrain the nature of the accretion flow, we calculate the ratios of bolometric to Eddington luminosities $L_{bol}/L_{Edd}$, and the radiative efficiency $\\eta$ based on the Bondi accretion rates. We find that $L_{bol}/L_{Edd} \\sim 10^{-4}-10^{-6}$ and $\\eta \\sim 10^{-2}-10^{-6}$ for all the objects in our sample, suggesting radiatively inefficient accretion flows. ",
        "introduction": "The recent discovery of \\hst\\ and ground--based observations that most nearby galaxies harbor supermassive black holes \\citep{kor95,ric98}, along with independent evidence that weak nuclear activity is very common in both spiral and elliptical galaxies \\citep{hec80,sad89,tad98}, affirms the fundamental link between ``normal'' and active galaxies.  According to the Palomar survey \\citep{ho97}, over 40\\% of nearby galaxies ($z\\approx0$) contain low--power Active Galactic Nuclei (AGN), with typical bolometric luminosities lower than $10^{44}$ \\lum.  Most of these, based on their low--ionization spectra, are classified as Low Ionization Nuclear Emission--line Regions (LINERs), or transition objects.  Since accreting supermassive black holes are widely held to be responsible for the nuclear activity in AGN, one naturally wonders why such intense nuclear activity is limited to a minority of galaxies.  A possible explanation is that obscuring material along the line of sight conceals the nuclear activity.  Alternatively, there could be inefficient accretion onto the black hole due either to a low accretion rate (the AGN--starved scenario), or to some kind of radiatively--inefficient accretion flow \\citep[RIAF; see e.g.][]{qua03}.  Weak AGN, such as those hosted by nearby low--power radio galaxies in giant ellipticals, may represent the link between powerful AGN and ``normal'' galaxies.  One class of objects that may represent this link are Weak--Line Radio Galaxies \\citep[WLRGs;][]{tad98}, which are classified by [O$_{~\\rm{III}}$]$\\lambda$5007 luminosities more than 10 times lower than those of the general population at similar redshifts.  The weakness of the [O$_{~\\rm{III}}$]$\\lambda$5007 emission line in WLRGs raises the possibility that these objects may be related to LINERs; indeed, there is evidence that WLRGs and LINERs are intimately related \\citep{lew03}.  If all LINERs can be regarded as true AGN, then they would constitute approximately 75\\% of all AGN \\citep{ho97}.  Put in this light, WLRGs/LINERs could prove essential to understanding the link between ``normal'' and active galaxies.  X--rays provide us with a powerful tool by which we can achieve progress in our understanding of the nature of accretion in low-power AGN, because they originate in the innermost, hottest parts of the accretion flow. In addition, soft X--rays map the distribution of the circumnuclear gas around the nuclei, allowing us to quantify the fuel reservoir on the nucleus.  In Section 2 we describe our sample of WLRGs/LINERs.  In Section 3 we describe the observations and data reduction.  Section 4 contains the results of the spatial analysis based on \\chandra\\ observations (4.1), and the results of the temporal and spectral analysis (Sections 4.2 and 4.3, respectively), both of which are based on \\chandra, \\xmm, and \\sax\\ observations.  In Section 5 we discuss the obscuration and nature of accretion in WLRGs/LINERs, as well as some comparisons with the results of previous work.  Finally, in Section 6 we summarize our results. Throughout the paper we use a Friedman cosmology with $H_0={\\rm 75~km~s^{-1}}$ and $q_0=0.5$. ",
        "conclusions": "This paper seeks to shed light on the role of WLRGs/LINERs in the context of uncovering the relationship between ``normal'' galaxies and powerful AGN.  The results of the X--ray analysis allow us to address several important questions towards this end: (1) What is the role of obscuration in WLRGs/LINERs? (2) What is the nature of accretion in such objects? (3) How do the X--ray properties of WLRGs/LINERs compare to those of other classes of radio galaxies?  \\subsection{The role of obscuration}  According to the unified model of AGN \\citep[e.g.,][]{urr95}, the central regions of AGN appear to contain a molecular torus that prevents the penetration of radiation from the nucleus along certain lines of sight.  This can cause intrinsically similar AGN to look remarkably different from different viewing angles.  Recently, however, the presence of an obscuring torus in low--power radio galaxies, specifically FRIs, has been called into question \\citep[henceforth C99C02]{chi99,cap02}, based on \\hst\\ observations of the 3C and B2 catalogs.  Since our sample is composed of low--power radio galaxies, and because our analysis is focused on the X--rays, we are in a position to contribute to this debate.  X--rays, due to their high penetrating power, are the ideal means by which the presence of an obscuring torus can be investigated.  The results of our spectral analysis show that 4 of the 9 objects in our sample require an absorption component on the order of $10^{22}$ \\nh, which qualifies as significant absorption.  Out of those 4 objects, 3 are confirmed LINERs, while 1 is classified as conflicting. In addition to these 4 objects there are 2 others that, while not heavily absorbed, require an absorption component on the order of $10^{21}$ \\nh; one of these additional objects is a confirmed LINER, the other is a possible LINER. Thus, 6 of 9 objects in our sample show signs of absorption, which suggests that some low--power radio galaxies may indeed house obscuring tori.  Moreover, all 6 of these absorbed objects are classified as FRIs or FRI/FRIIs , which is in contrast to the hypothesis that low--power radio galaxies, specifically FRIs, lack a molecular torus (C99C02).  The apparently surprising result that the only FRII in our sample (1637--77) is not absorbed may be explained if the object can be classified as a Low Excitation Galaxy (LEG); Chiaberge, Capetti, \\& Celotti (2002) explain that, while LEGs may be classified morphologically as FRIIs, their radio to optical nuclear properties are indistinguishable from FRIs.  The [O$_{~\\rm{II}}$]$/$[O$_{~\\rm{III}}$] ratio for 1637--77 places it right on the border of the LEG classification, but if 1637--77 can be classified as an LEG, we should expect it to be unabsorbed.  In order to address the apparent discrepancy between the optical/UV and the X--ray absorption properties, one must keep in mind that the inner regions of the accretion flow ($\\sim10-50~{\\rm{R_G}}$), where most of the radiation is supposed to originate, is inaccessible to any observatory. Even though the angular resolution of \\hst\\ ($\\sim0.1$ arcseconds) is much better than \\chandra\\ ($\\sim1.0$ arcseconds) and \\xmm\\ ($\\sim10$ arcseconds), the region probed by \\hst\\ is not closer to the black hole than the X--ray satellites.  Indeed, assuming a typical black hole mass of $5\\times10^8~{M_ {\\odot}}$ at a distance of $40~\\rm{Mpc}$ (corresponding to $z\\sim0.01$), the inner $50~{\\rm{R_G}}$ subtends a region of angular size $\\sim10^{-4}$ arcseconds, which is three orders of magnitude smaller than the spatial resolution of \\hst.  This indicates that the unresolved X--ray and optical sources simply reflect the PSF of the respective instrument rather than the physical size of the emitting regions.  In light of these considerations, it is entirely possible that the X--ray emission is produced in a region closer to the black hole than that of the optical/UV emitting region.  Indeed, theoretical arguments predict that the location of the X--ray emitting region is closer to the black hole in both accretion and jet models.  This hypothesis is supported by the difference in absorption properties found in the optical/UV and the X-rays, if it is assumed that two distinct media are responsible for the absorption \\citep{wei02,sam03,lew03}.  In this theory the optical/UV absorber is a dusty medium lying farther away from the nucleus than the X--ray absorber, which is a neutral or ionized gas largely free of dust.  Further support for the hypothesis that the X--rays are produced in a region interior to the optical/UV emitting region comes from the recent work of Donato, Sambruna, \\& Gliozzi (2004).  Those authors find that a substantial fraction of the compact cores detected in the optical \\citep{chi99}, do not have an unresolved X--ray counterpart, which suggests a different origin and/or nature for optical/UV and X-ray radiation.  Additionally, Donato et al. show that neither X--ray obscuration nor the detection of unresolved X--ray cores is related to the presence of dust lanes, which are responsible for the optical absorption. Finally, we note that there is no discrepancy between this paper and Donato et al. (2004) concerning obscuration; in that paper all the FRIs except the LINER 1216+06 (NGC 4261) show negligible absorption (N$_H<10^{21}~\\rm{cm^{-2}}$).  This suggests that the absorption mechanisms in WLRGs/LINERs may be different than those in other low--luminosity radio galaxies.   \\subsection{The nature of accretion}  The origin of X--rays in radio galaxies is a topic of considerable debate.  Specifically the question is: Do X--rays have their origin in the process of accretion or at the unresolved base of the jet?  The correlation between the radio and X--ray core luminosities in low and high luminosity radio galaxies \\citep[e.g.,][]{fab84,can99,wor94,har98} has often been used as an argument in favor of a common origin at the base of the jet.  However, this correlation has recently been shown by Merloni, Sebastian, \\& Di Mateo (2003) to derive from a more general correlation involving both the X--ray luminosity and the black hole mass---the so--called ``fundamental plane''---in which the radio is related to the jet, and the X--rays are related to inefficient accretion onto the black hole.  This idea is supported by the detailed study of 1216+06 (NGC 4261) by Gliozzi et al. (2003)---their argument is summarized by three main points: (1) they found there to be a high Bondi accretion rate along with a low AGN luminosity, suggesting the presence of a radiatively inefficient accretion flow; (2) they observed an Fe line at $\\approx7$ keV of nuclear origin; (3) they studied the temporal and spectral variability, and combined this with independent information on jet energetics, which lead them to conclude that the bulk of the X--ray emission originates from a radiatively inefficient accretion flow, with negligible jet contribution.  Based on the similarity between 1216+06 and the other objects in our sample, we explored the possibility that the X--rays for all the objects in our sample are related to accretion and not the base of the jet.  One way to explore the origin of nuclear emission is via the correlation (or lack thereof) with the radio core dominance ($R=L_{core}/L_{lobe}$)---this has been done in the optical by Kharb \\& Shastri (2004).  By demonstrating a strong correlation between optical point--like emission and radio core dominance, Kharb \\& Shastri found the optical emission to be strongly beamed, and thus produced by the jet, confirming the results of Chiaberge, Capetti, \\& Celotti (1999). We adopted this approach for our X--ray sample (see Fig.~\\ref{isoemission}, Table~\\ref{tbl-6}), but found there to be no correlation between the core X--ray emission and the radio core dominance, indicating that the emission is not beamed, but rather is isotropic and thus likely to be related to accretion.  On the other hand, there is some indication that in non--LINER FRIs at least a fraction of the X--ray emission is beamed, and thus related to the base of the jet \\citep{donprep}.  Many of the objects in our sample require a powerlaw component for their spectral fit, but the powerlaw components (ranging from $\\Gamma\\sim1.5$ to $\\sim2.1$, excluding the \\sax\\ observation with $\\Gamma\\sim2.7$) can be a good representation for both jet dominated and accretion dominated emission.  The lack of a strong \\feka\\ line suggests the lack of a standard accretion disk; however, both the radiatively inefficient accretion and the jet scenario can account for this result.  And while temporal analysis can yield insights into the nature of accretion, our results yield no such clues.  More insightful information comes from the calculations of the bolometric and Eddington luminosities.  We used estimates of the black hole masses based on stellar velocity dispersions to calculate $L_{{\\rm Edd}}$, and the relationship $0.1L_{{\\rm bol}}=L_{\\rm X}$ in order to calculate $L_{{\\rm bol}}$.  For Seyfert--like objects, which are thought to have a standard thin--disk, the ratio of $L_{{\\rm bol}}$ to $L_{{\\rm Edd}}$ is $\\sim10-20\\%$; an Eddington ratio much less than 10\\% could signify inefficient accretion.  Some accretion properties are listed in Table~\\ref{tbl-5}.  In our sample we find that the majority of sources have values of $L_{{\\rm bol}}/L_{{\\rm Edd}}$ between $10^{-4}$ and $10^{-7}$.  Even the \\xmm\\ observation of 0915--11, with its high X--ray luminosity, has $L_{{\\rm bol}}/L_{{\\rm Edd}}\\sim10^{-3}$.  This is a model--independent indication of Radiatively Inefficient Accretion Flow (RIAF).  Using the results of our spatial analysis, specifically the gas density profile obtained by the de--projection of the surface brightness profile, we calculated \\.M$_{{\\rm Bondi}}$, which serves as a rough estimate of accretion onto the black hole.  With {\\it \\.M}$_{{\\rm Bondi}}$ and $L_{{\\rm bol}}$ we can use the formula $L_{{\\rm bol}}=\\eta${\\it \\.M}$_{{\\rm Bondi}}c^2$ to make an estimate of the radiative efficiency of accretion $\\eta$.  We find values on the order of $10^{-5}-10^{-6}$, except for the object 0915--11, which using the \\chandra\\ observation yields $\\eta=0.01$, and using the \\xmm\\ observation yields $\\eta=0.09$.  Those objects with $\\eta<<0.1$ are very likely to have inefficient accretion, but even 0915--11, based on the \\chandra\\ observation, is a candidate.  However, the \\xmm\\ observation makes the accretion scenario for 0915--11 ambiguous.  \\subsection{Comparison with previous results}  As mentioned above, Donato et al. (2004) have analyzed a larger sample of FRIs from the B2 and 3C catalogues.  Those authors found there to be no absorption of X--rays in FRIs, except in the LINER 1216+06 (NGC 4261), which is consistent with our results.  While Donato et al. found that the majority of the X--rays in classical FRIs are related to inefficient accretion onto the black hole, their analysis suggests that at least a fraction of the X--rays are beamed and thus potentially related to the base of the jet.  This lends further support to the idea that WLRGs/LINERs form a distinct group compared to classical FRIs.  A more general X--ray analysis of radio--loud AGN was made by Sambruna, Eracleous, \\& Mushotzky (1999) using \\asca.  In that paper the authors found there to be correlations for the [O$_{~\\rm{III}}$]$\\lambda$5007, the 5 GHz lobe, and the 12$\\rm \\mu m$ luminosities to the 2--10 keV intrinsic X--ray luminosity.  Those correlations demonstrated a slight lobe radio excess in WLRGs.  We have updated the plots of Sambruna et al. to include those WLRGs that were not in that paper, and have introduced a plot of the 5 GHz core luminosity and the 2--10 keV intrinsic X--ray luminosity (see Table~\\ref{tbl-6}, and Fig.~\\ref{correlations}).  These updated plots show that the WLRGs do indeed demonstrate a lobe radio excess; moreover, they form a distinct class as compared to the other subclasses of radio--loud AGN.  To quantify the degree of linear correlation between the 2--10 keV intrinsic luminosity and the 5 GHz lobe radio power for the WLRGs (and the other correlations that follow), we calculated the linear correlation coefficient $r$ and computed the chance probability $P_c(r;N)$ that a random sample of $N$ uncorrelated pairs of measurements would yield a linear correlation coefficient equal to or greater than $|r|$.  We also calculated the Spearman and Kendall parameters, which confirm the $P_c(r;N)$ results. We found the probability of correlation between $L_{\\rm lobe}$ and $L_{\\rm 2-10~keV}$ for the WLRGs to be $\\approx98\\%$, with a least squares fit yielding the relationship $\\log L_{\\rm lobe}=0.46~\\log L_{\\rm 2-10~keV}+22.6$ ($0.46\\pm0.1$, $22.6\\pm5$), represented by the dotted line in Fig.~\\ref{correlations}.  The remaining subclasses of AGN formed a distinct group of their own with a probability of correlation $>99.9\\%$; however, the least squares fit of this group yielded a significantly different relationship: $\\log L_{\\rm lobe}=0.84~\\log L_{\\rm 2-10~keV}+5.26$ ($0.84\\pm0.1$, $5.26\\pm6$), represented by the dashed line in Fig.~\\ref{correlations}.  This lobe excess in the WLRGs raises the possibility that there once was much more intense jet activity in these objects than is shown today.  The plot of the core radio power vs. the intrinsic 2--10 keV luminosity shows that the WLRGs follow the correlation of the other subclasses of radio--loud AGN.  The probability of correlation is $>99.9\\%$, and a least squares fit yields the relationship $\\log L_{\\rm core}=0.51~\\log L_{\\rm 2-10~keV}+18.56$ ($0.51\\pm0.08$, $18.56\\pm3.3$).  The plot of the MIR luminosity vs. the intrinsic 2--10 keV luminosity suggests that there may also be an MIR excess.  However, the MIR data for WLRGs are largely upper limits, so this suggestion should be taken with caution; also, because there were so many upper limits, we were not able to better constrain the correlation found by Sambruna et al. (1999), thus Fig.~\\ref{correlations} plots that previously found relationship: $L_{\\rm 12\\mu m}=0.83~\\log L_{\\rm 2-10~keV}+8.28$. Finally, we found the probability of correlation between $L_{[{\\rm O_{III}}]}$ and $L_{\\rm 2-10~keV}$ to be $>99\\%$, with a least squares fit yielding the relationship $\\log L_{[{\\rm O_{III}}]}=1.2~\\log L_{\\rm 2-10~keV}-10.33$ ($1.2\\pm0.2$, $-10.33\\pm7$), which is consistent within the errors with Sambruna et al. (1999). We notice that the WLRGs tend be have a slight [O$_{~\\rm{III}}$] deficiency; however, due to the small number of data points, we could not get good statistics for a least squares fit of the WLRGs as a distinct class."
    },
    "0410/astro-ph0410448_arXiv.txt": {
        "abstract": "We have compiled an extensive dataset on potential parts of the Monoceros tidal stream and performed an exhaustive survey of dwarf galaxy semi-analytic orbits in order to constrain its orbital properties. The best-fit orbits are subsequently realized as self-consistent $N$-body simulations in order to reproduce the spatial and velocity distribution of satellite debris. We find that all kinematic and geometric constraints can be fit by a single stream allowing for multiple wraps. The orbital eccentricity and inclination of the progenitor are strongly constrained to be $e = 0.10 \\pm 0.05$ and $i = 25^\\circ \\pm 5^\\circ$. Ten new estimates of proper motions from the Sloan Digital Sky Survey (SDSS) clearly exclude all retrograde orbits. Particles lost by the satellite populate two, nearly concentric rings naturally explaining the detection of stream stars at both 6--8\\,kpc (Ibata et al. 2003, Newberg et al. 2003) and 12--18 kpc (the Tri/And stream; Rocha-Pinto et al. 2003) from the Sun. We have attempted to predict the present location of the Monoceros stream progenitor using different information: (i) the kinematical and spatial distribution of detections and (ii) the different mean metallicity in the innner and the outer rings. Due to the lack of observational data in the whole range of Galactic latitudes, the geometrical/kinematical constraints leads to a wide range of possible locations. By associating older parts of the model stream with lower metallicity parts of the observed data, we argue in favor of a current location of $l \\sim 245, b \\sim -18$., with a distance to the Sun $r_s\\simeq $15 kpc. The mass of the progenitor has been poorly constrained due to the slow orbital decay. Similar fits have been obtained for masses $(3-9)\\times 10^{8} M_\\odot$. \\\\ We have analyzed to possible common origin of the Canis Major dwarf and the Monoceros stream. Canis Major dwarf moves on a prograde, nearly circular orbit ($e\\simeq 0.16$) in the Milky Way disk ($i\\simeq 4^{+14}_{-4}$ deg.). This orbital inclination is too low to account for the large vertical dispersion of stream stars. However, the bimodal distribution of radial velocities in the central region found by Martin et al. (2004b) probably indicates that their selection criteria to indentifying dwarf stars lead to a contamination of background stars. In that case, the kinematical data outlined above might result to an underestimated orbital inclination. Lastly, the distance estimation to Canis Major dwarf is around a factor two smaller than that obtained from our model. Unfortunately, the possible identification of the Monoceros stream progenitor in Canis Major remains unclear. ",
        "introduction": "In a $\\Lambda$CDM universe, the inner regions of massive galaxies like the Milky Way gain a large fraction of their mass through tidal disruption and accretion of a large number of low-mass fragments (e.g van den Bosch, Tormen \\& Giocoli 2004). The fossil records of these merging processes may be observable nowadays in form of long tidal streams or large-scale stellar substructures around the parent spiral galaxies. Numerical galaxy formation simulations show that dynamical friction brings sub-structures from outer halo regions to the neighbourhood of the parent galaxies' disks. If haloes are flattened (oblate) and their axi-symmetry plane is that of the disk, orbits of non-polar satellite galaxies tend to become co-planar and circularise with time (Pe\\~narrubia, Kroupa \\& Just 2004). In this scenario, tidal debris of several disrupted satellite galaxies might have contributed to the formation of the stellar disk if they followed nearly circular orbits with a low orbital inclination at late times of their evolution (Navarro 2004 and references therein).  The Milky Way is an important laboratory in which the predictions of this cosmological scenario can be tested. In the last decade, large-scale surveys have proved the existence of tidal streams  (Sgr: Ibata et al. 1994; Monoceros: Newberg et al. 2002) in our Galaxy, providing strong observational evidence that disruption of dwarf satellites contributes to the assembly of some components of our Galaxy. These tidal streams offer an unique opportunity to study accretion events in considerable detail using the chemical, kinematic and spatial distribution of tidal stream's stars, which can be directly compared against predictions of N-body simulations of merging events (e.g Law, Johnston \\& Majewski 2004).   Recently, the Sloan Digitized Sky Survey team reported the discovery of a coherent ring-like structure  at low galactic latitude spanning about 100 degrees in the sky (Newberg et al. 2002; Yanny et al. 2003). Follow-up  observations (Ibata et al. 2003) found that this structure of low metallicity stars surrounds the Galactic disk at Galactocentric distances from $\\sim$ 15 kpc to $\\sim 20$ kpc. Different groups proposed two different scenarios to explain the nature of this feature in the outer Galactic disk (see Helmi et al. 2003): i) a new tidal stream (Yanny et al. 2003); and ii) a stellar ring resulting from perturbations in the disk similar to the remnants of ancient warps (Ibata et al. 2003). Tracing this structure with 2MASS M giant stars,  Rocha-Pinto et al. (2003) concluded that its structural characteristics are consistent with the interpretation of this structure as the fossil of a merging dwarf galaxy, similar to the Sgr dwarf galaxy (Ibata et al. 1994), in the Galactic plane.  Interestingly, Frinchaboy et al. (2004) and Crane et al. (2003) also suggested some nearby Galactic open and globular clusters with coordinated heliocentric radial velocities, indicating a possible common origin with the tidal stream.    Unlike the Sagittarius dwarf galaxy, the Monoceros stream has been detected prior to locating the main body of the parent galaxy. The location is still controversial. From the analysis of the 2MASS catalogue, Martin et al. (2004a) reported a strong elliptical-shaped stellar over-density in the constellation of  Canis Major, which is coincident in position and radial velocity with a small group of 4 globular clusters (see Sec.~\\ref{sec:cluster}). Using a fairly simple model, they identify the CMa over-density as the main body of the progenitor dwarf galaxy of the Monoceros stream (named Canis Major dwarf). Bellazzini et al. (2004) presented colour-magnitude diagrams in the surroundings of the CMa location, concluding that the system is situated at 8$\\pm$ 1 kpc from the Sun and it is composed by a metal-rich, intermediate-age population. Momany et al. (2004), however, comparing proper motions and radial velocities of Canis Major M-giant stars, obtain that the overdensity in this region mimics the thick disk kinematics. Moreover, they find that the star counts in that region are those expected in the standard Milky Way model if taking into account the warp and flare of the disk. However, a deep colour-magnitude diagram of the center of the Canis Majoris over-density by Mart\\'\\i nez-Delgado et al. (2004) shows a very well-defined main-sequence consistent with a limited extent in distance, confirming that this stellar population is associated with a distinct, possibly still bound stellar system with properties (surface brightness, absolute magnitude, stellar content) compatible with those of Local group dwarf spheroidal galaxies. In base of these last results, the Canis Major overdensity will be referred to the CMa dwarf galaxy in this work.   In this paper we present the results of an extensive search of possible disruption scenarios of satellite galaxies to explore: a) how many of the 'over-density signatures' can be attributed to a single stream and b) to constrain its progenitor orbit. With an objective criterion we select those orbits that reproduce the spatial and kinematical distribution of candidate debris. Beyond exploring the possible progenitor location, the results will: (i) help to carry out new surveys in different regions of the sky and (ii) constrain the thick disk formation history through the distribution evolution of stripped stars.  This contribution has been organised as follows: In Section~\\ref{sec:obs} we compile the observational data available so far. Section~\\ref{sec:tools} outlines the galaxy and satellite models that we use in our investigation. We also describe our semi-analytic orbit algorithm and the results of the orbital parameter survey. We obtain the distribution of debris via the N-body algorithm sketched in Section~\\ref{sec:nbody}. In Section~\\ref{sec:res} we compare the best-fit orbit with the observational data, whereas in Section~\\ref{sec:dis} we comment different implications of our results. ",
        "conclusions": "\\label{sec:conc} In this paper, we have used a combined semi-analytic/N-body technique to explore the nature and origin of the Monoceros tidal stream. This method has allowed us to explore a large parameter space in a systematic way, with the goal of constraining the orbital properties of the Monoceros stream. \\\\ We have found that the available observational data at the present day is sufficient to robustly constrain some aspects of the Monoceros stream history and progenitor. In particular:  (i) the heliocentric distance range at which the stream is observed plus the heliocentric radial velocity curve determines an  orbital eccentricity of $e\\simeq 0.10\\pm 0.05$, (ii) the range of galactic latitudes indicates an orbital inclination of $i\\simeq 25^\\circ\\pm 5^\\circ$ from the disk plane and, lastly, (iii) proper motions of stream stars are only compatible with a prograde orbit.  We found that it is considerably less straightforward to predict the main body position through numerical calculations. The range of distances, the radial velocity curve and the vertical extension of debris are relatively insensitive to the final location of the satellite galaxy, since they define a given volume in phase-space that the progenitor fills up after several wraps. Yet, the projected positions of stream tails are time-dependent reflecting, therefore, the position of the main system or, equivalently, the final integration time of our N-body model.\\\\ Besides the uncertainty in the satellite remnant position, the lack of detections at low galactic latitudes (owing to disk absorption within $b\\in [-20^\\circ,20^\\circ$]) as well as for a complete range of Galactic longitudes makes difficult to constrain the location of the main system. Owing to these limitations, our model matches the geometrical and kinematical distribution of debris if the main system remnants locate within the range $100^\\circ\\geq l \\geq 200^\\circ$, corresponding to integration times from 2.66 to 3.04 Gyr. \\\\ On the other hand, observations show that close detections of the stream have higher metalicities than the distant ones, which provides an additional constraint to our model.  We have used this constraint by assuming that the progenitor's stellar metalicities are related to their initial binding energy. In particular, we assume that stars with low binding energies move, in average, in the outer satellite regions and, therefore, should present lower metalicities than those with high binding energies. By comparing the initial binding energy (metallicity) distribution with observations for the range of locations commented above, we find that the metallicity gradient between the distant and the close detections can only be reproduced for integration times longer than 2.9 Gyr. In that case, the main system location ($l\\sim 245^\\circ$, $b\\sim -18^\\circ$) is similar to that of the Canis Major dwarf, with a heliocentric distance of $15.2\\pm 0.8$ kpc. The treatment employed to determine the distribution of metalicities in the stream is, however, approximate, because have no way of knowing the initial metallicity gradient in the stream progenitor. Moreover, we implicitly assume that detections with different metalicities reflect an initial property of the progenitor. Other scenarios are possible and, for example , we cannot discard that streams with different metalicities at different distances belong to different progenitors.   The best-fitting model predicts a halo axis-ratio of $q_h=0.6$. However, the poorly-constrained selection function of observations and the incomplete area coverage leads to degenerated solutions for different values of $q_h$. In particular, we find that similar fits can be obtained with halo axis-ratios $q_h\\in[0.6,0.8]$.   The model we present here is, therefore, far from being definitive. Future detections, mapping larger areas of the sky, will provide more constrains on the progenitor's orbit and a better determination of parameters listed in Table~\\ref{tab:param}. New detections are also necessary to constrain the progenitor location by means of numerical models.  With the recently available kinematical data of the Canis Major dwarf, which Martin et al. (2004a) claim to be the progenitor of the Monoceros stream, we have integrated the orbit backwards in time, comparing its orbital properties against those of our model. We find contradictory points in favour and against that suggestion:  (i) In favour: \\begin{itemize} \\item The orbital eccentricity of CMa ($e \\simeq 0.16$) and orbital sense of motion (prograde) are consistent with those of our model. \\item The CMa's projected location ($l=240^\\circ$, $b=-8^\\circ$) is consistent with the geometrical and metallicity constrains imposed by the Monoceros stream stars. \\item The orbital inclination ($i\\simeq 4^{+14}_{-4}$ deg.) is consistent with that of our model within $1.5\\sigma$ errors from proper motions. \\item The radial velocity of CMa ($v_r=109$ km/s) is also similar to that of our model ($v_r=135$ km/s) and presents a low dispersion incompatible with thick disk features. \\end{itemize}   (ii) Against: \\begin{itemize} \\item The orbits of globular clusters NGC 1851, 1904 and 2298, which Bellazzini, Ferraro \\& Ibata (2003) and Martin et al. (2004a) claim to be an evidence that the CMa over-density region is the remnant of the Monoceros stream progenitor, are inconsistent with the orbit of CMa, since NGC 1851 and NGC 1904  move on highly eccentric orbits and NGC 2298 on a retrograde, eccentric one.  \\item The bimodality in the velocity distribution of CMa's M-giant stars (Martin et al. 2004b) can be reproduced by our model in regions where the projected positions of two stream tails overlap. That bimodality in $v_r$ cannot be observed in the central part of our satellite model since the surface density is several orders of magnitude larger than that of overlapping stream tails. This fact seems to point out to a possible contamination of background giant stars in the sample used to measure proper motions and the radial velocity of the main system. As a result, the orbital inclination might have been underestimated. \\item The main body of our best-fit model's progenitor is $\\sim 15$ kpc from the Sun --- twice as distant as the observed Canis Majoris stellar structure. Since the mass of our model is poorly constrained, larger mass values cannot be rejected. In particular, dynamical friction would drive our satellite model down to 7.2 kpc in 3 Gyr for an initial mass of approximately $1.8 \\times 10^9 M_\\odot$, whereas our best-fitting model has an approximate initial mass of $(6\\pm 3)\\times 10^8 M_\\odot$. It remains unclear, however, whether the resulting distribution of debris from such a massive satellite undergoing tidal disruption would also reproduce the observations. \\end{itemize}  Unfortunately, with the results obtained in this work we cannot unambigously determine whether the orbit of the CMa dwarf is consistent with that of the Monoceros stream progenitor, partly due to the uncertainty in the available observational data. More strict criteria to select stars that belong to the CMa dwarf appear necessary to obtain reliable measurements of the kinematical properties of this system.    \\vskip1cm  We want to thank Dana I. Dinescu for helping us to analyse the information obtained from proper motions. We also thank H. Rocha-Pinto and J. D. Crane for giving us access to their data before being published.  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 (\\texttt{http://www.sdss.org/}).  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0410/astro-ph0410212_arXiv.txt": {
        "abstract": " ",
        "introduction": "Astronomy research images are now almost exclusively digital.  There is an increasing need to store these images and to send them over the Internet.  Average bandwidth and storage capacity, although increasing, continue to be a bottleneck that frequently limits scientific exploitation of these images.  Lossless compression of astronomical image files therefore creates a clear advantage over non-compression since it increases effective storage and bandwidth and so bolsters scientific utility.  Lossy compression is more controversial, however, as the scientific value of the data lost in the compressed images must be weighed against the scientific value of the data gained by the extra bandwidth and storage space.  One of the premier scientific uses for astronomical images is photometry.  Whether detecting the presence of, for example, distant supernovae, Local Group microlensing, binary star variability, or planetary transits, the need for photometric accuracy in astronomy images remains a primary objective for many astronomical research projects.  Frequently, the photometric value of an astronomical image is concentrated in the point sources in the image.  Conversely, the bulk of the image size is concentrated in the background behind these sources.  When the number of pixels taken up by sources is small compared to the number of pixels that compose the background, it becomes possible to significantly compress the image size while preserving a certain level of photometric integrity.  While lossless compression algorithms preserve 100\\% of the signal, lossy algorithms can provide a better compression factor while losing some of the signal \\citep{14,6,7}. However, lossy compression algorithms tend to convolve science with art.  Astronomy-specific lossy data compression algorithms are not new, and many have been proposed and widely used \\citep{3,4,5,13,14}. HCOMPRESS \\citep{13} is a commonly used FITS compression program based on a two-dimentional Haar wavelet transform. This compression program is fast and provides a relatively high compression factor. However, while providing the user control of the lossiness$\\backslash$compression tradeoff, HCOMPRESS provides a limited control over the type of the signal that is lost in the compression $\\backslash$ decompression process.  The most similar compression approach to that discussed here is the insightful FITS compression program FITSPRESS \\citep{14}. FITSPRESS is a wavelet based compression algorithm that has, among other things, sensitivity to preserving the brightest image pixels.  Lossy compression (e.g. JPG, HCOMPRESS) tends to be controlled by parameters that have no scientific or statistical meaning.  In contrast, the PHOTZIP FITS compression algorithm provides the user control of the preserved$\\backslash$lost signal in scientific terms. Like \\citep{18}, PHOTZIP should be used as a preprocessor to multi-purpose lossless compression algorithms (e.g. LZW) and allows those algorithms a higher compression factor by losing some of the signal, but ensures that only {\\it background} information of the image will lose some of its signal. The algorithm provides an interface which can be used in order to define, in terms of $\\sigma$, the criteria for a pixel to be considered as background, and guarantees a user-defined maximum absolute difference (also in terms of $\\sigma$) for any background pixel that loses signal.  In Section 2 we describe the algorithm, in Section 3 we present ways to improve the compression factor, in Section 4 we discuss photometric integrity and in Section 5 we discuss the performance of the algorithm. ",
        "conclusions": "When even a single pixel is quantized, photometry can be at least partially compromised.  A background estimated from the mean of surrounding pixels, however, will usually end up closer to the original background. We have shown that using even as little as 1600 pixels for background estimation can reduce the background error in the mean of the quantized pixels to below a single count.  Practically, this means that the background estimation and hence photometric measurements can be protected in the PHOTZIP compression$\\backslash$decompression process.  In sum, we present here a simple lossy compression algorithm for astronomical images. The main advantage of this algorithm is that it can preserve the signal of bright pixels, while symmetrically losing signal only from the background.  The criteria for preserving a pixel's value is user-defined (in terms of $\\sigma$), so the user can accurately control the compression factor/signal loss trade-off.  In addition, the algorithm guarantees a user-defined maximum absolute difference (also in terms of $\\sigma$), so the user can control the amount of lost signal for those areas in the frame that are not preserved.  The algorithm was implemented and tested on unsigned 16-bit integer FITS images, but we believe that the same approach can be applied also to signed integer and floating point images. The algorithm was tested on some typical astronomical 16-bit integer FITS images and appeared to be effective.  PHOTZIP is now in routine use on FITS data taken by the Night Sky Live project.  We thank an anonymous referee for numerous and insightful comments.  LS acknowledges fellowship support from Michigan Technological University.  RJN acknowledges support from the National Science Foundation."
    },
    "0410/astro-ph0410738_arXiv.txt": {
        "abstract": "The changing universe model (CUM) describes galaxy parameter relationships (SESAPS '03, session EB 2).  The CUM must be successfully applied to cosmological scale observations to be considered a cosmological model.  A major component of current cosmological models is the Hubble constant $H_\\mathrm{o}$.  An equation is derived using the CUM model relating redshift $z$ and the distance $D$ to galaxies and is applied to a sample of 32 spiral galaxies with $D$ calculated using Cepheid variable stars.  The equation predicts a galaxy may have $z<0$ in special circumstances.  Three elliptical galaxies with peculiar characteristics are discovered to be CUM Sinks.  The Sinks give a physical explanation of the ``Virgocentric infall'' and ``Great Attractor'' observations without a large, unobserved mass.  At low cosmological distances, the equation reduces to $z \\approx \\exp(KD) \\, -1 \\approx KD$, where $K$ is a constant, positive value.  The equation predicts $z$ from galaxies over 23 Gpc distant approaches a constant value on the order of 1000.  The CUM gives a physical basis for the Doppler shift of particle photons. ",
        "introduction": "The changing universe model (CUM) \\cite{hodg4} was derived from examining galaxy scale observations.  The CUM must be successfully applied to cosmological scale observations to be considered a cosmological model.  That the redshift $z$ of emitted photons from galaxies generally increases with the extragalactic distance $D$ (Mpc) between the emitting galaxy and observer is well known ($z - D$ relationship).  The $z - D$ relationship has different meanings in current cosmologies.  The idea of ``expanding space'' may be defined as a set of comoving coordinates wherein Space is like an expanding rubber sheet that carries matter along with it.  Space is expanding between particles.  Alternatively, expanding Space may be like an expanding gas in a vacuum.  Space exerts a force on particles causing matter to be carried along with the expanding Space.  An alternate idea is that Space is an unobtrusive medium through which matter moves in free particle motion with a certain initial velocity and inertia.  In these cases, photons have a dual particle and wave nature and the distribution of matter in the universe is assumed isotropic and homogeneous in cosmological scale volumes.  Currently fashionable, cosmological models attributes $z$ to be solely a Doppler shift of a light wave.  The ``standard model'' assumes the universe is homogeneous and isotropic.  The assumption of homogeneity is compatible with either a static galaxy distribution or with a very special velocity field obeying the Hubble Law \\cite[page~396]{binn} \\begin{equation} D =  \\frac{c}{H_\\mathrm{o}}z \\label{eq:1}, \\end{equation} where $H_\\mathrm{o}$ (km~s$^{-1}$~Mpc$^{-1}$) is the Hubble constant and $c$ (km~s$^{-1}$) is the speed of light.  The $H_\\mathrm{o}$ occupies a pivotal role in current cosmologies.  The methods of calculating supernova distances, the cosmological microwave background (CMB) power spectrum, weak gravitational lensing, cluster counts, baryon oscillation, expansion of the universe, and the fundamental aspects of the Big Bang depend on $H_\\mathrm{o}$ and the Hubble law.  However, the determination of $H_\\mathrm{o}$ has a large uncertainty and different researchers calculate different values.  Figure~\\ref{fig:1} shows the measured galactocentric redshift $z_\\mathrm{m}$ versus the calculated redshift $z_\\mathrm{H}$ using Eq.~(\\ref{eq:1}), $H_\\mathrm{o}=70$ km~s$^{-1}$~Mpc$^{-1}$, and $D$ calculated using Cepheid variable stars for 32 galaxies \\cite{free,macr}.  The correlation coefficient of $z_\\mathrm{H}$ versus $z_\\mathrm{m}$ is 0.80.  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{f1.eps} \\caption{\\label{fig:1} Plot of the measured redshift $z_\\mathrm{m}$ versus the calculated redshift $z_\\mathrm{H}$ after the third iteration using Eq.~(\\ref{eq:1}) and $D$ calculated using Cepheid variable stars for 32 galaxies \\cite{free,macr}.  The straight line is a plot of $z_\\mathrm{H} = z_\\mathrm{m}$.  The circles indicate the data points for galaxies with (l,b) = (290$^\\circ \\pm 20^\\circ$,75$^\\circ \\pm 15^\\circ$).} \\end{figure}  The deviation of the recession velocity of a galaxy from the straight line of the Hubble Law is ascribed to a ``peculiar velocity'' of the photon emitting galaxy relative to earth \\cite[page 439]{binn}.  The deviation from the Hubble Law is the only means of determining the peculiar velocity of a galaxy.  Therefore, to determine the distances to $n$ galaxies requires $n+1$ constants.  However, the peculiar velocity is assumed to be limited to less than 1000 km~s$^{-1}$, which is a small error at cosmological scale distances.  Also, the average peculiar velocity for all galaxies is assumed to be zero on the scale that the universe appears homogenous.  The 2dFGRS \\cite{peac} suggests this scale is $z > 0.2$.  The circles in Fig.~\\ref{fig:1} denote data for galaxies in the general direction of $( l, b)=(290^\\circ \\pm 20^\\circ, 75^\\circ \\pm 15^\\circ)$.  These galaxies are outliers.  \\citet{aaro2} found the peculiar velocity field in the local supercluster is directed toward NGC 4486 (Messier 087) with $( l, b) \\approx (284^\\circ, 74^\\circ)$ at a speed of 331$\\pm$41 km~s$^{-1}$.  This has been called the ``Virgocentric infall''.  NGC 4486 is a peculiar, large, elliptical galaxy with strong X-ray emissions.  In addition, \\citet{lilj} detected a quadrupolar tidal velocity field from spiral galaxy data in addition to the Virgocentric infall pointing toward (l,b) = (308$^\\circ \\, \\pm 13^\\circ $, 13$^\\circ \\, \\pm 9^\\circ $) at a speed of $\\approx 200$ km~s$^{-1}$.  NGC 5128 (Centaurus A) at (l,b) = (310$^\\circ$, 19$^\\circ$) and at a distance of $ 3.84 \\pm 0.35$ Mpc \\cite{rejk} is a galaxy with properties \\cite{isra} similar to NGC 4486.  \\citet{lynd} found elliptical galaxies at distances in the 2000-7000 km~s$^{-1}$ range are streaming toward a ``Great Attractor'' centered on $( l, b) = (307^\\circ \\pm13^\\circ, 9^\\circ \\pm8^\\circ)$.  Centaurus B at (l,b) = (310$^\\circ$, 2$^\\circ$) is a galaxy with properties similar to NGC 4486.  In a more recent analysis, \\citet{huds} suggested a bulk flow of 225 km~s$^{-1}$ toward $( l, b) \\approx (300^\\circ, 10^\\circ)$.  However, the total mass in these directions appears to be insufficient to account for the peculiar velocity fields using Newtonian dynamics.  The changing universe model (CUM) proposed that Sources at the center of galaxies are continually erupting the Space and matter of our universe \\cite{hodg4}.  Space is a constituent of our universe that ``flows'' from point Sources at the center of galaxies to Sinks.  The ``flow'' is like the nonviscous flow of heat from point sources to point sinks in a solid.  Space and the distribution of matter in Space are neither isotropic nor homogeneous.  The Space potential $\\rho$ is like temperature in a solid.  The gradient $\\bm \\nabla \\rho$ of the Space potential influences matter by exerting a force on the surface area of matter.  Photons are columns of basic particles called hods.  The hod is a two-dimensional surface with zero thickness.  As photons move through Space, hods may be stripped (redshift) from or added (blueshift) to the photon hod column.  The $z - D$ relationship might show the reality of the physical nature of Space, of Space expansion, of the homogeneity of the universe, and of light.  This Paper examines the $z - D$ relationship of 32 spiral galaxies within 25 Mpc using the CUM of Space and of photons.  The developed equation results in a higher correlation coefficient than the Hubble law.  Also, a physical basis for the observed Virgocentric infall, for the Great Attractor phenomena, and for the Doppler shift of particle photons is given.  The object of this article is to apply the CUM beyond galaxy scale distances by developing an equation for the calculation of $z$ and by testing the equation using observations of galaxies that have $D$ calculated using Cepheid variable stars.  In section~\\ref{sec:model}, the CUM $z$ calculation equation is developed.  The equation is applied to galaxies, and compared with galaxy observations in Section~\\ref{sec:results}.  Some implications of the equation are discussed in~\\ref{sec:impl}.  The results are discussed in Section~\\ref{sec:disc}.  Section~\\ref{sec:conc} lists the conclusions. ",
        "conclusions": ""
    },
    "0410/astro-ph0410024_arXiv.txt": {
        "abstract": "\\noindent Understanding the galaxy in which we live is one of the great intellectual challenges facing modern science. With the advent of high quality observational data, the chemical evolution modeling of our galaxy has been the subject of numerous studies in the last years. However, all these  studies have one missing element which is \\textsl{the evolution of close binaries}. Reason: their evolution is very complex and single stars only perhaps can do  the job. (Un)Fortunately at present we know that the majority of the observed stars are members of a binary or multiple system and that certain objects can only be  formed through binary evolution. Therefore galactic studies that do not account for close binary evolution may be far from realistic. \\\\ \\noindent Because of the large expertise developed through the years in stellar evolution in general and binary evolution in particular at the Brussels Astrophysical Institute, we found  ourselves in a privileged position to be the first to do chemical evolutionary simulations with the inclusion of detailed binary evolution. The  complexity of close binary evolution has kept many astronomers from including binary stars into their studies. However, it is not always the easiest way of living that gives you the most excitement and satisfaction.\\\\  \\noindent The paper is a 115 page review, published in New Astronomy Reviews 48, 861. An offprint (with color figures) can be obtained via NARev or by sending a request to Dany Vanbeveren, dvbevere@vub.ac.be. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410206_arXiv.txt": {
        "abstract": "We modeled the spectral energy distribution (SED) of high redshift radio quiet quasars (high-$z$ RQQs). We computed spectra in a patchy corona geometry where an accretion disk extends to the last stable orbit and the Comptonizing active regions (hot clouds) are distributed above the disk. We explored the model parameter space to find theoretical parameters that give spectra with the optical/UV luminosity, the X-ray loudness, and the X-ray photon index compatible with those of high-$z$ RQQs observed with {\\it Chandra}. We found that a range of solutions is possible, from high-$kT_{\\rm e}$ low-$\\tau$ to low-$kT_{\\rm e}$ high-$\\tau$ models. The solutions require low level of energy dissipation in the hot clouds and low disk covering factor. The modeled mass is of the order of $10^{10}\\, M_{\\odot}$ and the accretion rate is $\\dot{M} \\ge 0.2\\, \\dot{M}_{\\rm Edd}$. We compare our results to those obtained previously for hot inner flow geometry. ",
        "introduction": "Modeling quasars' broad band energy spectra provides important information about the geometry and environment of the quasars at different epochs. Such spectra are thought to originate in an accretion of matter onto a massive black hole (Rees 1984). In the standard scenario a thin accretion disk that is formed thermalizes the gravitational potential energy and radiates it as a disk black body giving a characteristic bump in the spectrum (Shakura \\& Sunyaev 1973). In the case of high redshift quasars, this bump is shifted into the optical/UV range (Shields 1978), which makes it easy to observe and study. The X-ray part of the spectrum is believed to originate in an inverse Compton process of the disk photons on energetic electrons. Modeling the X-ray component gives a direct information about the geometry of the accretion flow (i.e.\\ the relative location of the cold accretion disk and Comptonizing plasma).  Radio quiet quasars are good objects for such studies since their X-ray spectra are not ``contaminated'' by any additional components connected to the radio emission. Moreover, the radio quiet quasars comprise majority among the quasars population (Stern et al.\\ 2000). They are more common, and hence more typical. Recent X-ray observations of quasars located at redshift $z>4$ (e.g. Becker et al.\\ 2001; Brandt et al.\\ 2002a; Bechtold et al.\\ 2003, hereafter B03) provide for the first time a unique opportunity to study these objects at the early Universe and investigate their evolution.  Thus, we focus our study on the high redshift radio quiet quasars (high-$z$ RQQs).  The observations of accreting sources support the hypothesis that the accretion  flow is composed of a cold medium producing a thermal radiation and a hot medium radiating in X-rays. The puzzling question is that of a relative location of these two phases. A variety of geometries has been considered in order to explain the observed spectra and correlations. In Sobolewska, Siemiginowska \\& \\.{Z}ycki 2004 (hereafter Paper I) we shortly reviewed currently considered scenarios for the accretion flow which were successful in explaining spectra of galactic X-ray binaries (XRB) and Seyfert galaxies. We discussed possible tests of geometry that can be conducted based on modeling the broad band spectra. In particular, a correlation between the amplitude of reflection, $R$, and the X-ray photon index, $\\Gamma$, observed in X-ray binaries and Seyfert galaxies (Zdziarski, Lubi\\'nski, \\& Smith 1999) seems to be one of the criteria. Such correlations can be reproduced by the following models: (1) a truncated disk surrounding a hot inner flow (Zdziarski et al.\\ 1999), (2) a hot corona formed by active regions  above the disk, outflowing at relativistic speeds (Beloborodov 1999a, hereafter B99), (3) a cold disk covered by a hot ionized skin which decreases the effectiveness of thermalization (Nayakshin, Kazanas \\& Kallman 2000) (4) a hot spherical inner flow surrounded by cold clouds (Malzac 2001). In Paper I we investigated in details the first possibility. We assumed that the accretion flow in high-$z$ RQQs proceeds through a standard optically thick cold accretion disk covered with a hot corona. At a certain radius the disk was assumed to evaporate and a hot inner flow was formed. We computed spectra in such geometry and explored the model parameter space. We found that modeling both the optical/UV and X-ray component provides tight constraints on the Comptonizing plasma parameters through the observed X-ray loudness parameter, $\\alpha_{\\rm ox}$, X-ray photon index, $\\Gamma$, and the optical/UV luminosity, $l_{\\rm UV} \\equiv \\log(\\nu{L_{\\nu}})$ at $\\nu = c/{\\lambda}_{\\rm 2500\\, \\AA}$ in the rest frame. For most objects the truncation radius is required to be larger than the radius of the innermost stable orbit, $r_{\\rm tr}>3\\,R_{\\rm S}$, with an upper limit of $r_{\\rm tr} \\leq 40\\, R_{\\rm S}$ (where $R_{\\rm S} = \\frac{2GM}{c^2}$ is a Schwarzschild radius). We also concluded that the advection in the inner hot flow may play an important role. Moreover, the shape of the inner hot flow may deviate from spherical one. Instead, it may be considerably flattened. Our modeling suggested high accretion rates, $\\dot{M} \\ge 20\\% \\dot{M}_{\\rm  Edd}$ which may suggest similarity of high-$z$ RQQs to XRBs in the high state (HS) or very high state (VHS) of black hole X-ray binaries. However, in these spectral states the accretion disk is thought to extend to the last stable orbit. This provides motivation for investigating the geometry of a disk extending to the innermost stable orbit, with some kind of an X-ray active corona.  In this paper we investigate such an active corona geometry. The structure of the paper is as follows. In Section~\\ref{sec:model} we describe the model. In Section 3. we present the results. Section 4. contains application of the model to the sample of high-$z$ quasars from B03. Finally, in Section 5. we discuss the results and give the concluding remarks. ",
        "conclusions": "\\label{sec:concl}  In this paper we applied the patchy corona model to high-$z$ RQQs, and explored the model parameter space. We compared the results to the results of the truncated disk with the hot inner flow model studied in Paper I. Our results show that:  \\noindent 1. If the hard X-ray part of the observed high-$z$ RQQs spectra originates in the Comptonization process in the corona formed by hot clouds above an accretion disk, the clouds must be vertically elongated and intercept less than $\\sim 10$\\% of the accretion disk radiation.  \\noindent 2. Similarly as in the hot inner flow model, the observed X-ray loudness and the photon index put strong constraints on the modeled optical depth. However, the number of observables is too small to uniquely determine the Comptonizing plasma parameters, and the grid of solutions with different pair ($\\tau$, $kT_{\\rm e}$) is possible for each object.  \\noindent 3. Both models require a similar mass of the black hole (of the order of $10^{10}\\, M_{\\odot}$), and accretion rates ($\\dot{M} > 0.2\\, \\dot{M}_{\\rm Edd}$).  \\noindent 4. In the patchy corona model a higher fraction of energy is allowed to be dissipated in the hot plasma above the disk than in the hot inner flow model.  \\noindent 5. The objects in the sample with X-ray spectra characterized by the photon index of $\\Gamma \\le 1.5$ can be modeled only assuming the outflowing corona.  \\noindent 6. Future high quality observational data, both in the optical/UV and X-rays can significantly improve constraints on the model parameters and help to discriminate between different accretion flow geometries."
    },
    "0410/astro-ph0410176_arXiv.txt": {
        "abstract": "We have conducted a deep, very long baseline interferometry (VLBI) observation at 1.4~GHz of an area of sky located within the NOAO Bo\\\"otes field, using the NRAO Very Long Baseline Array and 100-m Green Bank Telescope.  Applying wide-field VLBI techniques, a total of 61 sources, selected from a Westerbork Synthesis Radio Telescope (WSRT) image, were surveyed simultaneously with a range of different sensitivities and resolutions. The survey covered a total of 1017~arcmin$^2$ $=$ 0.28~deg$^2$ divided into annular fields centered on $\\alpha=$ 14h 29m 27.0s and $\\delta=$ $+$35\\arcdeg 28\\arcmin 30.00\\arcsec\\, (J2000). The inner 0-2\\arcmin\\, of the field reached an unprecedented 1 $\\sigma$ rms noise level of $9~\\mu$Jy~beam$^{-1}$ and yielded two detections. VLBI J142923.6466 and VLBI J142934.7033 have brightness temperatures in excess $10^{6}$~K and locate the active nucleus of their host galaxies (NDWFS J142923.6$+$352851 and NDWFS J142934.7$+$352859 with $I \\sim 16.3^{m}$ and $I \\sim 19.6^{m}$ respectively). Further deep surveys of the inner 2-4\\arcmin\\, and 4-6\\arcmin\\, of the field, with 1 $\\sigma$ rms noise levels of 11-19~$\\mu$Jy~beam$^{-1}$, detected a previously known source, VLBI J142910.2224, a quasar with a brightness temperature in excess of $10^{9}$~K that was also used during these observations as an in-beam phase calibrator. The deep VLBI survey between 0\\arcmin\\, and 6\\arcmin\\, thus detected 3 radio sources, drawn from a total of 24 targets. A shallower VLBI survey, conducted between 6\\arcmin\\, and 18\\arcmin\\, of the field center, and with 1 $\\sigma$ rms noise levels of 37-55~$\\mu$Jy~beam$^{-1}$, detected a further 6 radio sources, drawn from 37 additional targets. Each of those 6 VLBI detections has a brightness temperature in excess of $10^5$~K; this hints that those 6 are accretion-powered, a suggestion reinforced by the double structure of 3 of them. Combining both the deep and shallow VLBI surveys, optical identifications are available for 8 of the 9 VLBI detections.  Only VLBI J142906.6095 remains unidentified ($I > 25.6^{m}$), quite unusual as its integrated WSRT flux density is 20~mJy. Two other sources are not detected in K-band ($K > 18.5^{m}$) suggesting that some significant fraction of these compact radio sources may be located at $z > 1$. The VLBI detection rate for sub-mJy WSRT radio sources is 8$^{+4}_{-5}$\\%.  The VLBI detection rate for mJy WSRT sources is higher, 29$^{+11}_{-12}$\\%. This observational trend is expected from a rapidly evolving radio source population.  Moreover, this trend supports deep radio surveys, at lower resolution, that infer the radio emission associated with fainter sub-mJy and microJy sources arises via processes associated with extended regions of star formation. The 9 VLBI detections reported here pin-point the precise location of active nuclei or candidate active nuclei, and their VLBI positions can help to anchor the NOAO Bo\\\"otes field to the International Celestial Reference Frame.  The simultaneous detection of several sub-mJy and mJy radio sources, in a single observation, suggest that their combined response may be used to self-calibrate wide-field VLBI data.  There is every prospect that future deep VLBI observations can take advantage of this wide-field technique, in any random direction on the sky, thereby generating large-area, unbiased surveys of the faint radio source population. ",
        "introduction": "Evolutionary models of extragalactic radio sources \\citep[e.g.,][]{con84} indicate that the overwhelming majority of sources with integrated flux densities in excess of $S_{\\rm I} \\sim 1$~mJy are typically identified with active galaxies, energised by accretion onto massive black holes.  These bright sources have a median angular size of 5-10\\arcsec\\, \\citep[e.g.,][]{bon03,con03} and are usually resolved by the VLA on arcsecond and sub-arcsecond scales. Often a significant fraction of the total radio emission resides within a compact and unresolved radio ``core'', with a size measured on milliarcsecond (mas) and sub-mas scales using the techniques of very long baseline interferometry (VLBI).  Fainter sub-mJy and microJy radio sources have a higher probability of being energised by, and coextensive with, star formation in galaxies located at moderate redshifts $z \\sim 0.3-1$ \\citep{ric00,sma02}. \\citet{gar02} has shown that the bulk of these distant systems obey the locally derived \\citep[e.g.,][]{con92} FIR-radio correlation.  Indeed, only about 20\\% of the sub-mJy radio source population appears to be optically identified with active galaxies; a further 10\\% are optically faint sources with $I > 25^{m}$ and are difficult to classify from their weak radio emission alone \\citep{mux04}. MERLIN-VLA observations of faint radio sources in the HDF show that the microJy radio sources are resolved and typically sub-galactic in their overall extent, with typical measured sizes of 0.2-3 arcseconds \\citep{mux04}. This is consistent with the analysis of \\citet{con92}, limiting the maximum brightness temperature of star forming galaxies to $< 10^{5}$~K at frequencies above 1 GHz.  A recent study conducted by \\citet{hai04} suggests that a significant fraction of the faint radio source population may be located at high redshift. At the level of 10~$\\mu$Jy they predict, in their simplest models, source surface densities of 10~deg$^{-2}$ at $z > 10$. However, at these faint sensitivity levels it is increasingly difficult to identify radio-loud active galaxies from the dominant star forming radio source population, especially in the case of optically faint systems.  High-resolution VLBI observations resolve out these extended star forming galaxies and, currently, any compact radio sources in these galaxies, such as luminous Type IIn SNe \\citep[e.g.,][]{wil02,smi98}, would be too faint to detect at cosmological distances.  On the other hand, VLBI is very well matched to detect the very compact radio emission associated with the relativistic outflows generated by accretion onto massive black holes in active galaxies.  High sensitivity VLBI observations, covering a large fraction of sky, thus provide a simple and direct method of identifying these faint, and possibly distant, radio-loud active galaxies.  Free from the effects of dust obscuration, deep and wide-field VLBI studies can therefore contribute to the cosmic census of active galaxies (and their energising massive black holes), and together with redshift information, can potentially probe the accretion history of the early Universe.  In addition, the positions of the compact VLBI detections can be measured very accurately, via the use of standard phase-referencing techniques. Astrometric precision at the mas level can be achieved routinely, and cross identification with sources detected at other wave-bands can anchor these non-radio observations to the International Celestial Reference Frame (ICRF). In crowded and deep fields, such precision astrometry may be useful in identifying the counterparts of faint or obscured sources at other wavelengths.  The observations presented here build on earlier wide-field VLBI studies of the Hubble Deep Field North (HDF-N), conducted by \\citet{gar01} using the European VLBI Network (EVN). Unlike traditional snapshot VLBI surveys, the wide-field approach permits many potential targets in the field to be imaged simultaneously, taking full advantage of the sensitivity and the coverage in the {\\em (u,v)} plane associated with the total duration of the observations and the characteristics of the array.  Using this observing strategy, it is possible to survey many sources (as in the case of snapshot observations) but with much greater sensitivity.  In this paper, we present the results of a deep VLBI survey that covers a region of 1017~arcmin$^2$ $=$ 0.28~deg$^2$, surveying 61 potential sub-mJy and mJy radio source targets at angular resolution ranging from about 10 to 100~mas. For sources located at cosmological distances, the linear resolution corresponding to 10~mas is 90~pc or finer for a {\\em WMAP\\/} cosmology \\citep{spe03}. The VLBI observations and correlation are described in section~\\ref{obs}, while section~\\ref{cal} describes the VLBI calibration and imaging.  Definition of the VLBI survey fields, survey depths, and target-source selection appear in section~\\ref{sel}.  Results of the VLBI survey are presented in section~\\ref{res} and discussed in section~\\ref{dis}.  The paper closes in section~\\ref{fut} with a synopsis of the future prospects for deep, wide-field VLBI surveys. ",
        "conclusions": "\\label{dis}  This paper presents the first attempt to characterise the sub-mJy and mJy radio source population at mas resolutions.  A total of 61 potential targets were selected based on their peak WSRT flux densities, and 9 were detected at or above the 6 $\\sigma$ rms noise level in the VLBI images. Most of the sources are probably located at moderate redshifts ($z \\sim 0.1-1$) but the non-detection of two objects in K-band and one object in both I and K-band, suggests some significant fraction of these faint, compact sources may be located at $z >1$. Our observations suggests that the percentage of mJy and sub-mJy radio sources that are powered by black-hole accretion processes in this Bo\\\"otes field is, at least, 15$\\pm$5\\%.  Only 3 {\\it sub-mJy} WSRT sources (VLBI J142934.7033, J143002.5631, and J142835.5389) are detected out of a total of 40 potential {\\it sub-mJy} targets, implying a detection rate of 8$^{+4}_{-5}$\\%.  The detection rate is higher for {\\it mJy} WSRT sources: 6 sources are detected out of a total of 21 potential {\\it mJy} targets, for a detection rate of 29$^{+11}_{-12}$\\%.   The disappointing detection rate in the sub-mJy radio source population serves to reinforce the rapid evolution that is believed to be taking place in the radio source population at these faint mJy and sub-mJy flux density levels \\citep{mux04}.  However, the detection rate reported here is probably further affected by resolution effects. Several of the sub-mJy sources in our sample are likely to be partially, rather than fully resolved.  In these cases, compact structure associated with an active nucleus may be missed by our survey if it falls below our 6 $\\sigma$ detection threshold.  Our detection rate of about 8\\% for sub-mJy sources is, therefore, very much a lower-limit on the fraction of sub-mJy sources with active nuclei.  The morphology of the sources detected also confirms previously observed trends for brighter sources to exhibit radio structure on much larger scales than fainter sources.  This can be understood in terms of synchrotron-self-absorption theory, wherein the overall radio source size is proportional to the square root of the peak flux density \\citep[e.g.,][]{sne00}."
    },
    "0410/astro-ph0410495_arXiv.txt": {
        "abstract": "We present the results of the timing and spectral analysis of X-ray high-state observations of the high-mass X-ray pulsar SMC X-1 with \\textit{Chandra}, \\textit{XMM-Newton}, and \\textit{ROSAT}, taken between 1991 and 2001. The source has $L_{\\rm X} \\sim (3$--$5) \\times 10^{38}$ ergs s$^{-1}$, and the spectra can be modeled as a power law plus blackbody with $kT_{\\rm BB}\\sim0.18$ keV and a reprocessed emission radius $R_{\\rm BB}\\sim 2 \\times 10^{8}$ cm, assuming a distance of 60 kpc to the source. Energy-resolved pulse profiles show several distinct forms, more than half of which include a second pulse in the soft profile, previously documented only in hard energies. We also detect significant variation in the phase shift between hard and soft pulses, as has recently been reported in Her X-1. We suggest an explanation for the observed characteristics of the soft pulses in terms of precession of the accretion disk. ",
        "introduction": "SMC X-1 is a high-mass X-ray binary system consisting of an X-ray source, a $1.6$ $M_{\\sun}$ pulsar, accreting from a $17.2$ $M_{\\sun}$ companion, Sk 160 \\citep{clar03}. It is also the only X-ray pulsar for which no spin-down episodes have been observed \\citep{kaha99}. The X-ray pulsar has a spin period of $0.71$ s and an eclipsing orbital period of $3.89$ days, and in addition it shows a superorbital variation in flux with period between 45 and 60 days. This is likely due to precession of a warped accretion disk \\citep{wojd98}.  The presence of the accretion disk is likely to have significant consequences for the observed pulse profile. In Her X-1, a low-mass X-ray pulsar with a regular 35 day superorbital period, reprocessing of hard X-rays by the disk gives rise to a pulsating, soft spectral component \\citep{endo00}.  \\citet{zane04} have shown that the phase difference between hard and soft pulses changes as the accretion disk precesses. If similar reprocessing is at work in SMC X-1, as argued by \\citet{hick04}, the pulse profiles should show a variation similar to that of Her X-1. To that end, we have analyzed pulse profiles of SMC X-1 in different epochs from \\textit{Chandra}, \\textit{XMM}, and \\textit{ROSAT}. In $\\S$ 2 we discuss the observations and our methods of data analysis. In $\\S$ 3 we present the results of our analysis, and in $\\S$ 4 we consider the implications of our results and future work. ",
        "conclusions": "\\begin{figure} \\epsscale{1.0} \\plotone{f4.eps} \\caption{Cross-correlation between hard and soft pulse profiles for \\textit{Chandra} and \\textit{XMM} data.} \\end{figure}  In these observations both hard and soft X-rays show pulsations, but differences in the hard and soft pulse profiles indicate a different geometric or physical origin of emission \\citep{paul02}.  Soft pulses in the most luminous X-ray pulsars are likely to originate at the inner edge of the accretion disk, where hard pulses from the neutron star are reradiated at lower energies by disk gases \\citep{hick04}. If this picture is valid for SMC X-1, the observed blackbody component should be consistent with emission from the inner disk. For reprocessing by a partial spherical surface around the neutron star, the blackbody radius $R_{\\rm BB}$ is given by \\citep{paul02} $$R_{\\rm BB}=\\sqrt{\\frac{L_{\\rm X}}{4\\pi\\sigma T^{4}_{\\rm BB}}}.$$ Using $kT_{\\rm BB}\\sim 0.18$ keV and $L_{\\rm X}\\sim 4 \\times 10^{38}$ ergs s$^{-1}$, we have $R_{\\rm BB}\\sim 1.7 \\times$10$^{8}$ cm.  In the standard model for X-ray pulsars, the inner disk radius $R_0$ is close to the magnetospheric radius $R_{m}$, where the magnetic pressure of the dipole field equals the ram pressure of the infalling gas.  However, $R_{m}$ is difficult to estimate for SMC X-1, because the surface $B$-field has not been directly measured. We note that $R_{\\rm BB}$ is close to the corotation radius $R_{\\rm cor} = (GMP^2/4 \\pi^2)^{1/3} \\simeq 1.3\\times 10^{8}$ cm, which is another estimate for $R_0$.  However, the lack of spin-down episodes for SMC X-1 suggests that $R_0 < R_{\\rm cor}$ \\citep{kaha99}. The situation is therefore unclear, but without additional constraints it is certainly possible that the soft pulses originate at the inner accretion disk.  In light of this picture we consider the following pulse profile characteristics: \\begin{enumerate} \\item The hard pulse profile is consistently double-peaked, typically with one dominant peak. \\item The soft pulse profile varies in shape, and has been observed with either a single asymmetric peak or two peaks with varying degrees of symmetry. \\item The pulse profiles exhibit large-scale variations in phase relationship. Closely-spaced observations have revealed identical phase shifts. \\end{enumerate}  \\citet{zane04} suggest that the varying hard-soft phase shift in Her X-1 may be due to the precession of the accretion disk, so we consider the possibility that such a disk could reproduce the above characteristics in SMC X-1. If the observational line of sight is sufficiently close to the plane of the magnetic axis, we might always see two hard pulses, one from each magnetic pole. However, the soft profile could change dramatically in shape and relative phase, as observed, if precession of the disk causes the visible part of the reprocessing region to vary.  We note however that we cannot conclusively verify this picture and its analogy to Her X-1.  First, we lack sufficient sampling of the full 60 day superorbital period.  Second, the hard-soft phase shift appears to vary more rapidly for SMC X-1 (a change of $\\simeq$90$\\degr$ between $\\phi_{\\rm SO} \\simeq 0.16$ and 0.42) than for Her X-1 \\citep[$\\simeq$20$\\degr$ between $\\phi_{\\rm SO} \\simeq 0.03$ and 0.17][] {zane04}.  Third, it is not completely clear even in Her X-1 that the phase variation is caused by the 35 day disk precession; soon after the observations of \\citet{zane04}, the source entered an anomalous low state, an event that may be caused by changes in the inner disk structure \\citep{boyd04,oost01}.  Such changes could alter the shape of the inner reprocessing region, and thus the soft profiles, in a way that is not directly related to the superorbital period.  We conclude that in SMC X-1, we have observed variations in the soft pulse profile that likely reflect changes in the reprocessing of hard X-rays by the accreting gas.  These variations may be related to the superorbital precession of the accretion disk, as has been proposed for Her X-1.  If it can be shown with more observations that the pulse profiles and superorbital period are correlated, then we will have the opportunity to apply significant constraints to the geometry of SMC X-1."
    },
    "0410/astro-ph0410340_arXiv.txt": {
        "abstract": "Using  our extensive radio, optical, near-IR and X-ray imaging and spectroscopy, we consider the reason for the unusually large number of radio detected galaxies, mostly found outside the cluster core, in Abell 2125($z=0.2465$, richness class 4). With 20-cm VLA data, we detect continuum emission from 90 cluster members. The multiwavelength properties of these galaxies suggest that most of the radio emission is due to an enhanced star-formation rate. The dynamical study of \\citet{m03} suggests that Abell 2125 is undergoing a major cluster-cluster merger, with our view within 30 degrees of the merger axis and within 0.2 Gyr of core passage. The combination of projection effects and the physical processes at work during this special time in the cluster's evolution seem likely to be responsible for the unusual level of activity we see in the cluster. We argue that tidal effects on individual cluster members, often far from the cluster core, are responsible for the increased star formation. Our results are consistent with the idea that disk galaxies during this phase of a cluster's evolution undergo rapid evolution, through a burst of star formation, on their way to becoming S0's. ",
        "introduction": "With the recent convergence of supernova and CMB research to a remarkable consensus on the geometrical parameters of the universe, understanding the evolution of galaxies with cosmological epoch is, perhaps, the major goal of extragalactic astronomical research. From the morphology-density relation and the Butcher-Oemler effect, it is clear that the large-scale process of cluster formation must be a key part of the picture. But understanding galaxy evolution may be more difficult than geometric cosmology because of the complex role of dissipational processes, and the confusing results of observing two-dimensional projections of three dimensional physical objects (e.g. galaxies and clusters alike). Using as large a range of wavelengths as possible is one way to improve our understanding of what we are seeing. In this paper, we describe a study, using radio, optical, near-IR and X-ray imaging and spectroscopy, to try to understand how a very rich, very radio-active cluster, Abell 2125,  fits into the bigger picture.   We began our study of the Abell 2125 field with moderately deep 20cm VLA C-array observations \\citep{dow,o99}. We compared Abell 2125 (a richness class 4, blue cluster at $z \\sim 0.25$) with Abell 2645 (an apparently similar cluster at the same redshift but with much redder galaxies). We found a much higher detection rate of radio galaxies (27 vs 4) in Abell 2125, while the detection rate in A2645 is consistent with a normal AGN population in a lower redshift cluster. The excess in Abell 2125 occurs entirely in objects with $L_{20cm} < 10^{23}$ W Hz$^{-1}$, and thus seems consistent with a star-forming population \\citep[e.g.][]{ccb}. However, only a small fraction of these radio-selected galaxies shows optical evidence of so much star formation, leaving the interpretation of the radio results in some doubt. Either the star formation in most of the low-luminosity radio sources is very well hidden at optical wavelengths, or we must seek some other explanation for the radio activity.  Closer examination of the optical and X-ray properties of Abell 2125 shows that it is not a typical, very rich Abell cluster. The galaxy distribution shows a central concentration together with an extension to the southwest, spanning at least 2 Mpc. X-ray observations \\citep{wang} show the same pattern, and that the total X-ray luminosity is low for such a rich cluster. \\citet{wang} suggested that we might be seeing diffuse X-ray emission from a superstructure, as is seen in many cosmological simulations of large scale structure formation. Much of the excess radio-emitting galaxy population from \\citet{o99}, like the diffuse X-ray emission, is contained in the southwest extension, not the cluster core. These factors suggest that the unusual galaxy population, especially as seen at radio frequencies, might be connected to the large-scale nature of Abell 2125 itself.  To study this cluster further, we have made much deeper and higher resolution radio observations using all four VLA configurations. These deep radio data have also motivated deep optical and NIR imaging, optical spectroscopy, deep submillimeter observations, and a long {\\it Chandra} exposure. While some of these ancillary observations were motivated by possibility of studying a background sample, we consider here what they tell us about Abell 2125.   Paper I \\citep{o03a} described our radio, optical and NIR observational program to study the region of the sky centered on Abell 2125. In this paper, we now discuss the properties of the cluster members, particularly in light of the unusually large number of radio detections from cluster's members. Physical quantities are calculated using a WMAP $\\Lambda$CDM cosmology with $H_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$ and $\\Omega_V=0.73$. ",
        "conclusions": "The excess levels of radio emission in Abell 2125 are dominantly related to star formation. Abell 2125 is a cluster-cluster merger seen in projection. Active star formation is not associated with the cluster core but spread fairly uniformly throughout a $2\\times 6$ Mpc region in projection. This activity must be taking place in intermediate density regions, mostly in groups associated with the outskirts of the merging clusters.  The cluster-cluster merger seen in projection accounts for the relatively low X-ray luminosity and temperature, combined with the very high richness and relatively high Butcher-Oemler blue fraction. The excess radio population also is at least partly explained by projection effects. However, the large radio population seems also to require some other mechanism to explain the four to five times larger radio fraction than is seen in other such clusters. This extra factor may be due to the cluster-cluster merger being seen very close to core passage where the interactions of the two cluster systems should be greatest.  It seems likely that star-formation has been stimulated in the outlying group environments associated with the clusters during the time near the core passage. The variations in the tidal field, experienced by each galaxy due to the changing substructure and its position relative to each galaxy seems to be the most likely origin for the enhanced SFRs seen in Abell 2125. The activity in Abell 2125, combined with its dynamical state close to core passage of a major cluster-cluster merger, suggests that this phase in cluster evolution may be particularly important for galaxy evolution.  The authors thank A. Klypin  for useful discussions about the large scale structure and J. Eilek for comments on the text."
    },
    "0410/astro-ph0410389_arXiv.txt": {
        "abstract": "A spectroscopic technique employing an iodine absorption cell (I$_2$) to superimpose a reference spectrum onto a stellar spectrum is currently the most widely adopted approach to obtain precision radial velocities of solar-type stars. It has been used to detect $\\sim$80 extrasolar planets out of $\\sim$130 known. Yet in its original version, it only allows us to measure precise radial velocities of single stars. In this paper, we present a novel method employing an I$_2$ absorption cell that enables us to accurately determine radial velocities of both components of double-lined binaries. Our preliminary results based on the data from the Keck~I telescope and HIRES spectrograph demonstrate that 20-30~ms$^{-1}$ radial velocity precision can be routinely obtained for ``early\" type binaries (F3-F8). For later type binaries, the precision reaches $\\sim$10~ms$^{-1}$. We discuss applications of the technique to stellar astronomy and searches for extrasolar planets in binary systems. In particular, we combine the interferometric data collected with the Palomar Testbed Interferometer with our preliminary precision velocities of the spectroscopic double-lined binary HD~4676 to demonstrate that with such a combination one can routinely obtain masses of the binary components accurate at least at the level of 1.0$\\%$. ",
        "introduction": "The idea of passing starlight through an absorption medium to superimpose a set of reference lines and subsequently precisely measure radial velocities was first proposed by \\cite{Grif:73::}. At that time, \\cite{Grif:73::} used the atmosphere and the telluric H$_{2}$O and O$_{2}$ lines. In order to overcome systematic errors in velocities caused by, among others, the atmospheric pressure changes and hence shifts in the position of the telluric lines, \\cite{Cam:79::} introduced for the first time an absorption cell (containing a toxic hydrogen fluoride gas). This technique was subsequently used to carry out a pioneering precision radial velocity survey of 16 stars over the period of six years \\cite[1981-1987;][]{Cam:88::}. It is conceivable that \\citeauthor{Cam:88::} would have detected the first extrasolar planet if they targeted a larger number of stars as their average velocity precision was $\\sim$13~ms$^{-1}$.  \\cite{Mar:92::} took the absorption cell technique to a new level by introducing an iodine (I$_{2}$) absorption cell with the modeling of a spectrograph point spread function (PSF). The gaseous iodine, besides being much less toxic, has a strong line absorption coefficient (and hence requires a path length of only a few centimeters) and offers a dense forest of absorption lines over the wavelength range of 500-630 nm. In their approach, the Doppler shift of a star spectrum $\\Delta\\lambda$ is determined by solving the following equation \\citep{Mar:92::} \\begin{equation} \\label{i2::} I_{obs}(\\lambda) = [I_{s}(\\lambda+\\Delta\\lambda_{s})\\,T_{I_{2}}(\\lambda+\\Delta\\lambda_{I_{2}})] \\,\\otimes\\,PSF \\end{equation} where $\\Delta\\lambda_{s}$ is the shift of the star spectrum, $\\Delta\\lambda_{I_{2}}$ is the shift of the iodine transmission function $T_{I_{2}}$, $\\otimes$ represents a convolution and $PSF$ a spectrograph PSF. The parameters $\\Delta\\lambda_{s}, \\Delta\\lambda_{I_{2}}$ as well as parameters describing the PSF are determined by performing a least-squares fit to the observed (through the iodine cell) spectrum $I_{obs}$. To this end, one also needs a high SNR star spectrum taken without the cell $I_{s}$ which serves as a template for all the spectra observed through the cell and the I$_2$ transmission function $T_{I_{2}}$ obtained with the Fourier Transform Spectrometer at the Kitt Peak National Observatory. The Doppler shift of a star spectrum is then given by $\\Delta\\lambda = \\Delta\\lambda_{s} - \\Delta\\lambda_{I_{2}}$.  To date, the iodine technique has been successfully used to detect $\\sim$80 out of $\\sim$130 known extrasolar planets \\citep{Sch:04::}. Most of these detections were accomplished by the California-Carnegie group \\citep{Mar:03::}. The iodine technique --- thanks to its conceptual simplicity --- is the most commonly adopted way to obtain precision radial velocities. Iodine absorption cells are available on many spectrographs --- HIRES at the 10m Keck I (Keck Observatory), Hamilton at the 3m Shane (Lick Observatory), UCLES at the 3.9m Anglo-Australian Telescope (Anglo-Australian Observatory), HRS at the 9m HET (McDonald Observatory), MIKE at the 6.5m Magellan (Las Campanas Observatory), UVES at the 8m Kueyen (Cerro Paranal), HDS at the 8.2m Subaru (National Astronomical Observatory of Japan) and many others --- and are used for planet detections. Unfortunately, in its classical version the iodine technique can only be applied to single stars or spectroscopic binaries whose secondaries are so faint that their spectral lines are undetectable. This is dictated by the need to supply an observed template spectrum of each binary component for Eq.~\\eqref{i2::}, but in the case of double-lined spectroscopic binaries it cannot be accomplished since their spectra are obviously always composite and time variable.  In this paper, we present a novel technique which employs an I$_2$ absorption cell to obtain precision radial velocities of both components of double-lined spectroscopic binaries. The technique and its tests are described in Section 2 and its applications to stellar astronomy and searches for planets in binary and multiple stellar systems are discussed in Section 3. ",
        "conclusions": "We have introduced a new method to measure precision radial velocities of double-lined spectroscopic binaries with an iodine absorption cell. Our initial data sets demonstrate 20-30 ms$^{-1}$ velocity precision for ``early\" (F3-F8) type and $\\sim$10~ms$^{-1}$ for later type binaries. If combined with interferometric astrometry, such velocities enable us to determine the masses of the components of binary stars accurate at least at the level of $1\\%$. Our method also makes it possible to perform a search for extrasolar planets in binary and multiple stellar systems where the brightness of the primary and the secondary is comparable. This allows us to probe a regime of the parameter space of the formation and evolution of extrasolar planets not covered by other surveys."
    },
    "0410/astro-ph0410044.txt": {
        "abstract": "We searched the literature to obtain a complete list of known Class\\,0 sources. A list of 95 confirmed or candidate objects was compiled. To the best of our knowledge, all published broad-band observations from 1\\,$\\mu$m to 3.5\\,mm have been collected and are assembled in a catalogue. These data were used to determine physical properties (T$_{\\rm bol}$, L$_{\\rm bol}$, L$_{\\rm smm}$/L$_{\\rm bol}$, M$_{\\rm env}$) and for a uniform classification. 50 sources possess sufficient observational data and are classified as Class\\,0 or Class\\,0/1 objects. The source properties are compared with different evolutionary models to infer ages and masses, and their correlations are investigated. About 25\\% of the sources are found to be in a quiet accretion phase or possess a significantly different time evolution of the accretion rate than the average. In Taurus, with its isolated star formation mode, this seems especially to be the case. ",
        "introduction": "Protostars begin their life soon after the collapse of their parental cloud core begins. Such stellar embryos have masses of less than 10$^{-2}$\\,M$_\\odot$ (Larson \\cite{l03}) and are deeply embedded in a dense and massive envelope of gas and dust. The central object grows in mass by accreting material from this envelope either by direct infall or via a flattened circumstellar disc. As long as the central object is less massive than the surrounding envelope it is called a Class\\,0 object. Otherwise it is a Class\\,1 protostar. Since this mass ratio is not directly observable, other criteria have to be adopted for a classification. Andr\\'e et al. \\cite{2000prpl.conf...59A} defined three observational properties needed to classify an object as a Class\\,0 source: (i) an internal heating source (compact centimeter radio continuum) or bipolar outflow as an indication of a central young stellar object; (ii) extended and centrally peaked sub-millimeter continuum emission, indicating the spheroidal envelope; (iii) a high ratio of sub-millimeter to bolometric luminosity (L$_{\\rm smm}$/L$_{\\rm bol}$\\,$>$\\,0.005), with L$_{\\rm smm}$ measured longward of 350\\,$\\mu$m. It is shown in Andr\\'{e} et al. \\cite{1993ApJ...406..122A} that this last criterion is equivalent to the mass ratio of the envelope and the central star being larger than one. Other authors use the bolometric temperature (T$_{\\rm bol}$) instead of L$_{\\rm smm}$/L$_{\\rm bol}$ (e.g. Chen et al. \\cite{1995ApJ...445..377C, 1997ApJ...478..295C} use T$_{\\rm bol}$\\,$<$\\,70\\,K).  The Class\\,0 stage can be characterised as the main mass accretion phase, in which the forming star gains the bulk of its final mass. This process is not yet well understood. How does the mass accretion rate depend on the age and the properties of the cloud, and is it governed by turbulence or ambipolar diffusion? What is the lifetime of the Class\\,0 sources? How are the source properties connected to the molecular outflow, that is inevitably driven by these objects? The answers to these questions will finally help us to understand the feedback of the star formation process to the surrounding molecular cloud and the initial mass function.  To seriously address these questions we need a large, homogeneously classified sample of these very young protostellar objects. Several samples have been compiled in the recent years (e.g. Chen et al. \\cite{1995ApJ...445..377C}, Andr\\'e et al. \\cite{2000prpl.conf...59A}, Shirley et al. \\cite{2000ApJS..131..249S}, Motte \\& Andr\\'{e} \\cite{2001A&A...365..440M}). All these works have specific problems: (1) they concentrate only on a limited number of sources; (2) different criterias for classification of Class\\,0 objects are used; (3) only parts of the available observational data are considered.  The first step to obtain a large, homogeneously classified sample of Class\\,0 sources is a complete search of the literature for such objects. Here we present an as far as possible complete list of the youngest known protostars to date. We will classify them by means of their spectral energy distribution (SED) using all (to the best of our knowledge) published broad-band photometric data. We derive basic properties of these sources (T$_{\\rm bol}$, L$_{\\rm bol}$). The sub-mm data are used to estimate envelope masses (M$_{\\rm env}$) and sub-mm slopes of the SED ($\\beta$). By comparing T$_{\\rm bol}$, L$_{\\rm bol}$, and M$_{\\rm env}$ with evolutionary models we determine ages and masses of these sources and investigate limitations of the current observational data and evolutionary models. This is of particular interest for forthcoming powerful telescopes like Spitzer and ALMA, which will enable us to improve our knowledge of these sources significantly. Given this expected growth in observational data in the near future, the present catalogue might be quickly outdated. Hence, all the data will be available on a webpage\\footnote{http://www.dias.ie/protostars} and updated regularly.  In Sect.\\,\\ref{sample} we introduce the source sample. The method of our data analysis is then put forward in Sect.\\,\\ref{analysis}, including a description of the source classification in Sect.\\,\\ref{classification}. Finally the results and implications are discussed in Section\\,\\ref{discussion}.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%  Source Sample  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "A literature search for photometric broadband observations of 95 confirmed or candidate Class\\,0 sources was conducted. To our best knowledge all available broadband photometric data was used to construct SEDs from 1\\,$\\mu$m to 3.5\\,mm. If possible we determined basic properties of the sources (sub-mm slope of the SED, T$_{\\rm bol}$, L$_{\\rm bol}$, L$_{\\rm smm}$/L$_{\\rm bol}$, and M$_{\\rm env}$). For 59 objects sufficient enough data are available for a proper determination of the source parameters. 27 of these are classified as Class\\,0, 23 as Class\\,0/1, and 9 as Class\\,1 protostars.  To investigate the statistical properties of the obtained sample of very young protostars we used all objetcs within 500\\,pc which are properly classified as Class\\,0 or Class\\,0/1 to determine age and final star mass. Therefor we used three different evolutionary models for protostars from Smith \\cite{1998Ap&SS.261..169S, 2000IrAJ...27...25S, s02}, Myers et al. \\cite{1998ApJ...492..703M}, and Andr\\'e et al. \\cite{2000prpl.conf...59A}. Considering the uncertainties in the measured source properties the predicted final star masses from the different models show a good correlation. The absolute ages are, however, very model dependent.  An investigation of the final star masses and the resulting mass function shows a good agreement with the IMF for stars with M\\,$>$\\,0.5\\,M$_\\odot$. Our obtained Class\\,0 sample is limited to sources with final masses above 0.4\\,M$_\\odot$. A number of objects (25\\,\\%) is found that possess a much lower luminosity than the rest of our sample, considering their bolometric temperature and envelope mass. These objects might be in a 'quiet' accretion phase or the time evolution of their mass accretion rate is significantly different from the majority. This might especially be the case in the Taurus star forming region, where all identified Class\\,0 or Class\\,0/1 sources belong to this group.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%   Acknowledgements   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0410/cond-mat0410557_arXiv.txt": {
        "abstract": "The field theory of a fluctuating d-wave superconductor is constructed and proposed as an effective description of superconducting cuprates at low energies. The theory is used to resolve a puzzle posed by recent experiments on superfluid density in severely underdoped $YBa_2 Cu_3 O_{6+x}$. In particular, the overall temperature dependence of the superfluid density at low dopings is argued to be described well by the strongly anisotropic weakly interacting Bose gas, and thus approximately linear in temperature with an almost doping-independent slope. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410426_arXiv.txt": {
        "abstract": "It has been an interesting question that why there are one-third of Kuiper Belt Objects (KBOs) trapped into the 3:2 resonance but, in contrast, only several KBOs are claimed to be associated with the 2:1 resonance. In a model proposed by Zhou et al. (2002), the stochastic outward migration of the Neptune could reduce the number of particles in the 2:1 resonance and thus the objects in the 3:2 resonance become more distinct. As a complementary study, we investigate the effect of proto-stellar discs on the resonance capture. Our results show that the gaseous drag of a proto-stellar disc can trap KBOs into the 3:2 resonance rather easily. In addition, no objects are captured into the 2:1 resonance in our simulation. ",
        "introduction": "The general picture of the outer solar system has been changed due to the discovery of KBOs (Jewitt and Luu 1993, Williams et al. 1995). The origin of the dynamical properties of KBOs has become a very interesting but controversial issue. Particularly, it is surprising that one-third of the population are engaged into the 3:2 resonance with the Neptune.  Malhotra (1995) proposed a mechanism of resonance sweeping to explain the large fraction of 3:2 resonance objects, plutinos, in the Kuiper Belt. In that model, as the Neptune radially migrate outward, its mean motion resonance swept ahead of it through the Kuiper Belt and captured the Pluto, along with plutinos, into the 3:2 resonance. This theory predicted that the populations of the 3:2 and 2:1 resonance should be of the same order. Moreover, the radial migration of the Neptune is assumed to be smooth and continuous. Zhou et al. (2002) improved this by introducing a stochastic term to make this outward migration more realistic. They showed that for particular chosen parameters, the number of particles in the 2:1 resonance can be reduced and thus the objects in the 3:2 resonance become more distinct.  Yeh and Jiang (2001) showed that the scattered planet shall move on an eccentric orbit and thus the pure radial migration is too naive, although it could be an approximation. For example, the Neptune's orbit is not circular during its outward migration as shown in the simulations by Thommes et al. (1999). However, the Neptune's current orbit is nearly circular and thus one needs to introduce a massive Kuiper Belt to circularize the Neptune's orbit. Please also see the simulations in Gomes (2003).   In fact, the total mass of Kuiper Belt is itself a controversial issue. In order to form  100 km size KBOs around the current region, one would need a much more massive Kuiper Belt initially because the planetesimal accretion rate has to be high enough to make KBOs formed in time. However, the material in the Kuiper Belt has to be depleted significantly during the evolution because the observational upper limit of the mass at the Kuiper Belt is only 0.1 Earth mass. To avoid this formation problem, Levison and Morbidelli (2003) proposed a pushing-out model that both the Kuiper Belt Objects and Neptune were forced to migrate outward from where they were formed initially. They showed that some KBOs were captured into 2:1 resonance and could move out together with the Neptune in their simulations.   On the other hand, the traditional planet formation theory, i.e. core-accretion model, is facing the competition from the disc-instability model (Boss 1998). Because the disk-instability model can form a Jupiter mass planet in a few thousand years, it could be that the formation of the Neptune and KBOs is not that difficult as the traditional model. Although the solid core formation through dust grain growth probably still play an important role, the condition to initiate the formation of Neptune and KBOs and thus the formation timescales could be very different. Furthermore, Bryden et al. (2000) showed that even by the core-accretion model, it is possible to form the Neptune around the current location in $10^6$ or $10^7$ years, which is comparable or shorter than the disc depletion timescale. In any case, the higher density gaseous blobs might shorten the dust-accretion process. Therefore, it is not clear how severe the initial mass problem is.  No matter what the exact formation scenario is, our model here could represent a process that starts at the time (a) when the proto-Neptune and proto-KBOs were already pushed out after a violent migration and settle down around current places (if it is  difficult to form at current locations)  or (b) after the Neptune and KBOs have grown and formed around current regions with only very gentle migrations.  We are particularly interested in the influence from the proto-stellar disc and thus the possible migration is ignored here. The model will be described briefly in Section 2 and the results would be in Section 3. There are some discussions in Section 4. Section 5 concludes the paper. ",
        "conclusions": "We have used a model of the disc-star-planet system to investigate the effect of a gaseous proto-stellar disc on the resonance capture. In addition to the force from the central star and planet, the test particle is also influenced by the gravitational and frictional forces from the disc. Particularly, we apply our model to the problem of resonant KBOs and focus on the 3:2 and 2:1 resonances.  Our results show that the drag force plays an important role for the resonant capture and many KBOs can get captured into the 3:2 resonance but not into the 2:1 resonance. This is consistent with the observation that about one-third of the whole population of KBOs are in the 3:2 resonance and only a few are claimed to be in the 2:1 resonance. Therefore, the mechanism we study here might be helpful to understand the history of KBOs' resonance capture."
    },
    "0410/astro-ph0410610_arXiv.txt": {
        "abstract": "We consider the problem of tidal disruption of stars in the centre of a galaxy containing a supermassive binary black hole with unequal masses. We assume that over the separation distance between the black holes the gravitational potential is dominated by the more massive primary black hole.  Also, we assume that the number density of stars is concentric with the primary black hole and has a power-law cusp. We show that the bulk of stars with a small angular-momentum component normal to the black-hole binary orbit can reach a small value of total angular momentum through secular evolution in the gravitational field of the binary, and hence they can be tidally disrupted by the larger black hole.  This effect is analogous to the so-called Kozai effect (Kozai 1962, Lidov 1961, 1962) well known in celestial mechanics. We develop an analytical theory for the secular evolution of the stellar orbits and calculate the rate of tidal disruption. We confront our analytical theory with a simple numerical model and find very good agreement.  Our results show that for primary black-hole mass $\\sim 10^{6}-10^{7}M_\\odot$, the black-hole mass-ratio $q > 10^{-2}$, cusp size $\\sim 1\\pc$, the tidal disruption rate can be as large as $\\sim 10^{-2}-1M_\\odot/\\yr$. This is at least $10^{2}-10^4$ times larger than estimated for the case of a single supermassive black hole. The duration of the phase of enhanced tidal disruption is determined by the dynamical-friction time scale, and it is rather short: $\\sim 10^{5} \\yr$.  The dependence of the tidal disruption rate on the mass ratio, as well as on the size of the cusp, is also discussed. ",
        "introduction": "The problem of the tidal disruption of the stars by a supermassive black hole has attracted the attention of many researches starting from the pioneering paper of Hills 1975, who proposed that the stellar gas released from the tidally disrupted star could power quasars and AGN's. Later, it has been realized that the tidal disruption event could manifest itself as a luminosity flare in the centre of a galaxy (e.g., Gurzadian \\& Ozernoi 1981, Lacy et al.\\ 1982, Rees 1988).  The observational detection of these flares have recently been reported (e.g., Komossa et al.~2003 and references therein).  In order to be tidally disrupted a star must have its orbital pericentre sufficiently close to the central black hole, and accordingly, its angular momentum must be sufficiently small. Thus, the presence of the central black hole give rise to a `loss cone' in the distribution of stars in angular-momentum space (e.g., Frank \\& Rees 1976), where the number of the stars is quickly exhausted due to tidal disruption, and the process of refilling the loss cone determines the rate of tidal disruption. Estimates of the rate of tidal disruption are typically based on the assumption of two-body gravitational interactions between the stars as the dominant process of angular-momentum change. Since the relaxation time associated with this process is large for realistic models of galactic centres, these estimates give rather low values of the disruption rate, of the order of $10^{-7}$ to $10^{-4}M_\\odot/\\yr$ depending on the properties of the stellar population of the galactic centres and the mass of the black hole residing in the centre (e.g., Frank \\& Rees 1976, Lightman \\& Shapiro 1977, Shapiro \\& Marchant 1978, Cohn \\& Kulsrud 1978, Magorrian \\& Tremaine 1999, Syer \\& Ulmer 1999, Wang \\& Merritt 2004 and references therein).  It has been pointed out by Komberg 1968 and by Polnarev and Rees 1994 that a secondary black hole of a smaller mass could intrude in the galactic centre as a result of the process of merging of two galaxies and perturb stellar trajectories, leading to non-conservation of the angular momentum and an enhancement of the tidal disruption rate (for a review of astrophysical processes occuring in galactic centres containing a supemassive black hole binary see e. g. Merritt $\\&$ Milosavljevic 2004 and references therein). The purpose of this paper is to calculate this enhanced tidal disruption rate in a simple model of interaction between the binary black hole and the stars.  We assume that the binary black hole consists of two black holes of unequal masses moving around each other in circular orbits deep inside the potential well of the primary black hole.  We take into account the gravity of the central stellar cluster in the form of small corrections to the motion of stars. It has been pointed out that when the gravity is dominated by the central black hole, the number density of stars increases toward the centre according to a power law and such a distribution of the stars has been called `the cusp' (Peebles 1972, Bachall \\& Wolf 1976, Young 1980).  We use this power law distribution as the initial condition for our problem. We consider only stars with orbits within the orbit of the binary (which we will call the `inner problem').  We show that in this case the main process leading to the replenishment of the loss cone is a secular effect increasing the eccentricities of the stars called `the Kozai effect' (Kozai 1962, Lidov 1961, 1962).  The qualitative explanation of the process is very simple.  Let us assume as a zeroth approximation that the stars move on Keplerian orbits in the gravitational field of the primary black hole and that gravitational interaction with the secondary black hole introduces slow secular changes to the orbital elements.  Neglecting the possibility of resonance interactions and considering the evolution of the orbital parameters on a time scale scale much larger than the orbital periods of the star and of the binary, it is well-known that the secondary black hole can be approximated by its orbital average, which is a mass ring.  Since the gravitational field of a ring is static and axially symmetric, the energy and the angular-momentum component normal to the ring plane are conserved. However, in the course of the secular evolution the total angular momentum varies with time, and it can be as small as the constant normal component.  Since the number of stars whose normal component of angular momentum lies within loss-cone size is much larger than the number of stars whose total angular momentum lies within the loss cone, the number of stars which can be tidally disrupted is significantly increased by the presence of the secondary black hole.  In the next Section we formulate the model problem, which takes into account the secular evolution of the orbital elements of the stars from both the secondary black hole and the stellar cluster. The analytical solution of the secular problem is discussed in Section 3. The tidal disruption rate based on this solution is evaluated in Section 4. In Section 5 we present a numerical approach to the problem and check the validity of our analytical results. In Section 6 we use the results obtained in the previous Sections to estimate the enhanced tidal disruption rate in astrophysical systems.  This Section is written in a fairly self-contained way, so readers interested in astrophysical applications but not in the derivations of the relevant equations can go directly to this Section. Section 7 presents our conclusions and discussion. ",
        "conclusions": "We have shown in this paper that the tidal disruption rate in a galactic cusp containing a circular black hole binary can be as large as $\\sim 0.1$--$1 M_\\odot/\\yr$ for a cusp size of the order of $1\\pc$. Cusps of this size have been observed in e.g. our own Galaxy and M32 (e.g. Kormendy $\\&$ Gebhardt 2001 and references therein). In principle, the rate of gas supply to the centre of the galaxy is sufficient to power a quasar or a powerful AGN. However, the estimated duration of the stage of the enhanced tidal disruption rate $\\sim 10^{5} \\yr$ seems to be much shorter than the expected life time of the central engine in quasars. Also, note that the tidal disruption rate sharply decrease with increase of the cusp size $r_{a}$: $\\propto r_{1}^{-2}$, where $r_{1}=r_{a}/1pc$. The obsevations suggest that the cusp size grows with mass of the central black hole with the typical sizes in the range $10-30\\pc$ for the masses of central black holes of the order of $10^{7}-10^{8}M_{\\odot}$ (e. g. Kormendy $\\&$ Gebhardt 2001). So for an extended cusp of size $\\sim 10\\pc$ we get rather moderate values of the tidal disruption rate of the order of $1\\cdot 10^{-3}-1\\cdot 10^{-2}M_\\odot/\\yr$.  The tidal disruption events in a system containing a supermassive black binary black hole could have observational consequences on its own. A single tidal disruption event could manifest itself as sharp rise of luminosity of a galactic centre on time scale $\\sim 1\\yr$ (Gurzadian \\& Ozernoi 1981, Lacy et al 1982, Rees 1988, Komossa et al 2003 and references therein).  The repeating luminosity flares in a particular galaxy could suggest that there is a binary black hole in the centre. On the other hand, one can expect the increased probability of the tidal disruption events in the supermassive binary black hole candidates (e.g., Sillanpaa et al 1988, Valtaoja et al 2000).  We assume that the stellar distribution has a pronounced density increase toward the centre. This assumption is well motivated from theoretical point of view and also has been confirmed by analysis of the distribution of stars inside our own Galaxy (e.g., Alexander 1999) and some other galaxies (e.g., Faber et al.\\ 1997).  However, after the intervention of the secondary black hole and coalescence with the primary, a large fraction of the stars is expelled from the centre, and the density increase is smeared out. This may explain rather shallow density profiles observed in the centers of bright galaxies (e.g., Faber et al.\\ 1997) Thus, subsequent coalescence events of two black holes within the centre of the same galaxy would be accompanied by smaller numbers of tidally disrupted stars.  We have considered only the `inner problem' where all stars going to the loss cone were assumed to be inside the binary orbit and showed that for the inner problem the bulk of the stars going to loss cone would have had a small value of the projection of the angular momentum on the axis perpendicular to the binary orbital plane before intervention of the secondary in the cusp. The mechanism responsible for the supply of the stars into the loss cone is analogous to the well known Kozai effect in celestial mechanics. However, the important effect of the oscillation of the angular momentum on the time scale of the order of the binary orbital period should also be taken into account in order to obtain the correct expression for the tidal disruption rate in a binary with a sufficiently large mass ratio. Also, we take into account the apsidal precession determined by gravitational potential of the stellar cluster itself and develop the analytical theory of secular evolution in the combined gravitational field of the binary and the stellar cluster. This theory might be useful in other astrophysical applications, such as e.g., the theory of dynamics of the triple stars (e.g., Kiseleva, Eggleton \\& Mikkola 1998 and references therein), and the theory of dynamics of extra-solar planets around the binary or a single star (e.g., Papaloizou \\& Terquem 2001).  The tidal disruption rate may be significantly enhanced if we take into account `the outer' stars with semi-major axes larger than the half of the binary separation distance. The evolution of these stars in the binary gravitational field is mainly governed by two processes: the secular non-resonant interaction with the binary similar to the interaction discussed in our paper, and the effect of close encounters with the secondary which change the orbital elements of the star in a complicated chaotic way. Unfortunately, this problem looks very complicated for a thorough theoretical treatment and numerical calculations are very important in this case.  Our estimate of the rate of supply of the stars to the tidal loss cone relied upon the effect of the secular increase of the eccentricity. We have analysed this effect in an idealised situation where only the Keplerian gravitational field of the binary and corrections due to the stellar cluster have been taken into account and the binary has had a circular orbit with the fixed separation distance between the primary and the secondary. In a real astrophysical situation the binary could have a substantial eccentricity (e.g., Quinlan $\\&$ Hernquist 1997, Aarseth 2003) and the orbital parameters of the binary are evolving with time due to the effect of dynamical friction. Therefore, a very important question arises: to what extent is the effect of the secular increase of the eccentricity of the stellar orbits sensitive to possible corrections due to non-zero eccentricity and to a change of the orbital elements of the binary?  First we would like to point out that when the eccentricity of the binary is not very large the effects determined by the eccentricity seem to be rather insignificant for the inner problem considered above. Indeed, in the quadrupole approximation used in our analysis the gravitational field of the eccentric binary is formally equivalent to the gravitational field of the circular binary provided that we use the parameter of the binary orbit $p_{B}=a_{B}(1-e_{B}^{2})$ instead of the separation distance $D$ in equation (14) defining the characteristic time-scale of the Kozai effect, $T_{K}$ (e.g Innanen, Zheng, Mikkola and Valtonen 1997).  Of course, in the case of a highly eccentric binary the influence of the next order corrections in the small parameter $a/D$ would be more prominent, but we expect that the qualitative character of the secular evolution would remain unchanged until the condition $a/D < 1$ holds for the stellar orbits of interest. When the quadrupole approximation is valid, a slow change of the orbital parameters of the binary also seems to be unimportant. In this case this time dependence enters the equations of the secular evolution (10--13) only through the time dependence of the quantity $T_{K}(t)$ in front of the time derivatives. Introducing the new time variable $t_{1}=\\int dt/T_{K}$ we can eliminate this time dependence, and the secular evolution is the same in terms of the new time variable $t_{1}$ as it was in the case of the non-evolving binary orbital elements.  In this Paper we consider the black holes as point like sources of Newtonian gravity. However, as we discussed in the previous Section, the relativistic effect of Einstein apsidal precession (and, possibly, the effect of gravitomagnetic precession) could be important. A study of these effect could be a very interesting problem deserving a separate study.  In conclusion, we would like to point out that for a quantitative calculation of the tidal disruption rate our approach is perhaps too oversimplified, and a thorough numerical study is needed.  However, our model can be used for the order of magnitude estimates discussed above. It can also be used for guidance on the choice of parameters in numerical studies."
    },
    "0410/astro-ph0410560_arXiv.txt": {
        "abstract": "\\noindent X-ray Flashes (XRFs) are, like Gamma-Ray Bursts (GRBs), thought to signal the collapse of massive stars in distant galaxies. Many models posit that the isotropic equivalent energies of XRFs are lower than those for GRBs, such that they are visible from a reduced range of distances when compared with GRBs. Here we present the results of two epoch {\\it Hubble Space Telescope} imaging of two XRFs. These images taken approximately 45 and 200 days post burst reveal no evidence for an associated supernova in either case. Supernovae such as SN~1998bw would have been visible out to $z \\sim 1.5$ in each case, while fainter supernovae such as SN~2002ap would have been visible to $z \\sim 1$. If the XRFs lie at such large distances, their energies would not fit the observed correlation between the GRB peak energy and isotropic energy release ($E_{p} \\propto E_{iso}^{1/2}$), in which soft bursts are less energetic. We conclude that, should these XRFs reside at low redshifts ($z<0.6$), either their line of sight is heavily extinguished, or they are associated with extremely faint supernovae, or, unlike GRBs, these XRFs do not have temporally coincident supernovae. ",
        "introduction": "X-ray Flashes (XRFs) occupy the extreme left (lower energy) wing of the peak energy distribution of Gamma-Ray Bursts (GRBs). While GRBs radiate the majority of their energy in $\\gamma$-rays ($E_{p}\\sim200$ keV, Preece \\etal\\ 2000), XRFs are characterized by peak energies below 50 keV and an X-ray fluence in excess of that observed in $\\gamma$-rays.  XRFs were discovered with the {\\it BeppoSAX} Wide Field Cameras (Heise \\etal\\ 2001a); they constitute approximately 1/3 of the total burst population (Lamb \\etal\\ 2004).  The physical mechanisms driving XRFs and their observed differences to GRBs are some of the key questions in the field today.  Of the many competing theories relating to the origin of XRFs, the most popular is that XRFs are produced when a classical GRB is observed ``off-axis'', so that the highly collimated ejecta (i.e., the highest energies and harder photons) are not seen (Yamazaki, Ioka \\& Nakamura, 2002; Rhoads 2003; Dado et al. 2004). Alternatively, models invoking either an increase of the baryon load within the fireball itself (Dermer, Chiang \\& B\\\"ottcher 1999; Ramirez-Ruiz \\& Lloyd-Ronning 2002; Meszaros \\etal\\ 2002; Huang, Dai \\& Lu 2002) or low efficiency (high baryon load) shocks (Zhang \\& Meszaros 2002; Barraud et al. 2003) can also produce XRFs.  One interesting consequence of all of the models above, is that they predict a distance distribution for XRFs which is foreshortened with respect to that of GRBs, since XRFs emit a fraction of the energy of the GRB itself.  This hypothesis seems to be supported by the analysis of GRBs with known redshifts, whereby the event peak energy $(E_{p}$) scales as $E_{iso}^{1/2}$, i.e. softer GRBs have lower isotropic energies (Amati \\etal\\ 2002; Lloyd-Ronning \\& Ramirez-Ruiz 2002). Qualitative arguments on XRF distances were presented by Kouveliotou \\etal\\ (2004), where a comparison of the energetics of the X-ray afterglows of GRBs and XRFs suggested that were the latter placed a high redshift (e.g. $z=1.5$), their luminosities would rival those of the brightest GRBs, suggesting a closer proximity for XRFs.   Qualitative distance arguments have also been presented by Kouveliotou \\etal\\ (2004), where a comparison between the energetics of GRBs and XRFs indicated that if XRFs were placed further than $z=1.5$ their luminosities would exceed those of the brightest GRBs, suggesting a closer proximity for XRFs.  However, determining the true distances of XRFs has proven to be exceptionally difficult observationally.  Only two XRFs have so far been associated with optical afterglows, XRF~020903 (Soderberg \\etal\\ 2004) and XRF~030723 (Fynbo \\etal\\ 2004). Of these only the former has a measured spectroscopic redshift ($z=0.25$; Soderberg \\etal\\ 2004) while the latter has only  a redshift limit ($z<2.3$) with the suggestion of low redshift based on the presence of a supernova like component it its late time light curve (Fynbo \\etal\\ 2004). A second possible case, GRB/XRF 031203 is not yet settled, as there is an ongoing debate on the true nature of this nearby event ($z=0.105$), which was originally classified as an XRF (Watson \\etal\\ 2004) based on its X-ray brightness as inferred from the discovery of a halo from dust scattering in our own Galaxy (Vaughan \\etal\\ 2004). However, an analysis of its high energy spectrum recorded with {\\it INTEGRAL} (Sazonov \\etal\\ 2004) revealed an $E_p\\sim190$ keV, apparently inconsistent with that inferred from the dust halo. The paucity of distance measurements thus prevents us from deriving definitive conclusions on the XRF distance distribution.  The spectroscopic signatures of type Ic supernovae seen in GRB~030329/SN~2003dh (Stanek \\etal\\ 2003; Hjorth \\etal\\ 2003), and the discovery of photometric bumps in many GRB optical afterglow light curves (e.g. Zeh, Klose \\& Hartmann 2004), strongly support the collapsar model (Woosley 1993), whereby the majority of long duration GRBs are the result of core collapse supernovae. Should XRFs be associated with supernovae, as seen in GRBs, but found at typically lower redshifts, it should be possible to constrain their distances by the detection of a supernova brightening in their X-ray located host galaxies. This can be the case where an optical afterglow is seen (e.g. XRF 030723; Fynbo \\etal\\ 2004) or in some cases where there was no apparent optical afterglow (e.g., this was recently the case for GRB/XRF 031203 (Bersier \\etal\\ 2004; Thomsen \\etal\\ 2004;Cobb \\etal\\ 2004; Gal-Yam \\etal\\ 2004; Malesani \\etal\\ 2004)).  Here we present results on our {\\it HST} data of two XRF host galaxies, identified earlier by a combination of {\\it HST} and {\\it Chandra} data (Fruchter \\etal\\ 2002a,b; Bloom \\etal\\ 2003). In the present study, we concentrate on the difference imaging between two {\\it HST} epochs separated by six months, and the constraints these results place on underlying supernovae. Finally we combine these results with the properties of the X-ray lightcurves of each XRF and their similarities with GRBs to reach conclusions on the likely nature of XRFs. ",
        "conclusions": "We have presented deep, multi-epoch \\hst\\ observations of the fields of two XRFs. These observations, reaching depths of $V \\sim 28.5$ would be sensitive enough to detect underlying supernovae such as SN~1998bw or SN~2002ap out to $z >1$ even with moderate host galaxy extinction. At very low redshifts the extinction required to remove the supernova flux is very large, and incompatible with the measured column density in the case of XRF~020427. At higher redshifts XRF 020427 does not fit the $E_p$ - $E_{iso}$ relationship. This situation can be reconciled if the XRFs are situated locally behind large dust absorbing columns, if the supernovae are very faint, or if the SN and XRF are not temporally coincident. A final option is that GRBs and some XRFs do not represent the same physical process, and that they may be due to similar but physically distinct phenomena, where a supernova is not required for the latter.  The recent XRF~040701 may allow greater insights into the XRF-SN association, or the lack thereof. Optical observations failed to locate either a supernova or an optical afterglow (Berger \\etal\\ 2004; Pian \\etal\\ 2004), while {\\it Chandra} observations did find a candidate X-ray afterglow (Fox \\etal\\ 2004) and a coincident galaxy system at low redshift $z=0.22$. Deep {\\it HST} observations of this object may significantly aid our understanding of the XRF phenomenon."
    },
    "0410/astro-ph0410083_arXiv.txt": {
        "abstract": "{ One of the least biased complete Seyfert sample available to date has been considered to perform an X-ray survey using the capabilities of XMM-Newton and Chandra satellites. The X-ray nuclear luminosities obtained have been used to perform a multi-wavelength analysis as a powerfull tool in the identification of heavily obscured objects. The new X-ray data, together with highest quality nuclear fluxes available from the radio to the optical band, allow us to obtain nuclear Spectral Energy Distributions (SEDs) for type 1 and type 2 objects in our sample. The results obtained from the analysis of the SED have been discussed in the framework of Unified Models. } \\authorrunning{F. Panessa et al.} \\titlerunning{Multi-wavelength nearby Seyfert galaxies} ",
        "introduction": "Active Galactic Nuclei (AGNs) produce a large amount of energy in small spatial scales and radiate their energy over a wide range of energies, from gamma-rays and X-rays through the ultraviolet, optical and infrared spectral regions to the far-infrared and radio-frequency regions. X-ray emission is arguably the most important aspect of understanding AGNs, because of its unambiguous association with genuine nuclear activity, and its important diagnostic capabilities for studying accretion mechanisms.  Nonetheless, the study of the continuum emission of Seyfert galaxies over a broad range of frequencies is an important tool to understand/discriminate the energy output in different wavebands, as well as a mean to distinguish galaxies of different activity classes and to test Unified Models which ascribe the differences between type 1 and type 2 Seyferts as merely due to orientation effects rather than fundamental physical differences. Previous works, however, were limited since they concentrated nearly exclusively on high-luminosity AGNs (L$_{bol}$ $\\geq$ 10$^{44}$ erg/s) and very little data exist on the spectral properties of low-luminosity sources (low-luminosity Seyfert galaxies and Liners, with L$_{bol}$ $\\sim$ 10$^{41}$-10$^{43}$ erg/s). To date, only a handful of objects have been adequately studied with multi-wavelength observations obtained with sufficient angular resolution in order to detect the nuclear emission and separate it from the host galaxy background.  We have performed an extensive X-ray study of a complete sample of AGNs unbiased against low-luminosity sources using the capabilities of the \\cha and \\xmm observatories. We obtained high quality nuclear SEDs for type 1 and type 2 objects in our sample from the radio to the hard X-ray band, using the combination of high angular resolution capabilities and unprecedented sensitivity of the new generation telescopes. We refer to forthcoming publications for details on the results here presented (Panessa et al. in preparation, Cappi et al. in preparation). ",
        "conclusions": ""
    },
    "0410/astro-ph0410252.txt": {
        "abstract": "The bipolar nebula, IRAS 19312+1950,  is a unique  SiO maser source exhibiting both properties as young and evolved objects. To clarify the nature of this object,  we made molecular line observations with the Nobeyama 45-m radio telescope. We detected emission from O-bearing (HCO$^{+}$,  SiO,  SO, and SO$_2$), C- and N-bearing molecules  (CN, CS, HCN, HNC, NH$_{3}$, N$_{2}$H$^{+}$, HC$_{3}$N, H$_2$CS, and CH$_3$OH), and their isotopic species (C$^{17}$O,  $^{13}$C$^{18}$O,  and C$^{34}$S). Line profiles consist of  a weak broad ($\\Delta v\\sim 30$ km s$^{-1}$) and/or  a strong narrow ($\\Delta v\\lesssim 5$ km s$^{-1}$) components, depending on molecular species. Strong time variations of H$_2$O and SiO masers were also observed.  Numerical modeling of the envelope with the LVG-code results in a good fit of the model with a mass loss rate of $2.6 \\times10^{-4} M_{\\odot}$ yr$^{-1}$  to the observed intensities for the broad-component lines.  On the other hand, non-O-bearing molecules, which have only the narrow profiles, were found to have abundances typical to  those in  cool dust clouds. No isotopic enrichment was found, indicating little evidence of the narrow cool component being an ejecta of  the central AGB star or a possible companion. These facts compelled us to conclude  that IRAS 19312+1950 is an exotic mass-losing evolved star embedded in a low-mass ($\\sim 20 \\; M_{\\odot}$)  dark cloud. ",
        "introduction": "IRAS 19312+1950  exhibits a prominent bipolar nebulosity of the size of about 30$''$ in near-infrared images  (\\cite{nak00}). The IRAS color of this source, $log(F_{25}/F_{12})=0.47$, indicates a presence of cool dust of a temperature of about 200 K toward this source. Detecting SiO and H$_{2}$O maser lines in this object, \\citet{nak00} concluded that this is an oxygen-rich protoplanetary nebula similar to OH231.8+4.2. The distance to this object was estimated to be about 2.5 kpc  (\\cite{nak04}; hereafter Paper I), based on the near- and mid-IR flux densities. It is slightly unusual that CO and the other thermal molecular lines of less abundant species are detectable in the envelope of an Asymptotic Giant Branch (AGB) or post-AGB stars at 2.5 kpc away from the sun. For example, the distance to the similar object, OH231.8+4.2, is about 1.3 kpc (\\cite{mor87}), and those of  other commonly observed AGB stars in radio are less than 1 kpc (\\cite{olo98}). Near-IR imaging of this object with the UH 88-inch telescope shows a bright central star and surrounding nebula, extending from NE to SW by about 30$''$ (figure 1) . It also shows a small ($\\sim 5''$) ring-like structure near the central star, which is elongated perpendicularly to the bipolar axis.  In paper I,  we reported the results of mapping observations by several molecular lines (CO, HCN, and CS) using the Nobeyama 45-m radio telescope; the line profiles exhibit broad and/or narrow components depending on molecular species. We found  that both broad and narrow components centered at $V_{\\rm lsr}\\simeq$ 35 km s$^{-1}$ are strongly peaked in intensity at the IRAS position of this source with a 16$''$ telescope beam. Without doubt, the broad component with a full width of more than 20 km s$^{-1}$ must originate from the outflowing envelope of the central star. However, the spatial association of the narrow component to this object is a puzzle. The mass of the narrow component was estimated to be about 10 $M_{\\odot}$ (within 16$''$ beam). The mass obtained is compatible with the core masses of dark clouds (for example, \\cite{ful91}).  In this paper, we present a part of  molecular-line observations toward this object, and discuss on it from  a view point of circumstellar chemistry. In usual cases, molecular species found in dark clouds are considerably different from those found in the evolved objects [see, \\citet{gla96} and \\citet{ehr00}]. Therefore, the origin of the cloud surrounding the central star of IRAS 19312+1950 can be resolved from this study in some degree. Although evidence of an evolved star is still not completely crucial for this object or rather conflicting, we believe that this is a unique evolved object surrounded by a remnant of raw material. We present the observational results and discussion on the nature of this object in this paper.  %%%%%%%%%%%%%%%%%%%%%%%%% Fig. 1 %%%%%%%%%%%%%%%%% \\setcounter{figure}{0} \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){fig1.eps} %{fig0.eps} %%% \\FigureFile(width,height){filename} \\end{center} \\caption{$JHK$ false-color composite image of the field of IRAS 19312+1950, using the UH 88-inch telescope with the SIRIUS near-IR array camera (by courtesy of M. Tamura and K. Murakawa, National Astronomical Observatory). The image size is $128''\\times 128''$. The north is up, and the east is left.}\\label{fig1} \\end{figure}   % \\newpage ",
        "conclusions": "We have detected a variety of molecular lines of O-, C- and N-bearing molecules toward IRAS 19312+1950, a bipolar nebula. The line profiles consist of narrow and broad components of the widths of $\\sim$4 and $\\sim$30 km s$^{-1}$, respectively. From the parabolic line shape, the broad component is inferred to originate from the circumstellar envelope of this object. The LVG calculations suggest the mass loss rate of about $2.6 \\times 10^{-4} M_{\\odot}$ yr$^{-1}$ for this object. However, the origin of the narrow component is still a puzzle; from the low temperature and variety of molecules detected, it is suggested to be a dark cloudlet with a mass of about 20 $M_{\\odot}$. If it were a trapped material by a hypothetical companion, the H$_2$/CO mass ratio must be a factor of 10 less than the value normally assumed.  If this is really a dark cloud with a mass more than 20 $M_{\\odot}$, how the AGB/post-AGB star can accommodate a  cloud more massive than the star mass ? Furthermore, if this is a low-mass young stellar object with SiO masers more than a few kpc away, it is hard to explain the apparent size of the bright nebulosity and luminosity of the central star.  It is hard to provide a reasonable interpretation for this object except considering somewhat exotic, previously unknown, star model as a hypergiant (\\cite{jur01}), or possibly an AGB/post-AGB object with the huge orbiting gas reservoir in which the H$_2$/CO ratio is small compared with a normal ratio. Another extreme is a low-mass young stellar object with SiO masers with yet undeveloped CO outflow but shining the surrounding dust of the size of 30$''$.  Neither of them could be acceptable, except with excuse if they were extreme objects.  \\vspace{5mm}  The authors thank Dr. M. Tamura, and K. Murakawa for providing them with the SIRIUS image of IRAS 19312+1950. They also thank Dr. T. Umemoto for useful suggestions, and students (H. Fukushi, and Y. Tamura) for help of observations.   %%%%%"
    },
    "0410/astro-ph0410470_arXiv.txt": {
        "abstract": "Even if dark matter particles are unambiguously discovered in experiments, there is no clear reason to expect that the dark matter problem has been solved. It is very easy to provide examples of dark matter scenarios ({\\it e.g.} in supersymmetric models) where nearly identical detector signals correspond to extremely different relic densities. Therefore, the density of the particles discovered must be determined before their cosmological relevance is established. In this talk, I will present a general method to estimate the local density of dark matter particles using both dark matter and hadron collider experimental data when it becomes available. These results were obtained in collaboration with Gordon Kane at the University of Michigan. ",
        "introduction": "\\vspace{-0.6cm}\\indent We are now confident that our universe contains a large amount of cold dark matter. The most popular particle candidates for dark matter are weakly interacting massive particles (wimps). These particles are being searched for directly and indirectly by dozens of experimental groups throughout the world. If we are fortunate, wimps may soon be discovered experimentally.\\\\ \\indent Although the discovery of wimps in the galactic halo would have enormous implications for our understanding of elementary particle physics, it implies very little about our understanding of dark matter. It it not reasonable to assume that the particle discovered represents all of the dark matter. This issue was first raised technically, though not resolved, in \\cite{Brhlik:2000dm}. Indeed, it is easy to give examples where particles discovered in future and current experiments consist of less than $1\\%$ of the total dark matter.\\\\ \\indent Even if weakly interacting massive particles are produced at colliders, it is still very important to directly determine their local density in the galactic halo. This can only be done with experimental data. During this talk, I will describe how this can be determined in the context of supersymmetry where the dark matter is the lightest supersymmetric particle (LSP). Similar analyses could be done for any dark matter candidate. \\\\ \\indent This talk is organized as follows. First, I will illustrate why a discovery of dark matter particles in the halo is insufficient to address the dark matter problem and describe some of the uncertainties of relating the local and relic densities of dark matter. I will then describe how dark matter is detected directly in experiments and offer the general form of the interaction rate. This will show what is required to determine the local density of wimps. I will give a very useful way to improve these expressions using data from different detector materials and at different recoil energies.\\\\ \\indent In order to deduce the local density of wimps, it is essential to know the wimp's mass. I will present two methods to determine the mass of a wimp using direct detection data alone. The first reviews a well known relationship and the second presents preliminary results of the author. All of this work is then combined in the framework of the most general minimally supersymmetric standard model where the neutralino is the wimp seen in dark matter experiments. A general procedure is presented to estimate the local density and explicit bounds are given. ",
        "conclusions": "\\indent We have seen that, by itself, a discovery of dark matter particles in our galactic halo cannot address the dark matter problem of the universe. However, combined with data from colliders to identify the particle and limit its interaction parameters, we can generally estimate its local density.\\\\ \\indent In the framework of supersymmetry, we have shown a robust and iteratively improvable method to estimate the local density of a neutralino LSP observed in direct detection experiments using any information available about the MSSM.\\\\ \\indent Therefore, although the dark matter problem may not be solved immediately when wimps are observed, there are clear and general ways to address their cosmological significance.\\\\  \\noindent{\\bf Acknowledgements}\\\\ \\indent This research was done in collaboration with Gordon Kane of the University of Michigan and was supported by the Michigan Center for Theoretical Physics and the National Science Foundation's 2004 REU program.\\\\ \\indent I would like to thank the organizers of the 42$^{\\mathrm{nd}}$ {\\it International School of Subnuclear Physics}, Antonino Zichichi and Gerardus 't Hooft, for the opportunity to present this work and for organizing such a wonderful school."
    },
    "0410/gr-qc0410092_arXiv.txt": {
        "abstract": "Improved Wentzel-Kramers-Brillouin (WKB)-type approximations are presented in order to study cosmological perturbations beyond the lowest order. Our methods are based on functions which approximate the true perturbation modes over the complete range of the independent (Langer) variable, from sub-horizon to super-horizon scales, and include the region near the turning point. We employ both a perturbative Green's function technique and an adiabatic (or ``semiclassical'') expansion (for a linear turning point) in order to compute higher order corrections. Improved general expressions for the WKB scalar and tensor power spectra are derived for both techniques. We test our methods on the benchmark of power-law inflation, which allows comparison with exact expressions for the perturbations, and find that the next-to-leading order adiabatic expansion yields the amplitude of the power spectra with excellent accuracy, whereas the next-to-leading order with the perturbative Green's function method does not improve the leading order result significantly. However, in more general cases, either or both methods may be useful. ",
        "introduction": "\\label{intro} Anisotropies in the cosmic microwave background (CMB) radiation and inhomogeneities in the large scale structures of the Universe have nowadays become a fundamental tool to study the early universe. This trend will persist in the next years due to the data releases of Wilkinson~Microwave~Anisotropy~Probe~\\cite{wmap}, Sloan~Digital~Sky~Survey~\\cite{sdss} and the launch of the Planck mission~\\cite{planck}. \\par In a primordial inflationary phase, matter and space-time fluctuations are generated which inherit characteristics from the particular inflationary model studied (for an introduction and review see, e.g.~Ref.~\\cite{infla}). It is therefore very important to be able to determine the power spectra of cosmological perturbations for a large variety of different inflationary models, so that comparison with available and future data will tell us which models satisfactorily represent the time evolution of the Universe. Since exact solutions are not available for cosmological perturbations in general inflationary models (except for the case of an exponential potential~\\cite{lythstewart}) approximation methods are very welcome. The method of approximation most used is the slow-rollover approximation, introduced by Stewart and Lyth~\\cite{SL} (see Ref.~\\cite{LLKCTBA} for a review), which was inspired by power-law inflation. \\par It is also worthwhile testing different approximation schemes. In this paper we present Wentzel-Kramers-Brillouin (WKB)-type approaches which improve the proposals of Refs.~\\cite{martin_schwarz,H_MP} (see also Refs.~\\cite{mukh,mukh_thory,wang_mukh,SGong,TW,FMVV_II} for other approximations). For power-law inflation, the usual WKB method to lowest order reproduces exactly the spectral indices of perturbations, but the amplitudes are still rendered poorly. Given the importance of the consistency equation between the ratio of tensor to scalar amplitudes and the spectral index of tensor modes for single Klein-Gordon scalar field inflation~\\cite{infla}, we believe that the prediction on the amplitudes should also be improved, as is done in Ref.~\\cite{HHHJM} by means of the so-called ``uniform approximation'' to next-to-leading order (for the details on the formalism, see Ref.~\\cite{olver}). \\par Let us begin by recalling the Robertson-Walker metric in conformal time $\\eta$, \\be \\d s^2=a^2(\\eta)\\,\\left[-\\d\\eta^2+\\d\\vec x\\cdot \\d\\vec x\\right] \\ , \\ee where $a$ is the scale factor of the Universe. Scalar (density) and tensor (gravitational wave) fluctuations are determined on this background by two functions, respectively denoted as $\\mu=\\mu_{\\rm S}:= a\\,Q$ ($Q$ is the Mukhanov variable~\\cite{mukh}) and $\\mu=\\mu_{\\rm T}:= a\\,h$ ($h$ is the amplitude of the two polarizations of gravitational waves)~\\cite{gris,staro}, which satisfy a one-dimensional Schr\\\"odinger-like equation \\be \\left[\\frac{\\d^2}{\\d\\eta^2} +k^2-U(\\eta)\\right]\\,\\mu = \\left[\\frac{{\\d}^2}{{\\d}\\eta^2}+\\Omega^2(k,\\eta)\\right] \\,\\mu=0 \\ , \\label{osci} \\ee where $k$ is the wave-number. The time-dependent frequency is given by \\be \\Omega^2(k,\\eta)=k^2-\\frac{z''}{z} \\ , \\label{freq} \\ee where $z=z_{\\rm S}:=a\\,\\phi'/{\\cal H}$ for scalar ($\\phi$ is the homogenous inflaton) and $z=z_{\\rm T}:=a$ for tensor perturbations, primes denote derivatives with respect to $\\eta$, and ${\\cal H}$ ($:= a'/a$) is the conformal Hubble parameter. Eq.~(\\ref{osci}) must be solved together with the condition that the modes are initially plane waves for wavelengths much shorter than the Hubble radius, \\be \\lim_{\\frac{k}{a\\,H}\\rightarrow +\\infty} \\mu(k,\\eta) =d\\,\\frac{{\\rm e}^{-i\\,k\\,\\eta}}{\\sqrt{2\\,k}} \\ , \\label{init_cond_on_mu} \\ee where $d=d_{\\rm S}:=1$ and $d=d_{\\rm T}:=\\sqrt{8\\,\\pi}/m_{\\rm Pl}$ for scalar and tensor perturbations ($m_{\\rm Pl}$ is the Planck mass). The dimensionless power spectra of scalar and tensor fluctuations are then given by \\begin{subequations} \\be \\mathcal{P}_{\\zeta}= \\displaystyle\\frac{k^{3}}{2\\,\\pi^{2}}\\, \\left|\\frac{\\mu_{\\rm S}}{z_{\\rm S}}\\right|^{2} \\ , \\ \\ \\mathcal{P}_{h}= \\displaystyle\\frac{4\\,k^{3}}{\\pi^{2}}\\, \\left|\\frac{\\mu_{\\rm T}}{z_{\\rm T}}\\right|^{2} \\ . \\label{spectra_def} \\ee The spectral indices and their runnings are defined at an arbitrary pivot scale $k_*$ as \\be n_{\\rm S}-1:= \\left.\\displaystyle\\frac{\\d\\ln \\mathcal{P}_{\\zeta}} {\\d\\ln k}\\right|_{k=k_{*}} \\ , \\ \\ n_{\\rm T}:= \\left.\\displaystyle\\frac{\\d\\ln \\mathcal{P}_{h}} {\\d\\ln k}\\right|_{k=k_{*}} \\ , \\label{n_def} \\ee and \\be \\alpha_{\\rm S}:=\\left. \\frac{\\d^{2}\\ln\\mathcal{P}_{\\zeta}} {(\\d\\ln k)^{2}}\\right|_{k=k_{*}} \\ , \\ \\ \\alpha_{\\rm T}:=\\left. \\frac{\\d^{2}\\ln \\mathcal{P}_{h}} {(\\d\\ln k)^{2}}\\right|_{k=k_{*}} \\ . \\label{alpha_def} \\ee \\end{subequations} \\par The evolution of the Universe is usually described by means of a set of flow equations~\\cite{martin_schwarz,terrero,tesi}. The zero horizon flow function is defined by \\be \\epsilon_{0}:=\\frac{H(N_{i})}{H(N)} \\ , \\label{eps_0} \\ee where $H:=\\dot{a}/a$ is the Hubble rate and dots denote derivatives with respect to cosmic time $\\d t=a(\\eta)\\,\\d\\eta$. In this expression, $N$ is the number of e-folds, $N:=\\ln (a/a_i)$ [where $a_i=a(\\eta_i)$] after the arbitrary initial time $t_i=t_i(\\eta_i)$. The hierarchy of horizon flow functions is then defined according to \\be \\epsilon_{n+1}:= \\frac{\\d\\ln |\\epsilon_{n}|}{\\d N} \\ , \\quad\\quad n\\geq 0 \\ , \\label{hor_flo_fun} \\ee and we recall that inflation takes place for $\\epsilon_1 < 1$. \\par The effective potentials $U=U_{\\rm S}$ and $U=U_{\\rm T}$ can now be expressed in terms of $\\epsilon_1$, $\\epsilon_2$, and $\\epsilon_3$ only as \\be && \\frac{U_{\\rm T}(\\eta)}{a^{2}\\,H^{2}}=2-\\epsilon_{1} \\nonumber \\\\ \\label{pot} \\\\ && \\frac{U_{\\rm S}(\\eta)}{a^{2}\\,H^{2}} = 2-\\epsilon_{1}+\\frac{3}{2}\\,\\epsilon_{2} -\\frac{1}{2}\\,\\epsilon_{1}\\,\\epsilon_{2} +\\frac{1}{4}\\,\\epsilon_{2}^{2} +\\frac{1}{2}\\,\\epsilon_{2}\\,\\epsilon_{3} \\ , \\nonumber \\ee and Eq.~(\\ref{osci}) can be solved exactly for $\\epsilon_1$ constant in time, which is the case of an exponential potential for the inflaton. When the $\\epsilon_n$ have an arbitrary time dependence, approximate methods are necessary. \\par A WKB analysis of this problem has been recently presented in Ref.~\\cite{martin_schwarz} to first order in the adiabatic (or ``semiclassical'') expansion (see next~Section and Ref.~\\cite{birrel} for precise definitions), in which the Langer transformation~\\cite{langer} \\be x:=\\ln\\left(\\frac{k}{H\\,a}\\right) \\ ,\\quad \\chi:=(1-\\epsilon_1)^{1/2}\\,{\\rm e}^{-x/2}\\,\\mu \\ , \\label{transf} \\ee introduced in Ref.~\\cite{wang_mukh} is used~\\footnote{Note that we define $x$ with the opposite sign with respect to that in Ref.~\\cite{martin_schwarz}. This of course implies that our $x_i\\to+\\infty$ and $x_f\\to-\\infty$, whereas $x_i\\to-\\infty$ and $x_f\\to+\\infty$ in Ref.~\\cite{martin_schwarz}.} in order to improve the accuracy. Such a transformation brings Eq.~(\\ref{osci}) into the form \\be \\left[\\frac{{\\d}^2}{{\\d} x^2}+\\omega^2(x)\\right] \\,\\chi=0 \\ , \\label{new_eq} \\ee where the (new) frequency $\\omega(x)$ in general vanishes at the classical ``turning point'' $x=x_*$. On sub-horizon scales ($x\\gg x_*$), the mode $\\chi$ oscillates ($\\omega\\simeq k$) and the standard WKB approximation can be applied without difficulty. On super-horizon scales ($x\\ll x_*$), the perturbations exponentially decay or grow and the standard WKB approximation can again be applied. One then matches the two approximate solutions at $x=x_*$ to obtain the super-horizon amplitudes at $x_f\\ll x_*$ which determine the CMB spectra. This procedure opens up the possibility of further improving the knowledge of the spectra by including subsequent adiabatic orders. \\par However, a straightforward generalization to higher adiabatic orders is seriously hindered by the lack of a general prescription for matching the two WKB branches with sufficient accuracy. This problem was thoroughly discussed in Refs.~\\cite{langer34,langer}, where it was also suggested to replace the standard (plane wave-like) WKB functions with Bessel's functions. The latter indeed remain good approximations at the turning point, whereas the former require the matching with yet another particular solution at some (unspecified) points both to the left and right of $x=x_*$. We shall follow the proposal of Ref.~\\cite{langer} for a linear turning point and show that the inclusion of higher adiabatic orders indeed improve the results, contrarily to what was found in Ref.~\\cite{hunt}. We shall also introduce a new (perturbative) expansion for general turning point, which makes use of the Green's function technique, and compare the corresponding corrections for the case of power law inflation. \\par The article is organized as follows: In Section~\\ref{WKB_power_spectra}, we shall review the application to inflationary cosmological perturbations of the standard WKB approximation to leading adiabatic order~\\cite{martin_schwarz} and its shortcomings. In Section~\\ref{schiffmethod}, we shall describe in full generality a method of improving the results to higher orders, both by the perturbative expansion discussed in detail in Section~\\ref{pert} and the adiabatic expansion reviewed in Section~\\ref{adiab}. In Section~\\ref{Power-law inflation}, we shall apply the improved WKB approximation to power-law inflation which, being exactly solvable, will allow us to test the improved WKB predictions against exact solutions. Finally, we shall give our conclusions and observations in Section~\\ref{conc}. ",
        "conclusions": "\\setcounter{equation}{0} \\label{conc} In this paper we have improved the WKB approximation for the purpose of estimating the spectra of cosmological perturbations during inflation. We found general formulae for amplitudes, spectral indices and $\\alpha$-runnings of the fluctuations to next-to-leading order both in the adiabatic expansion of Ref.~\\cite{langer} and a new perturbative expansion which makes use of the Green's function technique. \\par We have then applied our method to power-law inflation in order to test it against exact results. It is known that the spectral indices and their runnings are obtained exactly in leading order, hence we focussed on the spectra to next-to-leading order. We have found that the perturbative corrections remain too small to yield any significant improvement, whereas the adiabatic expansion to next-to-leading order reproduces the exact amplitudes with great accuracy (see Ref.~\\cite{HHHJM} for a similar result with the uniform approximation). This result does not however mean that the perturbative expansion will not lead to significant corrections in different inflationary scenarios. \\par One way of understanding the difference between the two expansions in the power-law case may be the following. It is known that the Born~approximation in quantum mechanics (analogous to our perturbative method with the Green's function) is good for high multipoles (angular momenta) $\\ell$ of the expansion in spherical harmonics. In our approach, the parameter $\\beta$ plays the role of the multipole index $\\ell$ for the energy levels of the hydrogen atom, as can be seen on comparing the frequency~(\\ref{omega2_PL}) with the expression given in Ref.~\\cite{langer} (see also Ref.~\\cite{martin_schwarz}). The approximation with the Green's function is therefore expected to yield more significant corrections for large values of $|\\beta|$, which coincide with the regime of fast-roll. Indeed, the leading order becomes more and more accurate for increasing $|\\beta|$. Since the interesting regime for inflation involves small values of $\\beta\\sim -2$, it is instead the adiabatic approximation which seems better since the horizon flow functions evolve slowly and/or the states are quasi-classical~\\footnote{Cosmological perturbations are amplified by inflation evolving from vacuum to highly squeezed states, which resemble classical states in the amplitude of fluctuations~\\cite{polarski}.}. This may be the reason whereby the adiabatic approximation works better in the next-to-leading order. However, in more general cases, either or both methods may contribute significant corrections. In this respect, let us remark that the adiabatic expansion can be straightforwardly applied only for a linear turning point~\\footnote{The same limitation seems to affect the method used in Ref.~\\cite{HHHJM}.}, whereas the perturbative Green's function method is not so restricted. \\par Having assessed the accuracy of the method, it is now natural to use it in order to improve current estimates for inflationary models which are not exactly solvable. We are in fact analyzing the slow roll approximation with our method and will report about it in a future publication~\\cite{new}."
    },
    "0410/astro-ph0410700_arXiv.txt": {
        "abstract": "Half a year after its outburst, the nova V4743\\,Sgr evolved into the brightest super-soft X-ray source in the sky with a flux maximum around 30\\,\\AA, exhibiting resonance lines of C\\,{\\sc v}, C\\,{\\sc vi}, N\\,{\\sc vi}, N\\,{\\sc vii}, and O\\,{\\sc vii}. We present preliminary results of an analysis of the  XMM-Newton RGS spectra by means of NLTE model-atmosphere techniques. ",
        "introduction": "Classical novae occur in close binary systems (main-sequence star + white dwarf). V4743\\,Sgr has been discovered as a nova on Sep 20, 2002 (Haseda et al\\@. 2002). A first CHANDRA observation, taken with ACIS-S in Dec 2002 (Ness et al\\@. 2003), shows a X-ray spectrum which is consistent with a view on the optically thick atmosphere of an extremely hot white dwarf with strong C\\,{\\sc v}, N\\,{\\sc vi}, and N\\,{\\sc vii} absorption features (indicating $T_\\mathrm{eff}$ of 1 -- 2 MK, Ness et al\\@. 2003) which are blue shifted (\\mbox{-2400}\\,km/sec) --- most likely expanding gas which consists out of CNO-processed material.  \\begin{figure}[ht] \\plotone{rauch2_fig1.eps} \\caption{Synthetic fluxes calculated from line-blanketed NLTE model atmospheres including H, He, C, N, and O. The abundances are mass ratios.} \\label{synflux} \\end{figure}  \\begin{figure}[hb] \\plotone{rauch2_fig2.eps} \\caption{Comparison of fluxes of $T_\\mathrm{eff}$ = 700\\,kK models with different abundance ratios (thick: like given in Fig.\\,\\ref{synflux}, thin: solar, dashed: black body).} \\label{comflux} \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410066_arXiv.txt": {
        "abstract": "{We present astrometry and $BVI$ photometry, down to $V\\simeq22$, of the very young open cluster \\ngc6530,  obtained from observations taken with the Wide Field Imager camera at the MPG/ESO 2.2\\,m Telescope. Both the $V$ vs. $B-V$ and the $V$ vs. $V-I$ color-magnitude diagrams (CMD) show the upper main sequence dominated by very bright cluster stars, while,  due to the high obscuration  of the giant molecular cloud surrounding the cluster, the blue envelopes of the diagrams at $V\\gtrsim 14$ are limited to the main sequence stars at the distance of \\ngc6530. This  particular structure of the \\ngc6530 CMD allows us to conclude that its distance is about $d \\simeq 1250$\\,pc, significantly lower than the previous determination of d=1800\\,pc.\\\\ We have positionally matched our optical catalog with the list of X-ray sources found in a Chandra-ACIS observation, finding a total of 828 common stars, 90\\% of which are pre-main sequence stars in \\ngc6530. Using evolutionary tracks of \\citet{sies00}, mass and age values  are inferred for these stars. The median age of the cluster is about 2.3\\,Myr; in the mass range (0.6--4.0)$\\,M_\\odot$, the Initial Mass Function (IMF) shows a power law index $x=1.22\\pm0.17$, consistent with both the Salpeter index (1.35), and with the index derived for   other young clusters ; towards smaller masses the IMF shows a peak and then it starts to decrease. ",
        "introduction": "} % Star formation regions and very young open  clusters are crucial systems for  understanding   the star formation process because they allow us to derive Initial Mass Functions that are  not affected by stellar and/or dynamical evolution effects. \\ngc6530 ($l=6.14, b=-1.38$) is an example of a very young open cluster  located in front of the M8 giant molecular cloud, also referred to as Lagoon Nebula \\citep{lada76} . The brightest part of M8 is the Hourglass Nebula, illuminated by the extremely young O star Herschel 36 (O7 V). Other stars of spectral type O, 9 Sgr (O4 V) and the binary  system HD 165052 (O6.5 V + O6.5 V) excite the Lagoon Nebula. The open cluster \\ngc6530 is located about   9\\Min8~ eastwards  of the Hourglass Nebula (corresponding to a projected distance from the Nebula of 3.6 pc, using the cluster distance $d=1250$\\,pc  estimated in this work) and can be recognized by its  brightest stars.  Since the work of \\citet{walk57}, several investigations have been devoted to study this cluster and to estimate its parameters. Using mainly photoelectric and photographic  observations, many attempts have been done to estimate the distance of \\ngc6530 from the  CMDs, which show a normal cluster main sequence down to about A0 stars; fainter stars lie above the main sequence,   indicating that the cluster is so young that its low mass members are still gravitationally contracting \\citep{walk57}. The lack of the complete cluster main sequence made it hard to obtain reliable estimates of the cluster distance. \\ngc6530 distance has been estimated by several authors to be in the range (1300--2000)\\,pc (cfr. Table \\ref{distance_tab}), by taking advantage of spectroscopic observations, to determine spectral types, and proper motions, to select cluster members. \\begin{table}[t] \\centering \\caption {Literature distance values   for \\ngc6530} \\vspace{0.5cm} \\begin{tabular}{ccc} \\hline \\hline Distance (pc)& $({\\rm m-M})_0$\t&Reference \\\\ \\hline 1380--2000\t&10.7--11.5&\\citet{walk57}\\\\%&photoelectric observations 1300\t\t&10.6\t   &\\citet{the60} \\\\%&photoelectric observations 1400\t\t&10.7\t   &\\citet{walk61}\\\\%&photoelectric observations 1580\t\t&11.0\t   &\\citet{hilt65}\\\\%&spectroscopy and photometry 1780\t\t&11.25\t   &\\citet{alte72}\\\\%&proper motions 1380\t\t&10.7\t   &\\citet{kila77}\\\\%&photographic photometry 1820\t\t&11.3\t   &\\citet{saga78}\\\\%&photoelectric observations 1900\t\t&11.4\t   &\\citet{chin81}\\\\%&$UBV$ photometry 1860\t\t&11.35\t   &\\citet{mcca90}\\\\%&photometry and spectroscopy 1600--2000\t&11.0--11.5&\\citet{anck97}\\\\%&multiwavelength study 1400\t        &10.74     &\\citet{lokt97}\\\\%&parallax 1800\t\t&11.25\t   &\\citet{sung00}\\\\%&$UBVRI$ and H$\\alpha$ Photometr 560--711\t&8.75--9.26&\\citet{lokt01}\\\\%&parallax \\hline \\hline \\end{tabular} \\label{distance_tab} \\end{table}  On the contrary, $UBV$ observations  have  allowed several authors to study the reddening law and to derive  cluster reddening values very similar to the average value $E(B-V)=0.35$  recently derived by \\citet{sung00}, who have assumed   a foreground reddening value, $E(B-V)_{fg}=0.17$, as estimated by \\citet{mcca90}. Finally, the age of the cluster  was  estimated to be in the range (1.5--2.0)\\,Myr by several authors \\citep{alte72,saga78,sung00}, while \\citet{dami04} give a median age of 0.8\\,Myr. Due to the young age of this cluster, low mass stars are expected to be still in pre-main sequence  phase, as already found in \\citet{sung00} by using photometric data down to the limiting magnitude $V\\sim 17$. 119 X-ray point sources in the Lagoon Nebula region have been recently detected by \\citet{rauw02} in  a 20ksec {\\it XMM--Newton} observation; they found that most of the X-ray sources are associated with pre-main sequence stars of low and intermediate mass. However, a larger list of point sources in the same region, with a much better spatial resolution, has been recently obtained  by \\citet{dami04}, using Chandra ACIS-I X-ray data. Damiani et al.  show that their source sample is made for at least $90\\%$ by cluster members and that out of the 884 X-ray sources found, only 220 have a counterpart in the \\citet{sung00} optical catalog and most of   this latter are pre-main sequence members of \\ngc6530. As suggested by Damiani et al., the remaining X-ray sources are likely stars with magnitude fainter than $V\\sim17$. The lack of published photometric data at magnitudes fainter than $V\\sim17$ has motivated the analysis of deep optical images  to study the low mass stellar population in the \\ngc6530 field and, in particular, the low-mass pre-main sequence stars of the cluster by means of cross-correlation with available X-ray data.  In this paper, we first present the observations and the data reduction procedure (Sect. 2) and the CMDs obtained using the optical catalog (Sect. 3). Therefore, we present  the cross-correlation of optical and X-ray data, from which membership is determined (Sect. 4), and \t     the cross-correlation of optical and 2MASS IR data (Sect. 5); mass and age determination of cluster members and the analysis of the spatial distribution are discussed in Sect. 6. Finally, the Luminosity and  the Initial Mass Functions  (Sect. 7) and our  conclusions (Sect. 8) are presented. ",
        "conclusions": ""
    },
    "0410/astro-ph0410250_arXiv.txt": {
        "abstract": "{ We report here the first evidence for stellar disk truncation at high redshift, based on surface photometry of a sample of 16 high redshift (0.6~$<$~z$~<$~1.0) disk galaxies from the GOODS HST/ACS data. The radial profiles are best fit by a double exponential profile. This result agrees with the profile of disks in local galaxies. The cosmological surface brightness dimming at this redshift range only allows us to detect galaxies with spatially ``early'' truncation, R$_{\\rm br}$/h$_{\\rm in}\\leqq$ 3.5. Six galaxies show the radial double exponential structure, with an average value of R$_{\\rm br}$/h$_{\\rm in}\\approx$~1.8. Such ``early'' truncated galaxies are missing in local samples so far. This result opens the ground for observing directly disk evolution through the study of the truncation radius as a function of redshift. ",
        "introduction": "The finite size of stellar disks was first reported by~\\cite{kruit79} who showed that disks were sharply truncated at around 4 scale-lengths. Recently, evidence has arisen that this truncation does not occur sharply but instead radial profiles are best fit by a double exponential shape \\citep{degrijs01,pohlen02}. For nearby objects, the profiles obtained from a sample of edge-on galaxies show a two-slope exponential with R$_{\\rm br}$/h$_{\\rm in}$ in the range $3-5$~\\citep{kregel}, where R$_{\\rm br}$, the break radius, is the radius where the slope changes and h$_{\\rm in}$ is the inner disk scale-length (see Pohlen et al., 2004; for a recent review on the truncation radius). This truncation might be related to a surface density threshold for star formation~\\citep{kennicutt} or might be related to the maximum specific angular momentum of the sphere from which the disk collapsed~\\citep{kruit87}. While galaxy formation models have been able to reproduce the observed double exponential profile~\\citep{bosch01,zhang}, none of these models have considered evolution of the truncation radius with redshift. At z~$\\approx$~1, galaxies are half the age of local objects and one can probe into a period where evolutionary effects may be significant ~\\citep{mao}. A study of the evolution of the truncation radius can help to discern between truncation arising as a relic from galaxy formation or due to a star formation threshold. Such a project has been carried out as a pilot survey using HST archival data and is reported in this Letter. ",
        "conclusions": ""
    },
    "0410/astro-ph0410299_arXiv.txt": {
        "abstract": "{ As the closest radio galaxy, Centaurus A is a powerful laboratory for the X-ray study of radio-emitting structures and their interactions with the hot interstellar medium (ISM).  This paper details our interpretation of the remarkable X-ray enhancement which caps the inner southwest radio lobe, at a radius of about 6~kpc from the galaxy center.  The shell of X-ray-emitting gas is hotter than the ambient ISM, and overpressured by a factor of 100.  We argue that it is heated compressed material behind the supersonically-advancing bow shock of the radio lobe, the first example of the phenomenon to be clearly detected.  The results demonstrate that Cen A is actively re-heating nearby X-ray-emitting gas.  The shell's kinetic energy is $\\sim 5$ times its thermal energy, and exceeds the thermal energy of the ISM within 15~kpc of the center of the galaxy.  As the shell dissipates it will have a major effect on Cen~A's ISM, providing distributed heating.  } ",
        "introduction": "There is much current interest in the possibility that radio sources heat the interstellar and intercluster medium.  Such heating would help to explain the weakness or absence of lines from gas cooling below 1~keV in the densest central regions of galaxies and clusters \\citep[e.g.,][]{peter01}.  A good place for heating would be behind the bow shock of a supersonically expanding radio structure.  In this paper we outline the simple theory and its application to Cen~A, the first source to show clear and direct X-ray evidence of supersonic expansion and heating. ",
        "conclusions": "\\begin{itemize}  \\item A hot shell of X-ray-emitting gas caps the southwest radio lobe of Cen~A.  \\item Shock jump conditions for an advance speed of $\\sim 0.008$c (${\\cal M} = 8.5$) are satisfied if the shell is in pressure balance with unseen gas at $kT \\sim 6.8$~keV behind the bow shock.  The emissivity of the 6.8 keV gas is too low to separate its X-rays from those of the ten-times-brighter gas of the shell.  \\item The radiative timescale for material in the shell ($\\sim 2 \\times 10^9$ yrs, using equation 5.23 of \\citet{saraz86}) is large compared with the lobe expansion time ($< 2.4 \\times 10^6$ years), so the material in the shell behaves as an adiabatic gas \\citep[e.g.,][]{alex02}.  \\item The cooling has improved the shell--ISM contrast by an order of magnitude, an important factor in the detection and modelling of the shell.  \\item The lobe is (or was) powered by energy deposition from a jet, and is overpressured relative to the ambient ISM.  Monitoring the intensity and searching for proper motions in the knots of the weak counterjet on this side of the source would provide important diagnostics for establishing the current state of the jet within the southwest lobe.  \\item To order of magnitude, the mass in the hot shell is consistent with material swept up from the ISM.  There must have been sufficient time since any earlier epoch of lobe expansion in the region for the ISM to be replenished.  \\item The shell is overpressured compared with the minimum-energy pressure in the radio lobe (in magnetic field and radiating electrons) by a factor of $\\sim 10$.  There is no particular reason to think that a dynamical object, seen in a snapshot, should be in a state of minimum energy.  However, if we do assume minimum energy in the lobe, and that the shell has reached equilibrium [but note that the sound-crossing time in the shell (thickness $\\sim 0.3$ kpc, $c_{\\rm} \\sim  9 \\times 10^{-7}$ kpc yr$^{-1}$) is about 15 per cent of the maximum time we estimate it has taken the lobe to reach its current size], the shell's  overpressure relative to the radio lobe could be balanced by the ram pressure from internal motions in the lobe for a moderate relativistic proton loading.  \\item The shell's kinetic energy is $\\sim 5$ times its thermal energy, and exceeds the thermal energy of the ISM within 15 kpc of the centre of the galaxy.  As the shell dissipates, most of the kinetic energy should ultimately be converted into heat and this will have a major effect on Cen~A's ISM, providing distributed heating.  \\end{itemize}"
    },
    "0410/astro-ph0410120_arXiv.txt": {
        "abstract": "We present the results of a survey for OH maser emission at 1612 MHz from dust-enshrouded AGB stars and supergiants in the LMC and SMC, with the Parkes radio telescope, aimed at deriving the speed of the superwind from the double-peaked OH maser profiles. Out of 8 targets in the LMC we detected 5, of which 3 are new detections --- no maser emission was detected in the two SMC targets. We detected for the first time the redshifted components of the OH maser profile in the extreme red supergiant IRAS\\,04553$-$6825, confirming the suspicion that its wind speed had been severely underestimated. Despite a much improved spectrum for IRAS\\,04407$-$7000, which was known to exhibit a single-peaked OH maser, no complementary peak could be detected. The new detection in IRAS\\,05003$-$6712 was also single-peaked, but for two other new detections, IRAS\\,04498$-$6842 and IRAS\\,05558$-$7000, wind speeds could be determined from their double-peaked maser profiles. The complete sample of known OH/IR stars in the LMC is compared with a sample of OH/IR stars in the galactic centre. The LMC sources generally show a pronounced asymmetry between the bright blueshifted maser emission and weaker redshifted emission, which we attribute to the greater contribution of amplification of radiation coming directly from the star itself as the LMC sources are both more luminous and less dusty than their galactic centre counterparts. We confirm that the OH maser strength is a measure of the dust (rather than gas) mass-loss rate. At a given luminosity or pulsation period, the wind speed in LMC sources is lower than in galactic centre sources, and the observed trends confirm simple radiation-driven wind theory if the dust-to-gas ratio is approximately proportional to the metallicity. ",
        "introduction": "Asymptotic Giant Branch (AGB) stars lose $\\sim$35 to 85\\% of their mass before ending up as white dwarfs. Much of this mass loss takes place near the tip of the AGB during the superwind stage, at mass-loss rates of up to $\\dot{M}\\sim10^{-4}$ M$_\\odot$ yr$^{-1}$ (van Loon et al.\\ 1999, 2003) leading to the early cessation of their stellar evolution and resulting in chemical enrichment of the interstellar medium (ISM). The mass loss occurs in the form of a stellar wind reaching typical outflow velocities of 5 to 30 km s$^{-1}$. These winds are driven by radiation pressure on dust grains that form at a height of several stellar radii (Goldreich \\& Scoville 1976) in the presence of strong radial pulsation of the base of the stellar atmosphere. Likewise, more massive stars ($M_{\\rm ZAMS}\\gsim8$ M$_\\odot$) may pass through a red supergiant phase and lose mass in a similar manner.  \\begin{table*} \\caption[]{The target sample plus the detections from Wood et al.\\ (1992). Bolometric magnitudes $M_{\\rm bol}$, pulsation periods $P$, K-band pulsation amplitudes ${\\Delta}K$, K-band magnitudes and colours $J-K$ are from Whitelock et al.\\ (2003), {\\it IRAS} 25 $\\mu$m flux densities $F_{25}$ (not colour-corrected) are from Trams et al.\\ (1999), and spectral types are from van Loon et al.\\ (in preparation), except where noted otherwise. We adopt distances to the SMC and LMC of 50 and 60 kpc, respectively (distance moduli of 18.5 and 18.9 mag, respectively). Also listed are the previously detected OH masers, and whether we detected OH maser emission here (no entry means never tried).} \\begin{tabular}{llcccccclrr} \\hline\\hline Object           & Alternative      & $M_{\\rm bol}$    & $P$              & ${\\Delta}K$      & $K$              & $F_{25}$         & $J-K$            & Spectral         & Previous         & This work        \\\\ name             & name             & (mag)            & (mag)            & (days)           & (mag)            & (Jy)             & & type             & OH               & OH               \\\\ \\hline {\\it SMC}\\\\ IRAS\\,00483$-$7347 & LI-SMC\\,61         & $-$7.20\\rlap{$^8$} & 1200\\rlap{$^8$}    & 0.61\\rlap{$^8$}    & 9.10\\rlap{$^3$}    & 0.46\\rlap{$^3$}    & 3.01\\rlap{$^3$}    & late-M             & & no                 \\\\ IRAS\\,00591$-$7307 & HV\\,11417          & $-$8.30\\rlap{$^2$} & \\llap{$\\gsim$}1000\\rlap{$^1$} & \\llap{$>$}2.00     & 8.30\\rlap{$^3$}    & 0.10\\rlap{$^3$}    & 1.33\\rlap{$^3$}    & M5e I\\rlap{$^1$}   & & no                 \\\\ \\hline {\\it LMC}\\\\ IRAS\\,04407$-$7000 & LI-LMC\\,4          & $-$7.11            & 1199               & 1.23               & 8.79               & 0.76               & 2.34               & M7.5               & yes\\rlap{$^6$}     & yes                \\\\ IRAS\\,04498$-$6842 & LI-LMC\\,60         & $-$7.72            & 1292               & 1.30               & 8.08               & 0.89               & 1.86               & M10                & & yes                \\\\ IRAS\\,04509$-$6922 & LI-LMC\\,77         & $-$7.28            & 1292               & 1.45               & 8.59               & 0.86               & 2.21               & M10\\rlap{$^6$}     & no\\rlap{$^9$}      & no                 \\\\ IRAS\\,04516$-$6902 & LI-LMC\\,92         & $-$7.11            & 1091               & 1.41               & 8.72               & 0.55               & 2.32               & M9\\rlap{$^6$}      & & no                 \\\\ IRAS\\,04553$-$6825 & WOH\\,G064          & $-$9.19\\rlap{$^5$} & \\hspace*{1mm}841   & 0.34               & 7.09               & \\llap{1}3.53       & 2.60\\rlap{$^4$}    & M7.5               & yes\\rlap{$^9$}     & yes                \\\\ IRAS\\,05003$-$6712 & LI-LMC\\,0297       & $-$6.20            & \\hspace*{1mm}883   & 1.59               & 9.95               & 0.33               & 2.95               & M9                 & & yes                \\\\ IRAS\\,05294$-$7104 & LI-LMC\\,1153       & $-$6.79            & 1079               & 1.20               & 9.21               & 0.56               & 2.97               & M8                 & & no                 \\\\ IRAS\\,05558$-$7000 & LI-LMC\\,1790       & $-$6.97            & 1220               & 1.42               & 9.25               & 0.80               & 3.27               & & & yes                \\\\ \\hline \\multicolumn{5}{l}{\\it Sample from Wood et al.\\ (1992) --- excluding IRAS\\,04553$-$6825} \\\\ IRAS\\,04545$-$7000 & LI-LMC\\,159        & $-$6.56            & 1216               & 1.57               & \\llap{1}0.13       & 0.83               & 5.70\\rlap{$^9$}    & & yes\\rlap{$^9$}     & \\\\ IRAS\\,05280$-$6910 & NGC\\,1984-IRS1     & $-$7.75\\rlap{$^9$} & & & 8.19\\rlap{$^9$}    & \\llap{2}4.18\\rlap{$^7$} & & & yes\\rlap{$^9$}     & \\\\ IRAS\\,05298$-$6957 & LI-LMC\\,1164       & $-$6.72\\rlap{$^9$} & 1280\\rlap{$^9$}    & 2.00\\rlap{$^9$}    & \\llap{1}0.29\\rlap{$^9$} & 1.38               & 3.54\\rlap{$^9$}    & & yes\\rlap{$^9$}     & \\\\ IRAS\\,05329$-$6708 & LI-LMC\\,1286       & $-$6.95            & 1262               & 1.50               & 9.90               & 1.23               & 5.10\\rlap{$^9$}    & & yes\\rlap{$^9$}     & \\\\ IRAS\\,05402$-$6956 & & $-$6.77            & 1393               & 1.80               & \\llap{1}0.40       & 1.02               & 4.46\\rlap{$^9$}    & & yes\\rlap{$^9$}     & \\\\ \\hline \\end{tabular}\\\\ References: $^1$ Elias, Frogel \\& Humphreys (1980), $^2$ Elias, Frogel \\& Humphreys (1985), $^3$ Groenewegen \\& Blommaert (1998), $^4$ Trams et al.\\ (1999), $^5$ van Loon et al.\\ (1999), $^6$ van Loon et al.\\ (1998a), $^7$ van Loon et al.\\ (2001b), $^8$ Whitelock et al.\\ (1989), $^9$ Wood et al.\\ (1992). \\end{table*}  During the superwind stage the circumstellar dust envelope becomes opaque at visual wavelengths, but the absorbed stellar light is re-radiated at infrared (IR) wavelengths. The resulting spectral energy distribution may be modelled in order to derive the mass-loss rate, provided we know the dust-to-gas ratio, dust grain properties, stellar luminosity and effective temperature, and outflow kinematics. Data on Magellanic Cloud and galactic dust-enshrouded AGB stars are consistent with a dust-to-gas ratio which depends linearly on metallicity, without a strong metallicity dependence of the total (gas+dust) mass-loss rate (van Loon 2000). This relies on the validity of simple dust-driven wind theory, in which the wind speed, $v_{\\rm exp}$, depends on the dust-to-gas ratio, $\\psi$, and luminosity, $L$, as (van Loon 2000; cf.\\ Habing, Tignon \\& Tielens 1994; Elitzur \\& Ivezi\\'{c} 2001): \\begin{equation} v_{\\rm exp}\\propto\\psi^{1/2}L^{1/4}. \\end{equation} It is thus of importance, both for deriving accurate mass-loss rates as well as for understanding the mass-loss mechanism, to measure wind speeds for stars of different luminosity and metallicity.  The wind speed may be measured from molecular emission line profiles at radio and (sub)mm wavelengths. This is best done for the abundant CO molecule, but its lines are too weak to detect outside the local galactic disk with presently available observatories. Oxygen-rich dust-enshrouded AGB stars (invariably of M spectral type) may also exhibit intense maser emission from the stimulated re-emission of stellar and circumstellar radiation by SiO, H$_2$O and/or OH molecules. These masers are radially stratified due to the spatial separation of the molecules in the outflowing circumstellar envelope (see Habing (1996) for a review). SiO is easily condensed into dust grains and therefore is abundant only in the dust-free inner cavity close to the star. Its maser emission is tangentially amplified resulting in a single peak which can be used to determine the stellar velocity. H$_2$O masers originate in the dust formation zone. In Miras these are predominantly amplified tangentially, providing another means of measuring the stellar velocity, but in the more extreme OH/IR stars also radial amplification is effective, resulting in a double-peaked line profile from which the outflow velocity can be measured (Takaba et al.\\ 1994). OH is formed by the dissociation of H$_2$O by interstellar UV radiation (Goldreich \\& Scoville 1976). As the OH maser emission is radially amplified it provides a means for measuring the speed of the fully developed wind (Elitzur, Goldreich \\& Scoville 1976) --- as opposed to the H$_2$O masers which probe a region where the wind has not yet reached its final speed (Richards, Yates \\& Cohen 1998, 1999).  OH masers have been widely detected in galactic metal-rich samples. The best opportunities to investigate the wind speed in a low-metallicity environment are provided by the Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC), with metallicities of the intermediate-age stellar populations of $\\sim\\frac{1}{3}$ and $\\sim\\frac{1}{5}$ solar respectively. The Magellanic Clouds have the additional advantage that their distances are reasonably well known, and hence the dependence of the wind speed on stellar luminosity can be determined. Prior to this study, only seven identified OH/IR stars were known in the LMC, of which five are AGB stars (the other two are supergiants), and none in the SMC. Amongst the five AGB stars, one was detected as a single-peaked OH maser line, and only one showed an uncomplicated ``classical'' double-peaked profile yielding a reliable value for its wind speed. We here present new detections of magellanic OH/IR stars, and an analysis of the dependence of the superwind speed on stellar luminosity and metallicity.  \\begin{table*} \\caption[]{OH maser detections, with the heliocentric velocities of the blue and red peaks (where identified), corresponding wind speed $v_{\\rm exp}$ from the separation of the peaks, peak flux density and integrated flux $F_{\\rm int}$. The rms noise levels in the flux densities and in the integrated fluxes are also listed as $\\sigma$ and $\\sigma_{\\rm int}$, respectively (for IRAS\\,04407$-$7000, IRAS\\,04498$-$6842, IRAS\\,05003$-$6712 and IRAS\\,05558$-$7000 these values refer to the smoothed spectra shown in Figs.\\ 1--4, with an effective channel width of 0.36 km s$^{-1}$).} \\begin{tabular}{lclrclrccc} \\hline\\hline Object                                & Integration                           & \\multicolumn{2}{c}{Velocity at peaks} & $v_{\\rm exp}$                         & \\multicolumn{2}{c}{Peak flux density} & $\\sigma$                              & $F_{\\rm int}$                         & $\\sigma_{\\rm int}$                    \\\\ name                                  & time (s)                              & \\multicolumn{2}{c}{(km s$^{-1}$)}     & (km s$^{-1}$)                         & \\multicolumn{2}{c}{(mJy)}             & (mJy)                                 & (mJy km s$^{-1}$)                     & (mJy km s$^{-1}$)                     \\\\ & & Blue                                  & Red                                   & & Blue                                  & Red                                   & & & \\\\ \\hline {\\it LMC}\\\\ IRAS\\,04407$-$7000 & 33000              & 240                & 270\\rlap{?}        & \\llap{1}5\\rlap{?}  & 50                & 8\\rlap{?}        & 5\\rlap{.5}       & \\llap{1}15         & 10                 \\\\ IRAS\\,04498$-$6842 & 28800              & 246                & 272                & \\llap{1}3          & 23                & 11                & 5\\rlap{.9}       & 34                 & 12                 \\\\ IRAS\\,04553$-$6825 & 14455              & 252, 262           & 300, 280           & \\llap{2}4          & \\llap{6}00, 180    & 70,  50           & \\llap{1}0\\rlap{.5} & \\llap{2,9}50       & 13                 \\\\ IRAS\\,05003$-$6712 & 32010              & 268                & & & 33                & $\\sim$13\\rlap{?}   & 5\\rlap{.1}       & 38                 & 12                 \\\\ IRAS\\,05558$-$7000 & 32500              & 265                & 285                & \\llap{1}0          & 17                & 17                & 6\\rlap{.9}       & 50                 & 14                 \\\\ \\hline \\multicolumn{6}{l}{\\it Sample from Wood et al.\\ (1992) --- excluding IRAS\\,04553$-$6825} \\\\ IRAS\\,04545$-$7000 & & 258                & 274                & 8                  & \\llap{1}40         & 30                & \\llap{1}7          & & \\\\ IRAS\\,05280$-$6910 & & 255                & 289                & \\llap{1}7          & 90                & 70                & \\llap{1}7          & & \\\\ IRAS\\,05298$-$6957 & & 271                & 292                & \\llap{1}0\\rlap{.5} & \\llap{2}40         & 130                & \\llap{1}7          & & \\\\ IRAS\\,05329$-$6708 & & 301                & 323                & \\llap{1}1          & 70                & 130                & \\llap{1}7          & & \\\\ IRAS\\,05402$-$6956 & & 263                & 284                & \\llap{1}0\\rlap{.5} & 80                & 60                & \\llap{1}7          & & \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We have presented the results of a new survey for OH maser emission at 1612 MHz in the circumstellar environments of magellanic dust-enshrouded giants.  Of the eight LMC sources selected, five were found to have OH masers, with at least one of the non-detections indicating that a detection is possible with increased integration time. This suggests that the sample was well chosen, due to the selection criteria employed: long pulsation period ($P\\gsim800$ d), large pulsation amplitude (${\\Delta}K\\gsim1.2$ mag), late oxygen-rich spectral type ($\\gsim$M8), red ($J-K\\gsim2$ mag), and bolometrically and mid-IR bright ($M_{\\rm bol}\\lsim-6$ mag and $F_{\\rm 25}\\gsim0.3$ Jy, respectively). We can thus be confident in our selection of future potential OH/IR stars. The strength of the OH maser peak in the LMC detections can be readily scaled from the strength of the OH maser peak in galactic centre sources on the basis of the differences in bolometric luminosity alone, without the need for an additional scaling due to the difference in metallicity. This provides some comfort for future searches for OH masers in the SMC or other low-metallicity environments.  The LMC sources show a pronounced asymmetry between the strength of the blue and redshifted emission, much more so than OH/IR stars in the galactic centre. We propose that this is due to a greater contribution of amplified radiation originating from the star itself, both in an absolute sense (bolometrically brighter stars) and because the lower metallicity of sources in the LMC leads to a smaller contribution of (amplified) circumstellar emission compared with OH/IR stars in the galactic centre.  Of the recipes for deriving mass-loss rates from the OH maser emission from Baud \\& Habing (1983) for the gas mass-loss rate and from Zijlstra et al.\\ (1996) for the dust mass-loss rate the LMC data follow the Zijlstra et al.\\ predictions better than the Baud \\& Habing predictions. This supports the interpretation of Zijlstra et al.\\ that the OH maser strength is a measure of the IR pumping flux and thus of the dust mass-loss rate.  We find that at the same luminosity or pulsation period, the difference in wind speed, $v_{\\rm exp}$, can be fully accounted for by metallicity effects. This agrees with the suggestion by Wood et al.\\ (1998) that the stars with higher $v_{\\rm exp}$ have higher metallicity. In addition, we show that the current observations are consistent with simple radiation-driven wind theory in which the wind speed depends on luminosity, $L$, and metallicity, $z$, as $v_{\\rm exp}{\\propto}z^{1/2}L^{1/4}$.  We compare the wind speed with the escape velocities, evaluated at the stellar surface, in the dust formation zone and further out in the wind. We find that the wind speed exceeds the local escape velocity at an ``escape radius'' of about 1--2$\\times10^4$ R$_\\odot$ from the star (just beyond the dust condensation radius) --- a value which is surprisingly similar for all OH/IR stars in the LMC, AGB stars and red supergiants alike (but different for galactic centre OH/IR stars)."
    },
    "0410/astro-ph0410316_arXiv.txt": {
        "abstract": "We discuss an extension to the Chandrasekhar-Fermi method for the evaluation of the mean magnetic field strength in molecular clouds to cases where the spatial orientation of the field is known. We apply the results to M17, using previously published data. ",
        "introduction": "There exist few techniques that allow for the measurement of quantities that characterize the magnetic field in molecular clouds. At millimeter and submillimeter wavelengths, the orientation of the magnetic field is most commonly traced using polarimetry measurements from dust continuum emission \\citep{Hildebrand 1988}. The strength of the magnetic field (in general, its line-of-sight component) can only be directly measured via the Zeeman effect (e.g., \\citet{Crutcher 1999, Brogan 2001}), usually at longer wavelengths. In order to gather as much information as possible about the magnetic field, the so-called Chandrasekhar Fermi (CF) method \\citep{Chandra 1953} is often used to infer the strength of the plane-of-the-sky component of the field. Because this is achieved with the same polarimetry data that give the orientation of the sky-projected magnetic field, the CF method can act as a bridge between the polarimetry and Zeeman observations to provide an estimate for the magnitude of the mean field strength in a given cloud.  In this paper, we discuss how a simple extension of the CF method can be used alone, i.e., without the need of Zeeman data, to infer the magnitude of the magnetic field; not only the strength of its plane-of-the-sky component. Furthermore, it will also be shown that, contrary to the original CF method which only really works well when the magnetic field is located close enough to the plane of the sky, our generalization is valid regardless of the field's orientation in space. However, this can only be accomplished if and when the spatial orientation of the magnetic field is known. That is to say, not only the orientation of its projection on the plane of the sky is needed (from polarimetry), but also its inclination to the line of sight. This last piece of information can be obtained through the technique of \\citet{Houde 2002} which relies on the availability of spectroscopic measurements from suitable neutral and ionic molecular species, as well as polarimetry.  Finally, we apply our extension to the CF method to already published data for the M17 molecular cloud \\citep{Houde 2002}, and infer a value for the magnitude of the mean magnetic field for this object. ",
        "conclusions": ""
    },
    "0410/astro-ph0410302_arXiv.txt": {
        "abstract": "We report first results of a multifrequency campaign from radio to hard X-ray energies of the prominent \\gray\\ blazar 3C~279 during the first year of the INTEGRAL mission. The variable blazar was found at a low activity level, but was detected by all participating instruments. Subsequently a multifrequency spectrum could be compiled. The individual measurements as well as the compiled multifrequency spectrum are presented. In addition, this 2003 broadband spectrum is compared to one measured in 1999 during a high activity period of 3C~279. ",
        "introduction": "The EGRET experiment aboard CGRO has identified about 90 blazar-type AGN emitting high-energy ($>$100~MeV) $\\gamma$-rays \\citep{Hartman99}. Among the most prominent ones is 3C~279, an optically violently variable (OVV) quasar, located at a redshift of 0.538. It shows rapid variability in all wavelength bands, polarized emission in radio and optical, superluminal motion, and a compact radio core with a flat radio spectrum. These properties put the quasar 3C~279 into the blazar sub-class of AGN. According to the unified model of AGN, blazars are sources which expel jets close to our line-of-sight. The blazar emission is predominantly of non-thermal origin, showing a typical two-hump spectrum from radio to \\grays. It is generally believed that in blazars the radio--through-optical/UV continuum is synchrotron radiation generated by relativistic electrons in a magnetized jet and the high-energy continuum from X-rays to \\grays\\ is due to Comptonization of soft photons by the non-thermal jet electrons. Most emission models have been developed within this scenario. Key predictions of the different models which can be tested by multifrequency observations are the predicted spectral shapes for blazars from radio to \\gray\\ energies, and the predicted temporal variations as function of energy.  3C~279 was detected by EGRET several times \\citep[e.g.,][]{Hartman01}, showing strong flaring activity with flux changes up to a factor of 100. 3C~279 was also detected by CGRO/COMPTEL at low-energy (1-30~MeV) \\grays\\ \\citep[e.g.,][]{Williams95} and CGRO/OSSE (50 keV - 1 MeV) at the transition from hard X-rays to \\grays\\ \\citep{NaronBrown95}, i.e., throughout the whole INTEGRAL band and above. Because of the low-significance OSSE detections, which only occurred on or near flaring periods \\citep{NaronBrown95}, not much is known on the hard X-ray properties of 3C~279. To improve on that INTEGRAL observed the blazar in its AO-1 period. To correlate the hard X-ray and multifrequency properties of 3C~279, the INTEGRAL observations were supplemented by multifrequency coverage, i.e., contemporaneous radio and mm, near-infrared and optical observations, as well as X-ray observations.  In this paper we present first results of our multiwavelength campaign on 3C~279 in 2003 (not all data are available yet). Because of the limited page number, we concentrate on presenting the main observational results. A more detailed presentation, including a discussion on the scientific implications of the new results, will be given in a later paper (Collmar et al., in prep.). ",
        "conclusions": "A multiwavelength campaign on the prominent and variable \\gray\\ blazar 3C~279, including INTEGRAL, was carried out in 2003. According to the longterm optical R-band lightcurve, the source was observed in its lowest optical activity state of the last decade. Because the source was detected by all participating instruments, a scientifically interesting low-state multifrequency spectrum could be compiled from radio to hard X-ray energies. It has an unprecedented coverage in the high-energy domain by Chandra and INTEGRAL, and provides new constraints for the modelling of 3C~279."
    },
    "0410/astro-ph0410628_arXiv.txt": {
        "abstract": "{ In a previous paper we were able to measure the abundance of uranium and thorium in the very-metal poor halo giant BPS \\object{CS\\,31082-001}, but only obtained an upper limit for the abundance of lead (Pb).  We have got from ESO 17 hours of additional exposure on this star in order to secure a detection of the minimum amount of lead expected to be present in \\object{CS\\,31082-001}, the amount arising from the decay of the original content of Th and U in the star. We report here this successful detection.  We find an LTE abundance $\\log\\mathrm{(Pb/H)}+12=-0.55 \\pm 0.15$ dex, one dex below the upper limits given by other authors for the similar stars \\object{CS\\,22892-052} and \\object{BD\\,+17\\degr\\,3248}, also enhanced in {\\it r}-process elements. From the observed present abundances of Th and U in the star, the expected amount of Pb produced by the decay of $^{232}$Th, and $^{238}$U alone, over 12 -- 15 Gyr is $-0.73\\pm 0.17$ dex. The decay of $^{235}$U is more difficult to estimate, but is probably slightly below the contribution of $^{238}$U, making the contribution of the 3 actinides only slightly below, or even equal to, the measured abundance. The contribution from the decay of $^{234}$U has was not included, for lack of published data. In this sense our determination is a lower limit to the contribution of actinides to lead production. We comment this result, and we note that if a NLTE analysis, not yet possible, doubles our observed abundance, the decay of the 3 actinides will still represent 50 per cent of the total lead, a proportion higher than the values considered so far in the literature. ",
        "introduction": "The detection of uranium in an old, very metal-poor star of the galactic halo, BPS \\object{CS\\,31082-001}, was first reported in Cayrel et al. (\\cite{CHB01}). A greatly improved analysis, Hill et al. (\\cite{HPC02}), (quoted as paper I) was made possible by a redetermination of crucial atomic data by Nilsson et al. (\\cite{NIJL02, NZL02}).  Hill et al. have determined the abundance of U~($\\log(\\mathrm{U/H})+12. =-1.92 \\pm 0.11$) and of Th ($\\log(\\mathrm{Th/H}) + 12 = -0.98 \\pm 0.05$), in the usual scale $\\log(n_H)=12.0$, in \\object{CS\\,31082-001}. These abundances have been used as cosmo-chronometers, comparing them to theoretical estimates of the initial production ratio. The time $ \\Delta t$ in Gyr elapsed from the formation of the two actinides to now, is linked to the production ratio $\\mathrm{(U/Th)_0}$ and the ratio measured in the star $\\mathrm{(U/Th)_{now}}$ by the simple relation : $ \\Delta t = 21.76[ \\log(\\mathrm{(U/Th)_0}) - \\log(\\mathrm{(U/Th)_{now}})]$ where the coefficient 21.76 is derived from the half-lives of $^{232}$Th and $^{238}$U.  The superiority of the pair U/Th over the pair Th/Eu has been amply demonstrated, for example in Goriely \\& Clerbaux (\\cite{GC99}), Goriely \\& Arnould (\\cite{GA01}), or Wanajo et al. (\\cite{WII02}, see their fig.7).  As both U and Th decay to the stable element lead, it is of great interest to know the abundance of lead in the star. In Hill et al. (\\cite{HPC02}), only an upper limit to the lead abundance was given, and we report here the result of a new observation obtained at ESO Paranal to get this abundance. The time requested was 17 hours, enough to detect the minimum amount of lead coming from the decay of the observed elements $^{238}$U and $^{232}$Th into $^{206}$Pb, and $^{208}$Pb, respectively. In addition, lead may come from other channels, in particular from the decay of $^{235}$U into $^{207}$Pb, and from more unstable nuclides decaying very quickly to lead, such as $^{234}$U. ",
        "conclusions": "In one of the very metal-poor stars showing a large enhancement of {\\it r}-process elements we have now a true determination of the lead abundance, instead of upper limits, only.  This abundance is very low,$-0.55 \\pm 0.15$\\ dex in \\object{CS\\,31082-001}, about one dex below the former upper limits in \\object{CS\\,22892-052} and \\object{BD\\,+17\\degr\\,3248}.  Also, our result shows that, in the purely $r$-process enriched photosphere of \\object{CS\\,31082-001}, most of lead results from the decay of $^{232}$Th , $^{235}$U, and $^{238}$U.  This places a limit on the amount of $^{235}$U which has contributed to the production of $^{207}$Pb, as well as on that of $^{234}$U. A non-LTE analysis of the spectrum is highly desirable, but hampered so far by the lack of photoionization cross-sections for Pb I."
    },
    "0410/astro-ph0410134_arXiv.txt": {
        "abstract": "We present results from VLBA observations of the 1667 and 1665~MHz OH lines in two OH megamaser galaxies IRAS\\,12032+1707 and IRAS\\,14070+0525. For IRAS\\,12032+1707 we also present Arecibo \\hi\\, absorption data. Almost all OH emission previously detected by single dish observations has been recovered in IRAS\\,12032+1707 and is found on a compact scale of $<$100 pc. The emission shows an ordered velocity field that is consistent with a single disk However, the data is also consistent with a scenario including two physically different gas components. We explore this possibility, in which the two strongest and most blueshifted maser features are identified as the tangent points of a circumnuclear torus.  The redshifted maser features would be extended in a direction perpendicular to the torus, which is in the direction of the merger companion. Thus, redshifted emission could be associated with an inflow triggered by a tidal interaction. \\hi\\ absorption covers the velocities of the redshifted maser emission, suggesting a common origin. In the second source, IRAS\\,14070+0525, a large fraction of the OH emission is resolved out with the VLBA.  We find no significant evidence of an ordered velocity field in this source, but these results are inconclusive due to a very low signal-to-noise ratio. ",
        "introduction": "Since the discovery of OH megamaser emission in Arp\\,220 \\citep{baan82}, the detection of $\\sim$95 OH megamaser have been reported \\citep{baan98,darling02,martin88,staveley87}. These single dish studies show that the shapes of the OH spectral lines are often complex, and cover velocities in the range 10\\kms\\ up to a few 1000 \\kms in each source.  Subsequent VLBI investigations of a few nearby OH megamaser galaxies, such as III\\,Zw\\,35 and Mrk\\,231, have demonstrated that the bulk of the OH maser emission arises in circumnuclear disks or tori \\citep{pihlstrom01,kloeckner03}.  However, there are also indications that a single disk component cannot account for {\\it all} maser emission, e.g.\\ in Mrk\\,273; \\citep{kloeckner04}. Early single dish studies suggested the presence of blueshifted wings in several OH megamaser spectra \\citep{baanetal89}. These have been interpreted as outflows, which possibly could provide an explanation for the emission that does not fit comfortably within a disk model.  The existence of outflows as well as inflows is not surprising, given that OH megamaser galaxies are exclusively associated with (Ultra)Luminous Infra-Red Galaxies ([U]LIRGs). (U)LIRGs probably represent short periods of nuclear starburst activity triggered by merger events. Such mergers will rapidly transport gas to the central regions, causing gas to fall inwards towards the nuclei. The combined effects of supernova explosions and stellar winds generated in this nuclear starburst can entrain the interstellar medium in outflows with velocities of $100-1000$ \\kms\\ \\citep{heckman90, alton99}.  It is also possible that OH maser emission occurs in both of the merging nuclei, as has been seen in Arp\\,220 \\citep{diamond89}.  With a slight offset in systemic velocity between the two merging nuclei, the combined spectrum observed with a single dish telescope would be relatively broad. However, for the very broadest OH megamaser lines (exceeding 1000 \\kms\\,), it is hard to interpret the lines either as originating from a single disk component or from a pair of nuclei.  A combination of orbital mechanics, possible disk rotation, molecular inflow and outflow, and the velocity separation of the two main lines at 1667~MHz and 1665~MHz is likely to make up the velocity range in such OH megamaser spectra.  Earlier global VLBI experiments have reported the presence of 100 \\kms\\, broad OH maser lines on parsec scales \\citep{diamond99}. In this paper we will concentrate on investigating the cause of OH megamaser lines with widths exceeding 1000 \\kms. We report on VLBA observations of two OH megamaser galaxies, IRAS\\,14070+0525 and IRAS\\,12032+1707, that have full width zero intensity velocities of $1500-2000$\\,\\kms. These objects are additionally interesting because of their high redshifts ($z=0.217$ and $z=0.265$).  Furthermore, their high OH luminosities of $L_{\\rm OH}=1.3\\times 10^4L_{\\odot}$ and $L_{\\rm OH}=1.2\\times 10^4L_{\\odot}$ respectively make the term 'gigamaser' suitable, and as such these galaxies are two of the most powerful OH megamaser galaxies known. The main aim of the current observations was to determine whether the broad lines are associated with inflows, outflows, rotating structures or violent kinematics. This paper also presents Arecibo \\hi\\, absorption data for one of the sources, IRAS\\,12032+1707. ",
        "conclusions": "\\subsection{IRAS\\,12032+1707 Continuum}\\label{contdisc} The continuum in IRAS\\,12032+1707 is only tentatively detected at the shortest u,v distances (Sect.\\ \\ref{cont}). In part, this result could be due to calibration errors, but it also indicates that the majority of continuum emission is found on scales larger than 260 pc (Sect.\\ \\ref{cont}). That implies an upper limit of the brightness temperature to be around 3$\\times 10^6$ K. Such a brightness temperature and spatial distribution would be consistent with a starburst origin similar to that seen in Arp\\,220 \\citep{smith98} and III\\,Zw\\,35 \\citep{pihlstrom01}.  Many ULIRGs fall on the radio-FIR correlation for starbursts \\citep{helou85, yun01}. This correlation can be understood by assuming that the radio emission is arising from starburst induced supernova remnants and from HII regions. Following the definitions from \\citet{helou85}, the correlation is quantified by a logarithmic ratio between the far-infrared flux density ($FIR$) and the 1.4 GHz radio flux density ($S_{\\rm 1.4 GHz}$):  \\begin{equation} q={\\rm log}\\left(\\frac{FIR}{3.75\\times 10^{12} \\rm W m^{-2}}\\right)-{\\rm log}\\left(\\frac{S_{\\rm 1.4GHz}}{\\rm W m^{-2}Hz^{-1}}\\right) \\end{equation}  where $FIR$ is defined by  \\begin{equation} FIR=1.26\\times10^{-14}(2.58S_{\\rm 60 \\mu m}+S_{\\rm 100 \\mu m}) \\rm W m^{-2} \\end{equation}  For the infrared selected galaxies in the $IRAS$ 2 Jy sample, the mean value of $q$ is $\\simeq2.35$ \\citep{yun01}. Given the FIR luminosity and radio flux density of IRAS\\,12032+1707, we estimate $q=1.77$. This value appears low compared with the mean of 2.35 and could imply the presence of an AGN. This appears contradictory to the non-detection of a compact radio AGN core. However, radio-excess objects in the $IRAS$ 2 Jy sample are defined as those objects having a radio luminosity that is greater than 5 times larger that predicted by the radio-FIR correlation.  Equivalently, this means objects for which q$\\leq$1.64. Furthermore, IRAS\\,12032+1707 has a very large infrared luminosity (log$(L_{\\rm IR}/L_{\\odot})=12.57$), placing the galaxy at the high tail of the ULIRG luminosity distribution \\citep{kim98}. For the most luminous sources in the $IRAS$ 2 Jy sample, scattering from the radio-FIR correlation is significantly larger than for the weaker sources with weaker FIR flux densities.  Therefore, the relatively low $q$-value for our source is not sufficient to indicate the presence of an AGN component. Other evidence that IRAS\\,12032+1707 is starburst-powered comes from its relatively cool IR color, $F_{\\rm 25\\mu m}/F_{\\rm 60\\mu m}=0.18$ \\citep{kim98}. It is only for warm colors ($F_{\\rm 25\\mu m}/F_{\\rm 60\\mu m}>0.25$) that an AGN component is required to explain the resulting dust temperatures.  \\subsection{IRAS\\,12032+1707 OH and HI}\\label{discuss}  The zero-intensity line width of the OH maser emission in IRAS\\,12032+1707 is around 2000 \\kms, which makes this source the host of one of the broadest OH megamaser lines. As mentioned in the introduction, this velocity range could be due to a combination of effects such as disk rotation and gas flows.  Due to the limited SNR in our data, it is difficult to determine unambiguously the cause of the masers in IRAS\\,12032+1707. Below we discuss a few possible scenarios that could agree with our observations.  \\subsubsection{Multiple disks}\\label{multipledisks}  It has been suggested that very broad OH megamaser lines might be a combination of maser emission occurring in {\\it both} of the merging nuclei, as is the case for the OH megamaser prototype Arp\\,220 \\citep{diamond89}. We can rule out that maser emission in IRAS\\,12032+1707 is the combined emission from the southern and northern nuclei, since it is clear that all maser emission occurs only in the northern nucleus (Sect.\\ \\ref{ohspatial}).  We can not however exclude the possibility that the northern nucleus itself is made up of two more closely interacting galaxies. In Arp\\,220 the nuclei have a projected separation of only 300 pc, which would correspond to an angular resolution of 86 mas at the distance of IRAS\\,12032+1707, much smaller than the near-infrared seeing limit of 0.73\\arcsec reported for the NIR image \\citep{kim02}. This issue requires further investigation by future higher resolution infrared and optical imaging, combined with VLBI maser data.  In the OH megamaser source Mrk\\,231 \\hi\\ absorption has been seen against a faint continuum disk \\citep{carilli98}, and the \\hi\\, velocity field displays a gradient in the same direction as the OH megamaser disk \\citep{kloeckner03}. In particular, the \\hi\\ absorption velocities encompass those of the OH lines. Assuming that \\hi\\ absorption in IRAS\\,12032+1707 has the same close correlation to OH as the \\hi\\ in Mrk\\,231, then the fact that \\hi\\ only partly covers the OH velocities implies that the \\hi\\, absorption is seen against only one of the disks.    \\subsubsection{A single nuclear disk}\\label{onedisk}  A clue to the nature of the OH maser emission must be given by the plot of the centroid position of the different velocity components. Clearly, there is a structured velocity field, with a bulk gradient in the north-south direction. If we assume that all maser emission belongs to a single disk-like structure, with a major axis going through the three most redshifted points at a position angle (PA) of 11$^{\\circ}$ (Fig.\\ \\ref{ir12_radecpos}), we can plot the data as velocity versus distance along this axis (Fig.\\ \\ref{velpos}). A least squares fit to the data points yields a gradient of 55 km\\,s$^{\\rm -1}$pc$^{\\rm -1}$ over 17 pc. In total, this would mean an enclosed mass of 2.2$\\times 10^8$sin$^{\\rm -2}i M_{\\sun}$ within a radius of 8.5 pc. Such a velocity gradient is surprisingly high given the gradients observed in other OH megamaser galaxies. For instance, for both III\\,Zw\\,35 and Mrk\\,231, the velocity gradient is close to 1.5 km\\,s$^{\\rm -1}$pc$^{\\rm -1}$ \\citep{pihlstrom01,kloeckner03}.  In contrast to the case of multiple disks suggested in Sect.\\ \\ref{multipledisks}, a scenario with a single disk for IRAS\\,12032+1707 makes it difficult to understand why the \\hi\\, absorption covers only part of the OH velocity range (assuming that the \\hi\\ and OH are parts of the same neutral gas structure).  \\begin{figure}[t] \\epsscale{0.9} \\rotatebox{-90}{ \\plotone{figure7.ps}} \\caption{IRAS\\,12032+1707: Velocity plotted as a function of position along the axis delineated by the three redshifted components.\\label{velpos}} \\end{figure}  \\subsubsection{A disk and an inflow}  An alternative interpretation of Fig.\\ \\ref{ir12_radecpos} could include two different gas components. A couple of facts support the possibility of two physically different OH maser components in IRAS\\,12032+1707. Firstly, single-dish variability studies show that the two strongest and most blueshifted peaks have a variable flux density while the redshifted emission displays no significant variability (J.\\ Darling, priv.\\ comm). Secondly, the \\hi\\ absorption velocities cover only the redshifted OH components (Sect.\\ \\ref{HIres}). If the \\hi\\ has a common origin with the redshifted OH gas, the blueshifted and redshifted masers are likely to be physically different gas structures.  Studying the relative location of maser components in Fig.\\ \\ref{ir12_radecpos}, we suggest that the two most blueshifted masers (which also are the brightest ones) could represent the tangent points of a circumnuclear disk.  With an inclined disk (or torus), the tangent points will have the longest paths of velocity coherent gas, resulting in regions of very strong maser emission.  This effect has been observed in III\\,Zw\\,35 \\citep{pihlstrom01}. The rotational velocity of such a torus or disk structure in IRAS\\,12032+1707 would be 118 sin$^{-2}i$ \\kms\\, at a radius of 3 pc, yielding a modest enclosed mass of 5$\\times 10^6$sin$^{-2}i M_{\\sun}$.  If the two blueshifted masers define the major axis of a disk component, the redshifted masers are directed more or less perpendicular to the disk (Fig.\\ \\ref{ir12_radecpos}). By studying the optical and NIR images of IRAS\\,12032+1707 \\citep{veilleux02} it can be seen that the southern optical nucleus is located slightly west of south, with respect to the northern nucleus. Tentatively, the redshifted masers could arise in gas that has been disturbed by tidal interactions between the merging nuclei. With the caveat of over-interpreting our low SNR data, it is tantalizing to note that there is a weak hint in Fig.\\ \\ref{velpos} of gas closer to the nucleus moving with higher velocity with respect to the nucleus, than does the gas further away.  In this scenario we suggest that the \\hi\\ and the redshifted OH originate in the same gas component. Any gas associated with the \\hi\\ absorption must thus necessarily be in front of the background continuum. Assuming the blue masers define a disk, it is plausible that the \\hi\\ is seen in absorption against a weak radio continuum emitted by this disk. Given a limited field of view in the VLBA observations, in addition to an only tentative detection of the radio continuum, it is justified to question whether the radio continuum in fact coincides at all with the maser emission. We note that VLA A-array observations have showed that all radio continuum emission originates in the northern optical nucleus (J.\\ Darling, priv.\\ comm.), and the OH maser emission is located in the same region. Since \\hi\\ absorption requires the presence of a background continuum, this is consistent with the \\hi\\ absorption and OH masers occurring in the same region, at least on scales of a few kpc (corresponding to the VLA A-array resolution).  The optical systemic velocity of the nucleus is 65055 \\kms\\, with an 1$\\sigma$ error of $\\pm70$\\kms\\,\\citep{kim98}. Comparing the optical systemic velocity with the OH and \\hi\\ velocities is hard since the optical redshift is measured for the whole IRAS\\,12032+1707 system, and thus does not separate the optical nuclei. In addition, the velocity is determined from optical emission lines that are often broad and might not properly reflect the systemic velocity since the emission lines may be associated with gas in motion with respect to the nucleus. Nevertheless, within a 3$\\sigma$ limit the systemic velocity is consistent with redshifted maser components {\\it falling in} towards the nucleus. \\citet{veilleux99} reports on optical emission lines in IRAS\\,12032+1707, that are broad compared to the typical ULIRG, consistent with the presence of an infalling gas component.  Neutral gas flows have been observed previously in other galaxies. For instance, VLA observations of the 1667, 1665 and 1612 MHz OH lines in the Galactic Center have revealed gas streaming inwards from the CND (the CircumNuclear Disk) toward SgrA* \\citep{karlsson03}. The streamer can be seen over a velocity range of $\\sim$100 \\kms.  Furthermore, \\hi\\ studies have extensively been used to trace merger dynamics, and one example is the detailed studies of \\hi\\ emission in `The Antennae'. \\citet{hibbard01} have mapped high \\hi\\ column densities exceeding $10^{22}$ cm$^{-2}$ in the region of the merging disks in this source. It is not unlikely that similar column densities could be probed by \\hi\\ absorption in a merger like IRAS\\,12032+1707.   \\subsection{IRAS\\,14070+0525}  A major difficulty with interpreting the OH megamaser emission in IRAS\\,14070+0525 is that we only detect a fraction of the total single dish emission. As already mentioned, a reason why so much emission is resolved out may be the lack of short baselines in our VLBA observations. The offset between the two detected peaks is not significant. Optical and NIR imaging cannot distinguish more than one nucleus in IRAS\\,14070+0525, and this source is interpreted as a merger where the nuclei have more or less coalesced \\citep{veilleux02}. Two nuclei could thus be located close to each other, and violent gas kinematics adding to the line width would be anticipated. More sensitive VLBI observations probing the emission on all spatial scales will be needed to gain a better understanding of the OH megamaser emission in this galaxy.  IRAS\\,14070+0525 is classified optically as a Seyfert 2 galaxy \\citep{kim98}. However, the color given by $F_{\\rm 25\\mu m}/F_{\\rm 60\\mu m}$ flux density ratio is only 0.13, and does not indicate the presence of an AGN. Contrary to IRAS\\,12032+1707, IRAS\\,14070+0525 falls on the radio-IR correlation with $q=2.4$ \\citep{condon98,kim98}, and no compact radio emission associated with an AGN was detected. Hence, the radio emission is diffuse and consistent with a starburst. Given the low 1.4 GHz flux density of 5\\,mJy \\citep{condon98}, a slightly resolved starburst would be expected to fall below our detection limit in these observations."
    },
    "0410/astro-ph0410244_arXiv.txt": {
        "abstract": "We report an optical photometric survey of 0.89~deg$^2$ of the Orion OB1b association centered on $\\sigma$~Ori.  This region includes most of the $\\sigma$~Ori cluster, the highest density region within Orion OB1b.  We have developed a statistical procedure to identify the young, low-mass, pre-main sequence population of the association.  We estimate that the cluster has $\\sim$160 members in the mass range (0.2$\\leq$M$\\leq$1.0~M$_{\\odot}$).  The cluster has a radius of $\\sim$3-5~pc and an estimated age of 2.5$\\pm$0.3~Myrs.  We estimate that the total mass of the cluster is 225$\\pm$30~M$_{\\odot}$.  This mass is similar to the estimated mass of the $\\sim$5$\\times$10$^5$ year old cluster NGC~2024.  NGC~2024 and $\\sigma$~Ori appear to be a well matched pair of clusters, except for the $\\sim$2~Myr difference in their ages. ",
        "introduction": "\\label{intro}  The Orion OB1 association is one of the most intensively studied local star forming regions.  Several stages of star formation are observed in Orion OB1 spanning ages from less than 10$^6$ years up to $\\sim$10$^7$ years.  Orion is notable as the largest nearby site of very recent, perhaps ongoing star formation.  Much of this activity is hidden within the main molecular clouds of the Orion Complex, the Orion A and Orion B clouds \\citep{gs89}.  There are several deeply embedded clusters in the Orion B cloud (L1630 Lada et al.\\ 1991) aside from the well known optically visible OB association.  \\subsection{The subgroups of Orion OB1}\\label{sec:intro_groups}  \\citet{blaauw64} divided the optically visible association into four groups labeled a-d in order of decreasing estimated age.  The groups were identified based upon their degree of concentration, their association with interstellar matter, age, and spatial positions.  The spatial boundaries of the groups are somewhat arbitrary in part because the photometric overlap between different groups makes it difficult to unambiguously assign individual O and B stars to a particular group (see figure 5 from Blaauw 1964).  Also, the large size and low spatial density of association members confuse efforts to determine membership from converging proper motion vectors \\citep{brown94}.  The proper motions are very small because Orion lies in the anti-direction of the solar reflex motion.  The nearest and oldest of the sub-associations, Orion OB1a, is roughly 10$^7$ years old and $\\sim$330~pc from the Sun \\citep{brown04}.  Orion OB1b contains the belt stars and the stars around them, including $\\sigma$~Ori.  Orion OB1b lies about 440~pc from the Sun and has an age of 2-5~Myrs \\citep{brown04}.  Orion OB1c contains the stars around the sword region.  Orion OB1d is the Orion Nebula Cluster (ONC).  The ONC lies at a distance of $\\sim$450~pc and is less than 1~Myr old. Much of the work on Orion OB1 has concentrated on the ONC (e.\\ g.\\ \\citet{hill97} in the optical: \\citet{garm00,flac03,feig03} in X-rays: \\citet{carp_hill01} in the NIR: \\citet{stassun99,herbst02,rebull01} for photometric variability: \\citet{scally99,simon99} for binaries). See \\citet{odell01} for a recent review.  As a region of active star formation, the ONC provides a look at the final stages of accretion.  Orion OB1b, including the $\\sigma$~Ori region, is a fossil star forming region which presents the end products of star formation.  \\subsection{Clustered Star Formation in Orion}  While the ONC is the best known example of clustered star formation in Orion, it is only one of nearly a score of concentrations of very young stars in the Orion OB1 and $\\lambda$~Ori associations.  Many of these clusters are deeply embedded within the giant molecular clouds ({GMC}s) of the star forming region.  In addition to the partially embedded ONC, \\citet{lada91} catalog four embedded clusters within the Orion A clouds and three within the Orion B cloud.  These embedded clusters all have radii of roughly 1~pc.  This is roughly the same size as the high density cores in GMCs \\citep{ll03}.  The number of members in these clusters ranges from 43 (L1641N) to $>$300 (NGC~2024).  \\citet{carp00} used 2MASS observations of the Orion A and Orion B clouds to estimate that more than 50\\% of all the stars in these clouds are located in clusters.  \\citet{gomez98} report three clusters in the $\\lambda$~Ori OB association with radii of $\\sim$30$^{\\prime}$ (3.5~pc).  Roughly 80\\% of the pre-main sequence ({PMS}) stars in the $\\lambda$~Ori region belong to one of these three clusters.  South of Orion's belt (${\\delta}<$-1.5) \\citet{gomez98} identify seven clusters of low-mass PMS stars among the strong H$\\alpha$ sources of the Kiso H$\\alpha$ survey \\citep{kiso1,kiso2,kiso3,kiso4}.  Five of these clusters correspond to known clusters of O and B stars: NGC 1977, the ONC, OMC-2, the upper sword, and the lower sword.  Two consist solely of low-mass stars.   \\citet{ll03} compiled a list of all published embedded clusters within two~kpc of the sun.  They found that there are several times too many young embedded clusters compared to the number of bound open clusters. They concluded that fewer than 10\\% of young embedded clusters remain as recognizable entities for more $\\sim$10$^7$ years after their natal molecular gas is dispersed.     \\subsection{The $\\sigma$~Ori Cluster}  \\citet{walt98} reported the identification of a concentration of low-mass PMS stars around the O9.5V star $\\sigma$~Ori.  \\citet{walt98} spectroscopically identified 104 PMS stars within 30$^{\\prime}$ of $\\sigma$~Ori.  Photometry of 0.15 deg$^2$ containing 45 spectroscopically confirmed PMS stars suggested another 65 likely PMS stars with V$<$19.  With 110 PMS stars in the 0.15 deg$^2$ area, \\citet{walt98} concluded that the spatial density of PMS stars was at least 700 deg$^{-2}$ in the region around $\\sigma$~Ori.  The spatial density of PMS stars in the spectroscopic survey decreased with distance from $\\sigma$~Ori, indicating that the low-mass PMS population near $\\sigma$~Ori formed part of a cluster with a radius of $\\sim$0.5$^{\\circ}$.  The $\\sigma$~Ori cluster has since proven to be a rich hunting ground for sub-stellar objects ranging from brown dwarfs \\citep{bejar99, bejar01, barr03} down to free-floating ``cluster planets''.  One candidate member has a mass as low as 3 Jupiter masses \\citep{zo00,zo02,mz03}, although \\citet{burg04} suggest that it is a foreground brown dwarf.  Sizes, luminosities, and ages depend upon the assumed distance to the $\\sigma$~Ori cluster.  Most papers on the brown dwarf population of the $\\sigma$~Ori cluster have used the Hipparcos distance for $\\sigma$~Ori.  Throughout this paper we use the 440~pc distance to Orion OB1b \\citep{brown04} as the distance to the cluster.  We prefer this value to the 350~pc Hipparcos distance to $\\sigma$~Ori \\citep{hipparcos} because the uncertainties of the Hipparcos measurements in Orion are large and the distance to Orion OB1b is averaged over many stars \\citep{dezeeuw99,brown04}.  \\subsection{Searching for New Low-Mass PMS Stars}  Several methods have been used to identify low-mass PMS stars in Orion and other star forming regions \\citep{walt00}.  H$\\alpha$ emission is an efficient means of finding classical T Tauri stars ({cTTs}).  Large scale surveys such as the KISO H$\\alpha$ survey \\citep{kiso1,kiso2,kiso3,kiso4} have identified many candidate low-mass PMS stars in Orion.  A drawback of low resolution H$\\alpha$ surveys is that they cannot distinguish between foreground dMe stars and low-mass PMS stars.  Also, H$\\alpha$ surveys are biased towards finding cTTs which have strong H$\\alpha$ emission.  Many, if not most of the low-mass members of the $\\sigma$~Ori cluster must be weak T Tauri stars ({wTTs}) since few of the members have strong H$\\alpha$ emission.  This is typical of regions where the natal gas has dispersed \\citep{breceno01}.  Young low-mass stars are magnetically active, which makes them bright X-ray sources.  X-ray surveys of star forming regions can detect both the wTTs and cTTs populations.  For large survey regions, the ROSAT all sky survey ({RASS}) may be used to search for low-mass PMS stars (i.\\ e.\\ Sterzik et al.\\ 1995; 2004).  The RASS has a limiting flux of $\\sim$10$^{-14}$~erg~cm$^{-2}$~s$^{-1}$.  Assuming a typical X-ray to V band flux ratio of 10$^{-3}$, only PMS stars brighter than V$\\sim$15 could be detected by the RASS \\citep{walt00}.  Also, the RASS sample is strongly contaminated by many young, X-ray active foreground stars which may be mistaken for PMS stars \\citep{breceno97}.  High to medium resolution spectroscopic observations unambiguously identify young low-mass PMS stars through the detection of the 6707~\\AA~line of Li I which indicates youth \\citep{walt00}. Spectroscopic observations also distinguish between cTTs and wTTs and yield radial velocities.  The primary disadvantage of using spectroscopy to find new low-mass PMS stars is that high or medium resolution spectroscopy of numerous faint stars over many square degrees requires significant amounts of time on large telescopes.  Broad band optical surveys with one meter class telescopes can detect the low-mass PMS stars in young OB associations with reasonable exposure times \\citep{wolk96,breceno01,breceno02,thesis}.  With broad band photometry, the challenge has always been to separate the PMS association members from main-sequence field stars.  Variability measurements are an extremely effective technique for identifying young, low-mass PMS stars \\citep{breceno01,breceno02}.  However, this method is potentially biased against PMS stars with only low amplitude photometric variations.  \\citet{wolk96} reported optical (UBVR$_C$I$_C$) photometry of X-ray-selected PMS stars near $\\sigma$~Ori.  He found that the X-ray-selected PMS stars occupied a distinct locus on the color-magnitude diagram ({CMD}).  Only about half of the stars in this locus were X-ray-selected PMS stars.  Spectroscopic observations of 26 non-X-ray detected stars in this PMS locus showed that about 70\\% were PMS stars.  This suggested that the low-mass PMS population in Orion OB1b could be efficiently identified by using single epoch photometry to select the stars which lie in the PMS locus.  Single epoch photometry alone cannot definitively identify any individual star as a PMS star, but it can be used to determine the size and spatial distribution of the low-mass PMS population of the association.  \\subsection{This work}  We have completed a BVRI survey of 0.89~deg$^2$ around $\\sigma$~Ori. Our data permit us to measure several properties of the $\\sigma$~Ori cluster, including the radius, the total mass, and the age of the cluster.  The mass and radius of the cluster determine the escape velocity of the cluster.  A small escape velocity would indicate that the cluster is not bound.  The total mass and radius of the cluster also provide a context which allows us to look at the $\\sigma$~Ori cluster as part of a hierarchy of star forming regions of various sizes and masses.  A reliable age for the cluster will place the cluster and its member stars on an evolutionary sequence with other young clusters such as NGC~2024 and the ONC.  Armed with an understanding of the cluster's mass, radius, and age, we can identify similar clusters at different evolutionary stages, or explore how the richness of clusters influences the evolution of protoplanetary disks. ",
        "conclusions": "It is possible to use single epoch optical photometry to statistically identify the low-mass PMS population of nearby, young clusters and associations.  This method is complementary to surveys such as that of \\citet{breceno01} which identify likely PMS stars from their variability.  Variability surveys provide a more secure identification of PMS stars, but may be biased to the most variable members of the PMS population.  Single epoch photometry can identify the whole PMS population on a statistical basis (at the price of greater contamination from field stars).  The size and spatial distribution of the PMS population may be determined through single epoch photometry, although individual stars can be securely identified as PMS stars only through follow up spectroscopy.  In a region such as the $\\sigma$~Ori cluster, the high density of PMS stars makes the number of field stars in the PMS locus relatively small.  In associations with a lower density of PMS stars there will be greater contamination of the PMS locus by field stars.  For associations with few stars or a very low surface density of low-mass members this method will not work well because the contrast between the PMS population and the field population will be very small.   The $\\sigma$~Ori cluster has an age of 2.5$\\pm$0.3~Myr.  The low-mass members of the cluster have a spatial distribution which is broadly consistent with a King model.  The relatively small number of stars in our sample can only weakly constrain the parameters of a King model. The radius of the cluster is about 30$^{\\prime}$ (3.5~pc).  We estimate that the total mass of the $\\sigma$~Ori cluster is $\\sim$225$\\pm$30~M$\\odot$.  This is very similar to the NGC~2024 cluster, but roughly 10 times less massive than the ONC.  With such a small mass, the cluster cannot be gravitationally bound unless it has an unexpectedly small velocity dispersion."
    },
    "0410/astro-ph0410072_arXiv.txt": {
        "abstract": "The determination of the metal abundances in the neutral interstellar medium (ISM) of dwarf star-forming galaxies is a key step in understanding their physical and chemical evolution. This type of investigation has been possible in the last 5 years thanks to FUSE. We will give a flavor of the issues involved by presenting the work that we are performing in this astrophysical field.\\looseness=-2 ",
        "introduction": "Dwarf galaxies are the most common type of galaxies in the local Universe and may have also played an important role in the past as building blocks from which larger systems formed by merging. Present-day dwarfs are the leftovers of this building process, and as such they could have had a completely different evolutionary path compared to their high-reshift counterparts. However, local star-forming dwarfs are particularly interesting from a cosmological point of view. They are chemically poorly evolved systems which are undergoing star formation (SF), thus the closest analogue to primeval galaxies. They are also perfect targets for FUSE, since SF means hot young and massive stars that emit the bulk of their energy in the rest-frame UV rich in transitions of the most abundant ions. These galaxies are characterized by a large reservoir of H\\,{\\sc i}. FUSE is thus particularly suited to study the physics, chemistry and kinematics of their ISM.\\looseness=-2 ",
        "conclusions": "The offset in metal content between neutral ISM and H\\,{\\sc ii} regions in dwarf star-forming galaxies is probably one of the unexpected great results of FUSE in the past five years. However, this area of research is still pretty new and many unknowns and uncertainties still affect the interpretation of the data. It is thus premature to draw conclusions and more targets with a data quality similar to that of NGC~1705, still need to be investigated before having all the pieces of this intriguing puzzle put together."
    },
    "0410/astro-ph0410558_arXiv.txt": {
        "abstract": "I review current results from searching for galaxies giving rise to damped \\lya\\ absorbers (DLAs) at $z<1$.  Using 14 confirmed DLA galaxies, I further show that intermediate-redshift galaxies possess large \\hI\\ envelope out to $24-30\\ h^{-1}$ kpc radius.  The photometric and spectral properties of these galaxies confirm that DLA galaxies are drawn from the typical field population, and not from a separate population of low surface brightness or dwarf galaxies. Comparisons of the ISM abundances of the DLA galaxies and the metallicities of the absorbers at large galactic radii suggest that some DLAs originate in the relatively unevolved outskirts of galactic disks. ",
        "introduction": "Damped \\lya\\ absorbers (DLAs) observed in the spectra of background QSOs probe neutral gas regions commonly seen in nearby galaxies (with \\hI\\ column density $N(\\hI)\\ge 2\\times 10^{20}$ \\cmjj).  In principle, these DLAs offer a means of studying the galaxy population at high redshift using galaxies selected uniformly based on known neutral gas content, rather than optical brightness or color.  Optical spectroscopic surveys for DLAs have demonstrated that DLAs dominate the mass density of neutral gas in the universe and that they contain roughly enough gas at $z=3.5$ to form the bulk of the stars in present-day galaxies (Wolfe \\etal\\ 1995; Storrie-Lombardi \\& Wolfe 2000; P\\'eroux \\etal\\ 2002; Prochaska \\& Herbert-Fort 2004).  But chemical abundance analyses of DLAs at redshift $z < 1.6$ yield sub-solar metallicities in the neutral gaseous clouds (Pettini \\etal\\ 1999; Prochaska \\etal\\ 2003), suggesting that they do not trace the bulk of star formation and are likely to represent a biased sample of galaxies.  Because metallicity varies as functions of morphological type and galactocentric distance, it is impossible to understand the origin of DLAs without first identifying the absorbing galaxies.  \\begin{tiny} \\begin{table}[th] \\caption{Summary of Confirmed DLA Galaxies at $z<1$} \\smallskip \\begin{center} {\\tiny \\begin{tabular}{p{0.8 in}crcrrrcc} \\hline \\hline & & & & \\multicolumn{1}{c}{$\\Delta\\theta$} & \\multicolumn{1}{c}{$\\rho\\times h^a$} & & $M_{AB}(B)^b$ & \\\\ \\multicolumn{1}{c}{QSO} & $z_{\\rm QSO}$ & $z_{\\rm DLA}$ & $\\log N(\\hI)$ & \\multicolumn{1}{c}{($''$)} & \\multicolumn{1}{c}{(kpc)} & \\multicolumn{1}{c}{$AB$} & $-5\\,\\log\\,h$ & Morphology \\\\ \\multicolumn{1}{c}{(1)} & (2) & (3) & (4) & \\multicolumn{1}{c}{(5)} & \\multicolumn{1}{c}{(6)} & \\multicolumn{1}{c}{(7)} & (8) & (9) \\\\ \\hline TON 1480 \\dotfill & 0.614 & 0.0036 & 20.34 & 114.0 & 5.94 & $B=11.5$ & $-18.7$ & S0 \\\\ HS1543$+$5921 \\dotfill & 0.807 & 0.009 & 20.35 & 2.4 & 0.31 & $R=16.5$ & $-15.3$ & LSB \\\\ PKS0439$-$433 \\dotfill & 0.593 & 0.101 & 19.85 & 3.9 & 5.13 & $I=17.2$ & $-19.6$ & disk \\\\ Q0738$+$313 \\dotfill & 0.635 & 0.2212 & 20.90 & 5.7 & 14.23 & $I=20.9$ & $-17.7$ & compact \\\\ Q0809$+$483 \\dotfill & 0.871 & 0.4368 & 20.80 & 1.5 & 5.9 & $R=19.9$ & $-20.3$ & disk \\\\ B2\\,0827$+$243 \\dotfill & 0.939 & 0.525 & 20.30 & 5.8 & 25.42 & $R=21.0$ & $-20.0$ & disk \\\\ PKS1629$+$120 \\dotfill & 1.795 & 0.532 & 20.70 & 3.0 & 13.24 & $R=21.6$ & $-19.2$ & disk \\\\ LBQS0058$+$0155 \\dotfill & 1.954 & 0.613 & 20.08 & 1.2 & 5.67 & $R=23.7$ & $-17.6$ & disk \\\\ HE1122$-$1649 \\dotfill & 2.400 & 0.681 & 20.45 & 3.6 & 17.66 & $I=22.4$ & $-18.8$ & compact \\\\ FBQS\\,0051$+$0041 \\dotfill & 1.190 & 0.740 & 20.40 & 3.3 & 16.87 & $I=22.4$ & $-18.6$ & compact \\\\ EX0302$-$2223 \\dotfill & 1.400 & 1.001 & 20.36 & 3.3 & 18.65 & $R=23.2$ & $-19.3$ & Irr \\\\ \\hline PKS 1127$-$145\\dotfill & 1.187 & 0.313 & 21.71 &  3.8 & 12.2 & $R=22.4$ & $-16.8$ & Irr \\\\ &       &       &       &  3.8 & 12.2 &     22.1 & $-17.4$ & compact \\\\ &       &       &       &  9.8 & 31.5 &     19.1 & $-20.1$ & disk \\\\ &       &       &       & 17.5 & 56.2 &     18.9 & $-20.3$ & disk \\\\ AO0235$+$164 \\dotfill & 0.940 & 0.524 & 21.70 & 2.1 &  9.4 & $I=20.2$ & $-20.3$ & compact \\\\ &       &       &       & 6.4 & 28.0 & $I=20.9$ & $-19.7$ & compact \\\\ FBQS\\,1137$+$3907 \\dotfill & 1.020 & 0.719 & 21.10 & 2.5 & 12.6 & $K=21.4$ & $-18.8$ & compact \\\\ &       &       &       & 1.5 &  7.6 & $K=21.7$ & $-18.5$ & Irr \\\\ \\hline \\multicolumn{9}{l}{a.\\ A $\\Lambda$ cosmology, $\\Omega_{\\rm M}=0.3$ and $\\Omega_\\Lambda = 0.7$ with $h = H_0/(100 \\ {\\rm km} \\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1})$ is adopted throughout.} \\\\ \\multicolumn{9}{l}{b.\\ $M_{{\\rm AB}_*}(B)-5\\log h = -19.6$ (Ellis \\etal\\ 1996).} \\end{tabular}} \\end{center} \\end{table} \\end{tiny}  \\begin{tiny} \\begin{table}[!ht] \\caption{Summary of Candidate DLA Galaxies at $z<1$} \\smallskip \\begin{center} {\\tiny \\begin{tabular}{p{0.8 in}crcrrrcc} \\hline \\hline & & & & \\multicolumn{1}{c}{$\\Delta\\theta$} & \\multicolumn{1}{c}{$\\rho\\times h$} & & $M_{AB}(B)$ & \\\\ \\multicolumn{1}{c}{QSO} & $z_{\\rm QSO}$ & $z_{\\rm DLA}$ & $\\log N(\\hI)$ & \\multicolumn{1}{c}{($''$)} & \\multicolumn{1}{c}{(kpc)} & \\multicolumn{1}{c}{$AB$} & $-5\\,\\log\\,h$ & References \\\\ \\multicolumn{1}{c}{(1)} & (2) & (3) & (4) & \\multicolumn{1}{c}{(5)} & \\multicolumn{1}{c}{(6)} & \\multicolumn{1}{c}{(7)} & (8) & (9) \\\\ \\hline Q0738$+$313 \\dotfill & 0.635 & 0.0912 & 21.18 & ... & ... & $K>17.8$ & $>-18.8$ & 1,2 \\\\ PKS0952$+$179\\dotfill & 1.472 & 0.2390 & 21.32 & ... & ... & ... & ... & 3 \\\\ PKS1229$-$021\\dotfill & 1.038 & 0.3950 & 20.60 & 1.4 & $\\approx 5.2$ & $R=22.1$ & $\\approx -18.4$ & 4 \\\\ Q1209$+$107\\dotfill & 2.191 & 0.6295 & 20.48 & 1.6 & $\\approx 7.7$ & $R=21.6$ & $\\approx -20.5$ & 5 \\\\ PKS1622$+$23\\dotfill & 0.927 & 0.6563 & 20.36 & ... & ... & $R>24.5$ & $>-16.9$ & 6 \\\\ Q1328$+$307\\dotfill & 0.849 & 0.692 & 21.30 & 2.0 & $\\approx 10$ & $I=22.1$ & $\\approx -19.3$ & 5 \\\\ PKS0454$+$039\\dotfill & 1.345 & 0.8596 & 20.76 & 0.8 & $\\approx 4.3$ & $R=24.2$ & $\\approx -19.0$ & 4,5 \\\\ \\hline \\multicolumn{9}{l}{1.\\ Turnshek \\etal\\ (2001); 2.\\ Cohen (2001); 3.\\ Rao \\etal\\ (2003); 4.\\ Steidel \\etal\\ (1994); 5.\\ Le Brun \\etal\\ (1997); }\\\\ \\multicolumn{9}{l}{6.\\ Steidel \\etal\\ (1997).} \\\\ \\end{tabular}} \\end{center} \\end{table} \\end{tiny}  Identifying the optical counterpart of the DLAs has been a challenging task, because these galaxies are faint and located at small projected distances to the background, bright QSOs (as implied by the intrinsic high column density of the absorbers).  Le Brun \\etal\\ (1997) reported candidate absorbing galaxies toward seven QSO lines of sight with known DLAs based on high spatial resolution HST/WFPC2 images.  These galaxies exhibit a wide range in morphological types, from luminous spiral galaxies, to compact objects, and to low surface brightness galaxies.  In two fields, they also identify possible groups of galaxies that are associated with the DLAs.  In contrast, Rao \\etal\\ (2003) collected a sample of 14 candidate or confirmed DLA galaxies, and argued that low-luminosity dwarf galaxies dominate the DLA galaxy population with $\\approx 50$\\% of the DLAs originating in galaxies of luminosity $L\\le 0.25\\,L_*$.  Despite extensive searches in the past decade, the number of confirmed DLA galaxies remains relatively small.  Of all the 23 DLA systems known at $z\\le 1$,  14 have been identified with their host galaxies using either spectroscopic or photometric redshift techniques (see section 2 for a complete list).  Here I will first review current results from searching for $z<1$ DLA galaxies.  Then I will compare the optical properties of the absorbing galaxies, such as $B$-band luminosity, stellar content, and rotation velocity, with known properties of the absorbers, such as $N(\\hI)$ and metallicity. ",
        "conclusions": ""
    },
    "0410/astro-ph0410714_arXiv.txt": {
        "abstract": "We show that charged-particles decaying in the early Universe can induce a scale-dependent or `running' spectral index in the small-scale linear and nonlinear matter power spectrum and discuss examples of this effect in minimal supersymmetric models in which the lightest neutralino is a viable cold-dark-matter candidate. We find configurations in which the neutralino relic density is set by coannihilations with a long-lived stau, and the late decay of staus partially suppresses the linear matter power spectrum.  Nonlinear evolution on small scales then causes the modified linear power spectrum to evolve to a nonlinear power spectrum similar (but different in detail) to models parametrized by a constant running $\\alpha_{s}=d n_{s}/d{\\rm ln} k$ by redshifts of 2 to 4.  Thus, Lyman-$\\alpha$ forest observations, which probe the matter power spectrum at these redshifts, might not discriminate between the two effects.  However, a measurement of the angular power spectrum of primordial 21-cm radiation from redshift $z \\approx 30$--$200$ might distinguish between this charged-decay model and a primordial running spectral index. The direct production of a long-lived charged particle at future colliders is a dramatic prediction of this model. ",
        "introduction": "While recent cosmological observations provide convincing evidence that nonbaryonic dark matter exists \\cite{CMBconstraints}, we do not know the detailed particle properties of the dark matter, nor the particle spectrum of the dark sector.  There has been considerable phenomenological effort towards placing model-independent limits on the possible interactions of the lightest dark-matter particle (LDP)\\footnote{In supersymmetric models the LDP is the lightest supersymmetric particle (LSP), but we adopt this more general notation unless we are speaking about a specific supersymmetric model.} in an attempt to try and identify candidates within detailed particle theories or rule out particular candidate theories.  For instance, models with stable charged dark matter have been ruled out \\cite{Gould:1989gw}, while significant constraints have been made to dark-matter models with strong interactions \\cite{Starkman:1990nj}, self-interactions \\cite{Carlson:1992,Spergel:1999mh}, and a millicharge \\cite{Davidson:2000hf,Dubovsky:2003yn}.  Recently, it was shown that a neutral dark-matter particle with a relatively large electric or magnetic dipole moment remains a phenomenologically viable candidate \\cite{Sigurdson:2004zp}.  Concurrent with this phenomenological effort, theorists have taken to considering physics beyond the standard model in search of a consistent framework for viable dark-matter candidates. The leading candidates, those that produce the correct relic abundance and appear in minimal extensions of the standard model (often for independent reasons), are the axion~\\cite{axionreviews} and weakly interacting massive particles (WIMPs), such as the neutralino, the lightest mass eigenstate from the superposition of the supersymmetric partners of the $U(1)$ and $SU(2)$ neutral gauge bosons and of the neutral Higgs bosons~\\cite{Jungman:1995df,larsrev}. However, other viable candidates have also been considered recently, such as gravitinos or Kaluza-Klein gravitons produced through the late decay of WIMPs \\cite{Feng:2003xh,Feng:2004zu}.  These latter candidates are an interesting possibility because the constraints to the interactions of the LDP do not apply to the next-to-lightest dark-matter particle (NLDP), and the decay of the NLDP to the LDP at early times may produce interesting cosmological effects; for instance, the reprocessing of the light-element abundances formed during big-bang nucleosynthesis \\cite{Feng:2003uy,karsten}, or if the NLDP is charged, the suppression of the matter power spectrum on small scales and thus a reduction in the expected number of dwarf galaxies \\cite{Sigurdson:2003vy}.  In this paper we describe another effect charged NLDPs could have on the matter power spectrum.  If all of the present-day dark matter is produced through the late decay of charged NLDPs, then, as discussed in Ref.~\\cite{Sigurdson:2003vy}, the effect is to essentially cut off the matter power spectrum on scales that enter the horizon before the NLDP decays.  However, if only a fraction $f_{\\phi}$ of the present-day dark matter is produced through the late decay of charged NLDPs, the matter power spectrum is suppressed on small scales only by a factor $(1-f_{\\phi})^2$.  This induces a scale-dependent spectral index for wavenumbers that enter the horizon when the age of the Universe is equal to the lifetime of the charged particles. What we show below is that, for certain combinations of $f_{\\phi}$ and of the lifetime of the charged particle $\\tau$, this suppression modifies the nonlinear power spectrum in a way similar (but different in detail) to the effect of a constant $\\alpha_{s} \\equiv d n_{s}/d{\\rm ln}k \\neq 0$.  Although these effects are different, constraints based on observations that probe the nonlinear power spectrum at redshifts of 2 to 4, such as measurements of the Lyman-$\\alpha$ forest, might confuse a running index with the effect we describe here even if parametrized in terms of a constant $\\alpha_{s}$.  This has significant implications for the interpretation of the detection of a large running of the spectral index as a constraint on simple single-field inflationary models.  The detection of a unexpectedly large spectral running in future observations could instead be revealing properties of the dark-matter particle spectrum in conjunction with a more conventional model of inflation.  We note that the Sloan Digital Sky Survey Lyman-$\\alpha$ data \\cite{Seljak:2004xh} has significantly improved the limits to constant-$\\alpha_{s}$ models compared to previous measurements alone \\cite{CMBconstraints,Seljak:2003jg,Tegmark:2003ud}. A detailed study of the ($\\tau$,$f_{\\phi}$) parameter space using these and other cosmological data would also provide interesting limits to the models we discuss here.  While, even with future Lyman-$\\alpha$ data, it may be difficult to discriminate the effect of a constant running of the spectral index from a scale-dependent spectral index due to a charged NLDP, other observations may nevertheless discriminate between the two scenarios.  Future measurements of the power spectrum of neutral hydrogen through the 21cm-line might probe the linear matter power spectrum in exquisite detail over the redshift range $z \\approx 30 - 200$ at comoving scales less than 1 Mpc and perhaps as small as $0.01$ Mpc \\cite{Loeb:2003ya}; such a measurement could distinguish between the charged-particle decay scenario we describe here and other modifications to the primoridal power spectrum. If, as in some models we discuss below, the mass of these particles is in reach of future particle colliders the signature of this scenario would be spectacular and unmistakable---the production of very long-lived charged particles that slowly decay to stable dark matter.   Although we describe the cosmological side of our calculations in a model-independent manner, remarkably, there are configurations in the minimal supersymmetric extension of the standard model  (MSSM) with the right properties for the effect we discuss here. In particular, we find that if the LSP is a neutralino quasi-degenerate in mass with the lightest stau, we can naturally obtain, at the same time, LDPs providing the correct dark matter abundance $\\Omega_{\\chi}h^2 = 0.113$ \\cite{CMBconstraints} and NLDPs with the long lifetimes and the sizable densities in the early Universe needed in the proposed scenario. Such configurations arise even in minimal supersymmetric schemes, such as the minimal supergravity (mSUGRA) scenario~\\cite{msugra} and the minimal anomaly-mediated supersymmetry-breaking (mAMSB) model~\\cite{mamsb}. This implies that a detailed study of the ($\\tau$,$f_{\\phi}$) parameter space using current and future cosmological data may constrain regions of the MSSM parameter space that are otherwise viable.  Furthermore, we are able to make quantitative statements about testing the scenario we propose in future particle colliders or dark matter detection experiments.   The paper is organized as follows: We first review in Section~\\ref{sec:charged} how the standard calculation of linear perturbations in an expanding universe must be modified to account for the effects of a decaying charged species, calculate the linear matter power spectrum and discuss the constraints to this model from big bang nucleosynthesis (BBN) and the spectrum of the cosmic microwave background (CMB).  In Section~\\ref{sec:nonlinear} we briefly discuss how we estimate the nonlinear power spectrum from the linear power spectrum and present several examples.  In Section~\\ref{sec:21cm} we discuss how measurements of the angular power spectra of the primordial 21-cm radiation can be used to distinguish this effect from other modifications to  the primordial power spectrum.  In Section~\\ref{sec:partmodel} we describe how this scenario can be embedded in a particle physics model, concluding that the most appealing scheme is one where long lifetimes are obtained by considering nearly degenerate LDP and NLDP masses. In Section~\\ref{sec:lifetimes} we compute the lifetimes of charged next-to-lightest supersymmetric particles (NLSPs) decaying into neutralino LSPs, and show that, in the MSSM, the role of a NLDP with a long lifetime can be played by a stau only. In Section~\\ref{sec:fraction} we estimate what fraction of charged to neutral dark matter is expected in this case, while in Section~\\ref{sec:msugra} we  describe how consistent realizations of this scenario can be found within the parameter space of mSUGRA and mAMSB.  Finally, in Section~\\ref{sec:colliders} we discuss the expected signatures of this scenario at future particle colliders, such as the large hadron collider (LHC), and prospects for detection in experiments searching for WIMP dark matter.  We conclude with a brief summary of our results in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have examined a scenario in which a fraction $f_{\\phi}$ of the cold-dark-matter component in the Universe is generated in the decay of long-lived charged particles and shown that a scale-dependent (or `running') spectral index is induced. The power spectrum on scales smaller than the horizon size when the age of the Universe is equal to the charged-particle lifetime gets suppressed by a factor $(1-f_{\\phi})^2$. Such a feature might be singled out unambiguously by future measurements of the power spectrum of neutral hydrogen through the 21-cm line, obtaining direct information on the charged particle lifetime and $f_{\\phi}$.  On the contrary, current and future tests for departures from a scale-invariant power spectrum based on Lyman-$\\alpha$ data may fail to uniquely identify the scenario we propose. In fact, we have estimated the modifications to the non-linear power spectrum at the redshifts and wavenumbers currently probed by Lyman-$\\alpha$ forest data, and shown that these resemble (but are different in detail) those in models with a constant running of the spectral index $\\alpha_{s}$.  We expect, based on this resemblance, models with $f_{\\phi}$ in the range $1/2 - 1/7$ (as predicted in some explicit models we have constructed) to be compatible with current cosmological data for lifetimes in the range $1 - 20$~yr. We have also verified that constraints from the primordial light-element abundances and distortions to the CMB spectrum are not violated.  From the particle-physics point of view, we have shown that the proposed scenario fits nicely in a picture in which the lightest neutralino in SUSY extensions of the standard model appears as the cold-dark-matter candidate, and a stau nearly degenerate in mass with the neutralino as the long-lived charged counterpart. A small mass splitting forces the stau to be quasi-stable, since the phase space allowed in the its decay process gets sharply reduced. At the same time, it implies that neutralino and stau are strongly linked in the process of thermal decoupling, with the charged species playing the major role.  Owing to these coannihilation effects, the current neutralino thermal relic abundance is compatible with the value inferred from cosmological observatations and, at early times, the stau thermal relic component is at the correct level.  We have described several explicit realizations of this idea in minimal supersymmetric frameworks, including the minimal supergravity scenario, namely the supersymmetric extension of the standard model with smallest possible parameter space. We have pointed out that charged dark matter fraction from 1/2 to 1/7 (or even lower) can be obtained and that stau lifetimes larger than 1~yr are feasible. We have also shown that some of the models we have considered may be detected in future WIMP direct searches, and discussed the prospects of testing the most dramatic feature of the model we propose, i.e. the production of long-lived staus at future high energy particle colliders."
    },
    "0410/astro-ph0410008_arXiv.txt": {
        "abstract": "We outline the results from a FUSE Team program designed to characterize O~VI absorption in the disk of the Milky Way.  We find that O~VI absorption occurs throughout most of the Galactic plane, at least out to several kpc from the Sun, and that it is distributed smoothly enough for the column density to decline with height above the disk and with distance in the plane. However, the O~VI absorbing gas is clumpy, and moves at peculiar velocities relative to that expected from Galactic rotation. We conclude that the observed absorption is likely to be a direct indicator of the structures formed when violent, dynamical processes heat the ISM, such as blowout from multiple supernovae events. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410522_arXiv.txt": {
        "abstract": "We analyze multi-band photometry of a subsample of low mass stars in the associations Ori OB1a and 1b discovered during the CIDA Orion Variability Survey, which have ages of 7 - 10 Myr and 3 - 5 Myr, respectively. We obtained $UBVR_cI_c$ photometry at Mt Hopkins for 6 Classical T Tauri stars (CTTS) and 26 Weak T Tauri stars (WTTS) in Ori OB1a, and for 21 CTTS and 2 WTTS in Ori OB1b. We also obtained $L$ band photometry for 14 CTTS at Mt. Hopkins, and 10$\\mu$m and 18$\\mu$m photometry with OSCIR at Gemini for 6 CTTS; of these, all 6 were detected at 10$\\mu$m while only one was detected at 18$\\mu$m. We estimate mass accretion rates from the excess luminosity at $U$, and find that they are consistent with determinations for a number of other associations, with or without high mass star formation. The observed decrease of mass accretion rate with age is qualitatively consistent with predictions of viscous evolution of accretion disks, although other factors can also play a role in slowing accretion rates. We compare the excesses over photospheric fluxes in $H-K$, and $K-L$, and $K-N$ with the younger sample of Taurus and find an overall decrease of disk emission from Taurus to Ori OB1b to Ori OB1a. This decrease implies that significant grain growth and settling towards the midplane has taken place in the inner disks of Ori OB1. We compare the SED of the star detected at both 10$\\mu$m and 18$\\mu$m with disk models for similar stellar and accretion parameters. We find that the low $\\le 18 \\mu$m fluxes of this Ori OB1b star cannot be due to the smaller disk radius expected from viscous evolution in the presence of the FUV radiation fields from the OB stars in the association. Instead, we find that the disk of this star is essentially a flat disk, with little if any flaring, indicating a a significant degree of dust settling towards the midplane, as expected from dust evolution in protoplanetary disks. ",
        "introduction": "Most star formation studies are based on a few nearby associations without massive stars, as Taurus and Chameleon or in highly populated clusters with significant high mass formation, like the Orion Nebula Cluster (ONC). In recent years, several teams have begun systematic studies of the populations of the OB associations. These associations cover large volumes in the sky and include populations covering the entire mass range. OB associations have been subject of studies with Hipparcos (Brown et al. 1997; de Zeeuw \\etal 1999), which has helped pinpoint their distances and mean ages. In a number of associations, the populations display a gradient of ages, from $\\sim$ 1 Myr to a few $\\sim$ 10 Myr, which has been interpreted as the result of triggered star formation (Blaauw 1964). This range of ages makes the OB associations the most suitable laboratory for studies of protoplanetary disk evolution, because it encompasses the time span in which giant planets are expected to form (Pollack et al. 1996; Alibert et al. 2004), and significant solid evolution is expected to take place, as indicated by meteoritic evidence (Podoseck \\& Cassen 1996; Wood 2004).  We are carrying out a photometric and spectroscopic survey of $\\sim $ 128 degrees in the Orion OB1 association in order to identify the low mass population.  Initial results of this survey have been presented by Brice\\~no et al. (2001, 2004 = Paper I). In Paper I, we identified 197 new low mass members in the subassociations Ori OB1a and 1b, spanning the mass range from $\\sim$ 0.2 to 1.4 $\\msun$. We found ages of $\\sim$ 3 - 5 Myr for Ori OB1b and $\\sim$ 7 - 10 Myr for Ori OB1a, confirming age determinations based on the OB stars (Blaauw 1964;  Warren \\& Hesser 1977, 1978; Brown et al. 1994). We also found that Ori OB1b could be identified with a population associated with a ring of low density gas and dust, probably formed by a supernova event, with most of the Classical T Tauri stars (i.e., stars accreting from circumstellar disks; CTTS) near the ring, similar to the case of $\\lambda$ Ori (Dolan \\& Mathieu 1999, 2001, 2002). In contrast, Ori OB1a is in a region devoid of gas and dust, in agreement with its older age. We also found that the number of CTTS decreases sharply between Ori OB1b and 1a, indicating the few disks remain by $\\sim$ 10 Myr, in agreement with findings from young clusters, some containing high mass stars, and low mass associations as TW Hya (Haisch, Lada, \\& Lada 2003; Hillenbrand, Carpenter, and Meyer 2004, in preparation; Sicilia-Aguilar et al. 2004, in preparation; Muzerolle et al. 2000, 2001).  In this work, we present multiwavelength photometry of a subset of the stars in Paper I, including stars in Ori OB1a, 1b and 1c. In particular, we have obtained $U$ photometry of a significant number of CTTS in these subassociations, from which we have determined mass accretion rates for comparison with disk evolution theories. We have also obtained $L$ band and mid-infrared (IR) measurements at 10$\\mu$m and 18$\\mu$m of a small subset of stars in Ori OB1, which has allowed us to make assessments of the state of the dust in their disks in comparison with younger populations. In \\S 2 we present the photometry, and in \\S 3, we examine the properties of the disks in Ori OB1a and 1b in comparison with the Taurus populations. We discuss the implications of these results for disk evolution in \\S \\ref{sec_discussion}. ",
        "conclusions": "\\label{sec_concl}  We have detected clear signatures of disk evolution in samples of accreting stars in the Ori OB1a and 1b associations, with ages 7 - 10 Myr and 3 - 5 Myr, respectively. The mass accretion rates of these stars are consistent with viscous evolution of accretion disks. In addition, we find a significant overall decrease of infrared emission with age. Comparison with disk models indicates that grain growth and settling, processes expected from dust evolution,  are responsible for this decrease. Future observations of larger samples of Ori OB1 stars with SPITZER will provide more detail on the nature of the dust evolution process.  Acknowledgments. We thank the referee for many useful suggestions. This work was supported by NSF grant AST-9987367 and NASA grant NAG5-10545. PD acknowledges grants from PAPIIT, DGAPA (UNAM) and CONACyT, M\\'exico. This research 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 paper is based on observations obtained with the mid-infrared camera OSCIR, developed by the University of Florida with support from the National Aeronautics and Space Administration, and operated jointly by Gemini and the University of Florida Infrared Astrophysics Group. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France."
    },
    "0410/astro-ph0410378_arXiv.txt": {
        "abstract": "We consider a contribution of microlensing to the X-ray variability of high-redshifted QSOs. Cosmologically distributed gravitational microlenses could be localized in galaxies (or even in bulge or halo of  gravitational macrolenses) or could be distributed in a uniform way. We have analyzed both cases of such distributions. We found that the optical depth for gravitational microlensing caused by cosmologically distributed deflectors could be significant and could reach $10^{-2} - 0.1$ at $z\\sim 2$.  This means that cosmologically distributed deflectors may contribute significantlly  to the X-ray variability of high-redshifted QSOs ($z>2$). Considering that the upper limit of the optical depth ($\\tau\\sim 0.1$) corresponds to the case where dark matter forms cosmologically distributed deflectors,   observations of the X-ray variations of unlensed QSOs can be used for the estimation of the dark matter fraction of microlenses. ",
        "introduction": "X-ray flux variability has long been known to be a common property of active galactic nuclei (AGNs), e.g. Ariel 5 and HEAO 1 first revealed long-term (days to years) variability in AGNs  and by uninterrupted observations of EXOSAT  rapid (thousands of seconds) variability was also established as common in these sources (see, for example reviews by \\cite{Mush93,Ulrich97} and references therein). X-ray flux variations are observed on timescales from $\\sim$1000 s to years, and amplitude variations of up to an order of magnitude are in the $\\sim$ 0.1 - 10 keV spectral band. Recently, \\cite{Man02} analyzed the variability of a sample of 156 radio-quiet quasars taken from the ROSAT archive, considering the trends in variability of the amplitude with luminosity and with redshift. They found that there was evidences for a growth in AGN X-ray variability amplitude towards high redshift ($z$) in the sense that AGNs of the same X-ray luminosity were more variable at $z>2$. They explained the $\\sigma$ {\\it vs.} $z$ trend assuming that the high-redshifted AGNs accreted at a larger fraction of the Eddington limit than the low-redshifted ones.  On the other hand, the contribution of microlensing to AGN variability was considered in many papers (see e.g. \\cite{Hawk93,Hawk02, Wamb01,Wamb01b,Zakh97}, and references therein). Moreover, recently X-ray microlensing of AGN has been considered \\citep{Popov01a,Tak01,Chart02a,Popovic03,Popovic03b,Dai03}. Taking into account that the X-rays of AGNs are generated in the innermost and very compact region of an accretion disc, the X-ray radiation in the continuum as well as in a line can be strongly affected by microlensing \\citep{Popovic03}.\\footnote{Simulations of X-ray line profiles are presented in a number of papers, see, for example, \\cite{Zak_rep02,Zak_rep02a,Zak_rep02_xeus,Zak_rep02_Gamma,Zak_rep03_Lom,Zak_rep03_ASR,Zak_rep03_Su,Zak_rep03_Sakh} and references therein, in particular \\cite{ZKLR02} showed that an information about magnetic filed may be extracted from X-ray line shape analysis; \\cite{Zak_rep03_AA} discussed signatures of X-ray line shapes for highly inclined accretion disks, \\cite{ZMB03} calculated shapes of spectral lines for non-flat accretion flows.} Recent observations of three lens systems seem to support this idea \\citep{Osh01,Chart02a,Cha04,Dai03}.  \\cite{Popovic03,Popovic03b}  showed that objects in a foreground galaxy with very small masses can cause strong changes in the X-ray line profile.  This  fact may indicate  that the observational probability of X-ray variation due to microlensing events is higher  than in the UV and optical radiation of AGNs. It is connected with the fact that typical sizes of X-ray emission regions are much smaller than typical sizes of those producing optical and UV bands. Typical optical and UV emission region sizes could be comparable or even larger than Einstein radii of microlenses and therefore microlenses magnify  only a small part of the region emitting in the optical or UV band (see e.g. \\cite{Pop01b,Aba02}, for UV and optical spectral line region). This is reason that it could be a very tiny effect from an observer point of view, in spite of  this fact recently \\cite{Richards04} observed microlensing of the C IV line in SDSS J1004+4112.   Microlenses in quasar bulge/halo give a small contribution into optical depth, therefore it would be reasonable to evaluate a contribution  from cosmological distribution of microlenses \\citep{ZPJ03,ZPJ04}. ",
        "conclusions": "In Fig.  we show the optical depth of cosmologically distributed microlenses assuming three cases of cosmoligically distributed microlenses: i) small fraction of baryonic matter (25 \\%) forms microlenses ($\\Omega_L = 0.01$); ii) almost all baryonic matter forms microlenses ($\\Omega_L = 0.05$); iii) about 30\\% of non-baryonic (dark) matter forms microlenses ($\\Omega_L = 0.1$).  As  was mentioned earlier by \\cite{Popovic03,Popovic03b} the probability of microlensing by stars or other compact objects in halos and bulges of quasars is very low (about $10^{-4} - 10^{-3}$). However, as one can see from Fig.~1, for cosmologically distributed microlenses it could reach $10^{-2} - 0.1$ at $z\\sim 2$. The upper limit $\\tau \\sim 0.1$ corresponds to the case where compact dark matter forms cosmologically distributed microlenses. As one can see from Fig. 1, in this case the optical depth for the considered value of $\\Omega_0$  is around 0.1  for $z > 2$. This indicates that   such a phenomenon could be  observed frequently, but only for distant sources ($z \\sim 2$).  To investigate  distortions of spectral line shapes due to microlensing \\citep{Popovic03,Popovic03b} the most real candidates are multiply imaged quasars. However, these cases the simple point-like microlens model may  not be very good approximation \\citep{Wamb01,Wamb01b} and one should use a numerical approach, such as the MICROLENS ray tracing program, developed by J. Wambsganss % or  some analytical approach for magnification near caustic curves like folds \\citep{Schneider92a,Fluke99} or near singular caustic points like cusps \\citep{Schneider92,Mao92,Zakharov95,Zakharov97,Zakharov99} as was realized by \\cite{Yonehara01}.  If  we believe in the  observational arguments of \\cite{Hawk02} that the variability of a significant fraction of distant quasars is caused by microlensing,  the analysis of the properties of X-ray line shapes due to microlensing \\citep{Popovic03} is a powerful tool to confirm or rule out Hawkins' (2002) conclusions.  As it was mentioned, the probability that the shape of the  Fe $K\\alpha$ line is distorted (or amplified) is highest in gravitationally lensed systems. Actually, this phenomena was discovered by \\cite{Osh01,Dai03,Chart02a,Cha02b,Cha04} who found evidencees for such an effect for   QSO H1413+117 (the Cloverleaf, $z=2.56$), QSO~2237+0305 (the Einstein Cross, $z=1.695$), MG J0414+0534 ($z=2.64$) and possibly for BAL QSO~08279+5255 ($ z=3.91$). One could say that it is natural that the discovery of X-ray microlensing was made for this quasar, since the Einstein Cross QSO~2237+0305 is the most \"popular\" object to search for microlensing, because the first cosmological microlensing  phenomenon was found by \\cite{Irwin89} in this object and  several groups have been monitoring the quasar QSO~2237+0305 to find evidence for microlensing.  Microlensing has been suggested for the quasar MG J0414+0534 \\citep{Angonin99} and for the quasar QSO H1413+117 \\citep{Remy96}. Therefore, in future a chance may be to find X-ray microlensing for other gravitationally lensed systems that have  signatures of microlensing in the optical and radio bands. Moreover, considering the sizes of the sources of X-ray radiation, the variability in the X-ray range during microlensing event should be more prominent than in the optical and UV. { Consequently, gravitational microlensing in the X-ray band is a powerful tool for dark matter investigations, as the upper limit of optical depth ($\\tau\\sim 0.1$) corresponds to the case where dark matter forms cosmologically distributed deflectors."
    },
    "0410/astro-ph0410652_arXiv.txt": {
        "abstract": "\\label{abs}  Strange stars with low masses are suggested to exist in reality, the origin of which could be via accretion-induced collapse (AIC) of white dwarfs. Such a strange star is likely bare, and would thus spin very fast, even to a period of $<0.1$ ms. Strange stars with low masses may differ from those with solar-masses in various astrophysical appearances. Observations to test this ``low-mass'' idea are proposed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410464_arXiv.txt": {
        "abstract": "Mass loss from massive stars ($\\ga 8 \\msun$) can result in the formation of circumstellar wind blown cavities surrounding the star, bordered by a thin, dense, cold shell. When the star explodes as a core-collapse supernova (SN), the resulting shock wave will interact with this modified medium around the star, rather than the interstellar medium. In this work we first explore the nature of the circumstellar medium around massive stars in various evolutionary stages. This is followed by a study of the evolution of SNe within these wind-blown bubbles. The evolution depends primarily on a single parameter $\\Lambda$, the ratio of the mass of the dense shell to that of the ejected material. We investigate the evolution for different values of this parameter. We also plot approximate X-ray surface brightness plots from the simulations. For very small values $\\Lambda << 1$ the effect of the shell is negligible, as one would expect. Values of $\\Lambda \\la 1$ affect the SN evolution, but the SN 'forgets' about the existence of the shell in about 10 doubling times or so. The remnant density profile changes, and consequently the X-ray emission from the remnant will also change. The initial X-ray luminosity of the remnant is quite low, but interaction of the shock wave with the dense circumstellar shell can increase the luminosity by 2-3 orders of magnitude.  As the reflected shock begins to move inwards, X-ray images will show the presence of a double-shelled structure. Larger values result in more SN energy being expended to the shell. The resulting reflected shock moves quickly back to the origin, and the ejecta are thermalized rapidly. The evolution of the remnant is speeded up, and the entire remnant may appear bright in X-rays. If $\\Lambda >> 1$ then a substantial amount of energy may be expended in the shell. In the extreme case the SN may go directly from the free-expansion to the adiabatic stage, bypassing the Sedov stage. Our results show that in many cases the SN remnant spends a significant amount of time within the bubble. The low density within the bubble can delay the onset of the Sedov stage, and may end up reducing the amount of time spent in the Sedov stage. The complicated density profile within the bubble makes it difficult to infer the mass-loss properties of the pre-SN star by studying the evolution of the resulting supernova remnant. ",
        "introduction": "Core collapse supernovae (SNe) arise from the collapse of massive stars with mass M $\\ga 8 \\msun$. As these stars evolve along the main sequence and beyond it, they lose a considerable amount of mass in the form of winds \\citep{kp00}. The collective action of the winds from different evolutionary stages may sweep up the material in the ambient medium to form a wind-blown cavity, bordered by a thin, dense, cold shell of material, which contains most of the mass of the swept-up medium. When the star explodes as a SN, the shock wave resulting from the explosion evolves in this highly modified circumstellar medium, and not in the pristine interstellar medium \\citep{m88}. The evolution of the supernova remnant is therefore very different from the classical evolution in a constant density interstellar medium, where the remnant successively advances through the phases of free-expansion, adiabatic or Sedov stage, radiative stage and dispersal into the surrounding medium \\citep[see for example][]{w72}.  Evidence has built up over the last few years that the medium around a pre-SN star is considerably modified by the action of winds, radiation and stellar outbursts. This is highlighted by the famous Hubble Space Telescope picture of the three-ring nebula surrounding the SN 1987A. The combined action of winds and radiation in sculpting this medium is undeniable \\citep{bl93, cd95}. Other known cases of SNe exploding in wind-driven shells are less spectacular. There are indications that the Cygnus Loop is the remnant of an old supernova explosion that went off in a wind-blown bubble \\citep{l97}. Circumstellar interaction models were earlier proposed for N49 \\citep{sdkw85} and N132D \\citep{h87}. \\citet{bsbs96} have modeled the X-ray emission from Cas A as arising from a SN explosion within a wind-blown cavity.  Examples of wind-blown bubbles around massive stars are far more numerous \\citep{c03, cec03}. Ring nebulae are seen around main sequence O and B stars, OfPe/WN9 stars \\citep{n99}, and many Wolf-Rayet stars in general. Bipolar nebulae are observed around many LBV stars \\citep{w01}, the best-known and studied one being $\\eta$ Carinae \\citep{dh97}. Circumstellar nebulae are observed around red supergiant stars such as VY CMa \\citep{shd01, hds02}, and around blue supergiant stars such as NGC6164-5 \\citep{lck87}. The progenitor of SN 1987A is thought to be a blue-supergiant (BSG). A similar structure is seen around the star Sher 25, suggesting that it may be an analogue of SN 1987A \\citep{c03}.  The evolution of remnants within wind-blown cavities has merited some attention in the past. The problem was discussed numerically by \\citet{cd89}. It was placed on a firm theoretical footing with analytic calculations by \\citet{cl89}. Early numerical calculations are described in a series of papers by Tenorio-Tagle and collaborators, wherein they discussed the 1-dimensional evolution \\citep{tbf90}, the 2-dimensional evolution \\citep{trfb91}, and off-center explosions within pre-existent bubbles \\citep{rtfb93}. A review is presented in \\citep{ftbr91}. \\citet{bl93} and \\citet{ma95} carried out hydrodynamical simulations of the nebula surrounding SN 1987A, and \\citet{cd95} incorporated the effects of the ionizing radiation from the pre-SN star in order to explain the linearly increasing X-ray and radio emission from SN 1987A, as well as the considerable slowing down of the shock front. \\citet{bsbs96} modeled the X-ray emission from Cas A under the assumption that the SN expanded in a wind-blown bubble.  In recent years the availability of high spectral and spatial resolution data, both space and ground-based, has added considerably to our knowledge of circumstellar (CS) wind-blown bubbles surrounding massive stars. Observations from HST, Chandra, XMM, and now SIRTF have enriched our knowledge of the structure, formation and evolution of wind-blown nebulae. Simultaneously, advances in understanding the structure of the ejected material in supernovae \\citep[for e.g.][]{tm99, mm99} have further increased our understanding of supernova evolution. The worldwide effort gone into studying SN 1987A has enabled us to explore in real time a SN expanding within a wind-blown bubble on its way to becoming a SNR.  The recent confirmation that gamma-ray bursts arise from very massive stars led to the realization that the shock waves in this case may also evolve in the pre-outburst stellar winds from the progenitor, rather than in the ISM. There is some speculation that gamma-ray burst afterglows may be explained in this manner \\citep{clf04}. The extreme interest in gamma-ray bursts as cosmological phenomena has given further impetus to the study of supernova shock waves in stellar winds.  Given the considerable increase in our observational knowledge, accompanied by improved theoretical understanding of SN evolution and SN progenitors, we wish to revisit the evolution of SNRs in circumstellar wind-blown bubbles. We have planned a comprehensive effort to explore the dynamics, kinematics and radiation signatures of SNe expanding in CS wind-blown bubbles. In this first paper we will explore the various aspects of the interaction, define the parameters on which the evolution depends, and carry out a broad parameter survey. Our initial models will be idealized in order to understand the effects of various parameters. In subsequent papers we will take more realistic stellar evolution calculations into account, which provide the wind-properties from a star at each stage. These will help us to explore the structure and kinematics of more complicated objects. Finally we will apply the knowledge gained to investigations of specific astrophysical objects, especially SN 1987A.  The remainder of this paper is organized as follows. In \\S \\ref{sec:massloss} we briefly discuss mass-loss from massive stars. Hot and cool stars are considered separately.  In \\S \\ref{sec:wbb} we describe the basic structure of CS wind-blown bubbles, and in \\S \\ref{sec:snejecta} we outline the density structure of the ejected SN material. \\S \\ref{sec:ts} describes a simple analytic solution that helps to understand the basic ideas, and sets the stage for the numerical calculations. A very brief overview of the code is given in \\S \\ref{sec:numinfo}.  In \\S \\ref{sec:snbub} we study the basic features of the interaction of SN shock waves with circumstellar wind-blown bubbles, using spherically symmetric simulations. Two specific cases are chosen to illustrate the basic aspects of the interaction. X-ray brightness profiles and luminosity plots are presented.  \\S \\ref{sec:disc} discusses the implications of our research for the onset and growth of the Sedov stage. Finally, \\S \\ref{sec:summary} summarizes the basic results, and outlines future work. ",
        "conclusions": "\\label{sec:disc}  It is clear that the medium around the progenitor star plays a large role in shaping the further evolution of the remnant. Depending on the nature of the progenitor star, the density of the medium into which the SN shock expands may vary considerably. RSG progenitors will have a region of very high density, surrounded by a low-density bubble formed by the main-sequence star, bordered by a dense shell. WR progenitors will have a low density wind-blown bubble surrounding the star, bordered by a dense circumstellar shell. The interior of the shell may show density variations depending on the various stages of evolution prior to the SN stage.  As we have seen in the case of SN 1987A, BSG progenitors will form a cavity with density somewhat midway between the above two cases. Factors such as ionization, rotation, binary evolution and magnetic fields could all affect these results.  Quite often the SNR spends a not insignificant fraction of its life within the bubble. It is then of interest whether the SN reaches a Sedov stage while inside the bubble or not. Note that, depending on the density profile, the swept-up mass must be several times that of the ejecta mass before the remnant enters the Sedov stage \\citep[see for example][]{dc98}. Whether the remnant sweeps up enough material is a function of the ejected mass and the bubble interior density. The bubble density depends on the nature of the exploding star, which is a function of its initial mass and mass-loss rate.  If the density into which the SN shock evolves is very low, then the remnant will not sweep up much material, and will remain in the free-expansion phase for a much longer time than it would in the ISM. This is probably the case with WR progenitors.  For remnants expanding in a low-density bubble it is unlikely that they will reach the Sedov stage before they collide with the dense shell. As mentioned in \\S \\ref{sec:largel} in cases where the mass of the shell significantly exceeds that of the ejecta, the remnant may evolve from the free-expansion stage directly to the radiative stage, completely bypassing the Sedov stage.  If the progenitor is a RSG, then the surrounding density could be much higher than that of the ISM. The remnant will sweep up a large mass of material in a short time. Whether it evolves to the Sedov stage then depends on the mass of the ejected material and the size of the RSG wind-region, which depends primarily on the time spent in the RSG stage. It is likely that the dense RSG wind will be surrounded by a low density main-sequence bubble, whose density could be a few orders of magnitude lower than that of the RSG wind. If the remnant has already reached the Sedov stage then the evolution will continue as a Sedov-type blast wave in a low density medium. If the remnant has not reached the adiabatic stage then it is unlikely that it will do so in the much lower density medium, and it will probably continue to be in free-expansion until the shock wave collides with the circumstellar shell.  The ejecta mass depends on the initial (zero-age main sequence) mass of the star, the mass of the progenitor star just prior to the explosion, and the nature of the compact stellar remnant. All of these factors depend on the metallicity. According to recent work by \\citet{hfw03} and \\citet{whw02}, stars in the mass range between 10 and 25 $\\msun$ and solar metallicity will leave behind a neutron star, and will have an ejecta mass about 10 solar masses or larger. However much more massive stars can form larger mass black holes, and the ejecta mass is reduced to just a few solar masses. A very broad conclusion that can be drawn is that RSG stars will result in a large ejecta mass, whereas WR stars will have a much lower ejecta mass (this is not true for stars between 25 and 35 $\\msun$). As we discussed before SNe from WR progenitors will evolve in a much lower density medium that those from RSG progenitors. Thus larger (smaller) ejecta mass SNe will evolve in denser (low density) winds. A general conclusion is that it will be difficult in most cases for the SNR to sweep up enough material to reach the Sedov stage before colliding with the dense shell.  The situation is opposite for SNe whose progenitor stars lie between about 25 and 35 $\\msun$, which will explode as RSGs but leave behind a black hole with mass much larger than that of a neutron star, and a very small mass of ejected material. These stars will most likely be surrounded by a higher-density medium, and therefore will tend to reach the Sedov stage quickly. When this Sedov remnant collides with the dense shell, the reflected shock will traverse an almost evacuated medium and move very quickly back to the center.  In higher mass stars the situation will be further complicated. The WR stage is a highly evolved stage of a high-mass star greater than about 35 $\\msun$. It is likely to be preceded by a RSG stage or perhaps even a Luminous Blue Variable (LBV) stage, if the initial mass of the star is much larger \\citep{lhl94}. LBVs can have very large mass-loss rates, and therefore very high-density winds, even compared to RSGs. What this implies is that the WR stage will generally follow a phase of very high-density winds. The WR phase itself has a high mass-loss rate and wind velocity. The momentum of the WR wind is generally going to be much larger than the RSG wind, and may even exceed that of the LBV wind, although this depends on the specific parameters. If the WR wind was preceded by a low-velocity RSG wind, the WR wind momentum will carry the RSG wind material along with it. This combined material may collide with the main sequence shell and reflect back, forming more complicated structures than described herein. The low density WR wind may end up containing higher-density material from the RSG stage. The density profile will not be representative of the wind-properties of the pre-SN progenitor star.  It is apparent that the medium into which a core-collapse SN evolves cannot be easily categorized into either a constant density medium such as the ISM, or a wind with a density $\\rho \\propto r^{-2}$. A wind-blown bubble formed by two-wind interaction includes both elements, a density that decreases as $r^{-2}$ followed by a more-or less constant density region between the wind termination shock and the contact discontinuity. However in stars larger than 35 $\\msun$ mass-loss in the various evolutionary stages preceding the SN explosion can give rise to very complicated density profiles beyond the wind-termination shock. Close in to the star the density generally follows a 1/r$^2$ law resulting from the last phase of evolution, but further beyond the density may vary considerably depending on the evolutionary sequence. These more realistic structures will be the subject of future papers.    In this paper we have studied the evolution of supernova shock waves in the circumstellar structures formed by pre-supernova mass loss. The evolution depends mainly on one parameter $\\Lambda$, the ratio of the mass of the dense surrounding shell to the mass of the ejected material. We have explored the changes in the evolution as the value of this parameter increases, and found that the evolution is significantly different from that for an explosion in the interstellar medium. In particular as the value of $\\Lambda$ increases we find that the energy imparted to the shell increases, the evolution of the remnant is speeded up, and the appearance of the remnant may change.  In contrast to the work of \\citet[][hereafter TBFR90]{tbf90}, we have used a much steeper density profile for the outer ejecta and a higher ejecta mass. We have also explored more thoroughly the region around $\\Lambda \\sim 1$, using two different values of $\\Lambda=0.14$ and $\\Lambda=3.7$. Simulations were also carried out, although they are not explicitly reported herein, with $\\Lambda >> 1$. Our results are consistent with the work of TBFR90 for $\\Lambda >> 1$, as would be expected. However, unlike TBFR90, we have explored two different cases around the value $\\Lambda \\sim 1$ and find considerable difference between the two cases. These differences can be noted in the evolution as well as the X-ray appearance. In one case a double-shelled structure is visible for a considerable amount of time after the shock shell interaction. In the other case the double shell is mostly absent, but if the interior density is high the entire remnant may be visible in X-rays. However differences are not so easily apparent in observations of the forward or outer shock front. Thus we caution that inferring the properties simply by observing the outer shock parameters can often lead to erroneous conclusions.  We have presented the evolution of the expansion parameter in each case. It is worth pointing out again that the expansion parameter varies significantly over the evolution, unlike in the case of explosion within a constant density medium or a wind. This is very important, because it implies that any instantaneous value of the expansion parameter is not representative of the evolutionary stage of the remnant, and will lead to an incorrect determination of properties. And in particular, assuming that the remnant is in the Sedov stage for computing the remnant parameters could lead to large errors.  The initial evolution of the bubble, within a low-density medium, can result in a large amount of time spent in the free-expansion phase, as the swept-up mass remains quite small. The presence of the circumstellar shell on the other hand, results in significant deceleration of the SN shock wave, and decreases the time spent in the Sedov stage as compared to the radiative stage. The net result is that the Sedov stage could be significantly shortened, and in some cases (see \\S \\ref{sec:largel}) may be entirely by-passed. It is plausible however that the forward shock, after emerging from the CS shell, may reach the Sedov stage. It is also quite likely though that given the deceleration of the shock front, which we have seen will occur even when the shell mass is about 10\\% of the ejecta mass, that the remnant may still reach the radiative stage much faster than in the absence of the shell.  Our X-ray profiles show that for $\\Lambda \\la 1$, a double-shelled structure is seen in X-rays, with emission coming mainly from shocked shell and from behind the reverse shock. It is possible that such a situation is present in Cas A. At later times when the outer shock has separated significantly from the shell, the emission from the outer shock begins to predominate. In the case with larger $\\Lambda$ the emission from the shocked shell always seems to predominate. It is clear then that as the remnant evolves and the outer shock separates, the X-ray emission from the shocked shell will appear to come from deep within say the radio emission arising from behind the outer shock. Precisely such a situation has been surmised for the remnant W44 \\citep{kh95}.   In this paper we have not considered multi-dimensional effects such as hydrodynamical instabilities and turbulence. These aspects will be dealt with in upcoming papers. Thin shells are known to be unstable, due to many instabilities ranging from the well-known Rayleigh-Taylor and Kelvin-Helmholtz type, to the range of instabilities lumped under the heading of Vishniac-type instabilities.  These instabilities may cause the shells to fragment, and lead to hydrodynamic behavior that has not been explored herein. Furthermore, variations in pressure can lead to turbulent behavior within the interior of the dense shells. Also of significance is that, as the reflected shock in Case 1 moves back to the center, there is a region behind the reflected shock that is Rayleigh-Taylor unstable. This instability is of importance because it occurs in the interior of the remnant, far away from the dense shell and forward shock, and can result in ejecta clumps and filaments in the inner regions of the remnant.  Finally, the real test of these results will be to compare and validate them with observed SNRs. We are indeed fortunate to be able to closely witness the formation and evolution of a SN which occurred within a wind-blown bubble, SN 1987A. The enormous body of observational data that has been collected should enable us not only to study this SN in the utmost detail, but by extension should provide considerable information on the general theory of SN evolution in wind-blown cavities. Of course SN 1987A has its own share of complexities, and a thorough understanding of the circumstellar medium requires 3-dimensional radiation hydrodynamics modeling, which has not yet been carried out. Specific comparisons between our models and observed remnants will be described in later papers in this series."
    },
    "0410/astro-ph0410187_arXiv.txt": {
        "abstract": "A straightforward formalism to evaluate the impact of Type Ia Supernovae (SNIa) on large scale astrophysical issues is presented, together with analytical formulations for the SNIa rate following an instantaneous burst of Star Formation (SF), for the variety of SNIa progenitors. Some applications of the parametrized formalism are illustrated. The observations seem to favour the Double Degenerate (DD) systems as SNIa precursors. ",
        "introduction": "Type Ia supernovae are an important source of Iron and energy to the interstellar medium, and for this reason the evolution of their rate over cosmic times has profound implications on a variety of astrophysical issues, including the chemical evolution of galaxies, the Fe content of clusters of galaxies, the evolution of the gas flows in Ellipticals (Es). For example, the delayed release of Fe from SNIa with respect to the prompt release of $\\alpha$ elements from Type II supernovae implies that the shorter the formation timescale of a stellar system, the higher the $\\alpha$/Fe abundance ratio recorded in its stars. Following this argument, short formation timescales have been inferred for E galaxies from the supersolar $\\alpha$ to Fe ratios traced by the Mg and Fe indices in their spectra (e.g. Matteucci 1994). As another example, the interplay between the rate of mass return from evolved stars in Es and the rate at which this matter is heated by SNIa explosions crucially determines the dynamical evolution of the gas in these galaxies (Renzini 1996), i.e. whether from inflow to outflow or from outfow to inflow.  Our ability to address these (and other) issues is hampered by the uncertainty which still affects the SNIa progenitors. Current models include Single Degenerate (SD) and Double Degenerate (DD) systems, both dealing with a close binary in which the primary component ($m_1$) evolves into a CO white dwarf (WD). When the secondary ($m_2$) evolves off the MS, Roche Lobe Overflow (RLO) occurs: in the SD model, the WD accretes and grows in mass; in the DD model, the envelope of the secondary is lost from the system, leaving a close double WD, which eventually merge due to the emission of gravitational radiation. In both cases, explosion may occur either when the accreting WD reaches the Chandrasekhar limit and C deflagrates (Ch), or when a He layer of $\\sim$ 0.15 $M_\\odot$ has accumulated on top of the CO WD so that He detonates (S-Ch) (Hillebrandt \\& Niemeyer 2000). The various models correspond to markedly different secular evolution of the SNIa rate, with great variance of the consequences on related issues.   To assess the impact of the SNIa rate on any astrophysical problem it is important to (i) adequately characterize the progenitor model and (ii) correctly couple this model to the SF history of the specific system. In this paper I will go through these steps, and show how the scenarios for the SNIa progenitors can be constrained by their impact on large scales. To do that I will use analytical relations for the SNIa rate for both SD and DDs derived in Greggio (2004), as a convenient parametrization of the progenitors' models.  \\section {The Model}  An effective formulation of the SNIa rate for the modelling of stellar systems rests upon the definition of two key quantitites: the distribution function of the delay times $f_{\\rm Ia}(\\tau)$, where $\\tau$ is the time between the birth of a star and its death as a SNIa, and the number fraction of SNIa events out of one stellar generation, $A_{\\rm Ia}$. Denoting with $\\tau_{\\rm i}$ and $\\tau_{\\rm x}$ the minimum and maximum delay time, and if $A_{\\rm Ia}$ is constant, the SNIa rate at epoch $t$ in a system with a SF rate $\\psi$ is:  \\begin{equation} \\dot{n}_{\\rm Ia}(t) = k_{\\alpha} \\cdot A_{\\rm Ia} \\cdot \\int_{\\tau_{\\rm i}} ^{\\min(t,\\tau_{\\rm x})}{\\psi(t-\\tau) \\cdot  f_{\\rm Ia}(\\tau) \\,\\, {\\rm d} \\tau} \\label {eq_rate} \\end{equation}  \\noindent where $k_{\\alpha}$ is the number of stars per unit mass in each stellar generation (e.g. $k_\\alpha= 2.83$ for Salpeter IMF between 0.1 and 120 $M_\\odot$), considered constant in time.      Eq.(\\ref{eq_rate}) gives the dependence of the observed SNIa rate on the SF history:  for systems with an almost constant $\\psi$, like late type galaxies, at late epochs ($t \\simeq \\tau_{\\rm x}$):  \\begin{equation} \\dot{n}_{\\rm Ia}^{\\rm L}(t) \\simeq \\langle \\psi \\rangle \\cdot k_{\\alpha} \\cdot A_{\\rm Ia} \\label{eq_nial} \\end{equation}  \\noindent since $f_{\\rm Ia}(\\tau)$ is normalized to one. Thus, the current SNIa rate in late type galaxies essentially gives information on the realization probability of the SNIa scenario, since it results from the contribution of systems with all possible delay times. By approximating $\\langle \\psi \\rangle$ with the ratio of the galaxy mass over its age, and assuming an $M_{\\star}/L_{\\rm B}$ ratio of the order of unity, Eq.(\\ref{eq_nial}) requires $A_{\\rm Ia} \\sim 10^{-3}$ to match the observed value of 0.2 SNUs (Cappellaro, Evans, \\& Turatto 1999).  In early type galaxies, where most SF has occurred in an initial burst of relatively short duration $\\Delta t$, the late epoch SNIa rate is:  \\begin{equation} \\dot{n}_{\\rm Ia}^{\\rm E}(t) = M_{\\star} \\cdot k_{\\alpha} \\cdot A_{\\rm Ia} \\cdot \\langle f_{\\rm Ia} \\rangle_{t-\\Delta t,t} \\label{eq_niae} \\end{equation}  \\noindent i.e. it is proportional to the value of $f_{\\rm Ia}$ at delay times close to the age of the galaxy, and to the total mass of stars {\\em formed} in the burst ($M_{\\star}$).  If $A_{\\rm Ia}$ and $k_{\\alpha}$ are the same in early and late type galaxies, the ratio between the current SNIa rates measured in SNU in the two galaxy types is:  \\begin{equation} \\mathcal{R}_{\\rm SNU} = \\frac{{\\dot{n}^{\\rm E}}_{Ia,SNU}}{{\\dot{n}^{\\rm L}}_{Ia,SNU}} \\simeq \\frac {(M_\\star/L_{\\rm B})^E} {(M_\\star/L_{\\rm B})^L} \\times \\frac {\\langle f_{\\rm Ia} \\rangle_{t-\\Delta t,t}} {\\langle f_{\\rm Ia}\\rangle_{\\tau_{\\rm i},t}}. \\label{eq_rsnu} \\end{equation}  \\noindent Using Maraston (1998) solar metallicity models, at an age of 12 Gyr a system formed at a constant rate has  $M_\\star/L_{\\rm B} \\simeq  2$, while a single burst stellar population has $M_\\star/L_{\\rm B} \\simeq 13$. Since $\\mathcal{R}_{\\rm SNU} \\sim 1$ (Cappellaro et al. 1999), Eq. (\\ref{eq_rsnu}) implies:  \\begin{equation} \\frac {\\langle f_{\\rm Ia} \\rangle_{t-\\Delta t,t}} {\\langle f_{\\rm Ia}\\rangle_{\\tau_{\\rm i},t}} \\sim  0.15 \\label{eq_frat} \\end{equation}  \\noindent showing that, on the average, the distribution of the delay times must be decreasing with $\\tau$, i.e. {\\em the majority of SNIa precursors are relatively short lived}.   \\vspace {0.3cm}  The theoretical SNIa rate can be obtained with numerical simulations, which follow the evolution of a population of binary systems, under some prescriptions concerning the outcome of the mass transfer phases (e.g. Han 1998). Typically, the predicted $A_{\\rm Ia}$ ranges within $10^{-4}$ and $10^{-3}$, while $f_{\\rm Ia}$ is charaterized by an early steep rise, followed by a decline at later epochs. These results depend on a variety of input parameters, whose role is difficult to gauge, and therefore they are hardly suitable for a thorough exploration of the parameter space. On the other hand, the shape of the $f_{\\rm Ia}(\\tau)$ function  can be characterized on the basis of general considerations founded on stellar evolution. Along this line, Greggio (2004) derives analytical relations for the $f_{\\rm Ia}$ function for both the SD and the DD model, briefly described below.  \\begin{figure}[!ht] \\plottwo{greggio_fig1a.eps}{greggio_fig1b.eps} \\caption{The SD model: $n(m_{2})$ is derived from the distributions $n(m_1) \\propto m_{1}^{-\\alpha}$ and $n(q) \\propto q^\\gamma$ where $q=m_2/m_1$, and summing over the relevant combinations. The right panel shows the $f_{\\rm Ia}^{SD}$ function on arbitrary units, as from Eq. (\\ref{eq_fiasd}). $\\tau$ is the MS lifetime of the mass $m_2$. $\\epsilon$ is the ratio between the accreted mass and the envelope mass of the donor. Curves are shown for $\\epsilon=1$ and $\\epsilon=0.5$ (labelled). For the Sub-Ch models, it is assumed that $m_{1} \\geq 3$.}  \\label{fig1} \\end{figure}  \\begin{figure}[!ht] \\plottwo{greggio_fig2a.eps}{greggio_fig2b.eps} \\caption{The DD model: dependence on $\\tau_{\\rm n,x}$ (left), and on the slope of the distribution of the gravitational delays (right). See text for more details.} \\label{fig2} \\end{figure}   In the SD model, the delay time is equal to the MS lifetime of the secondary, plus the duration of the accretion phase, which is negligible. Thus:  \\begin{equation} f_{\\rm Ia}^{SD}(\\tau) \\propto n(m_{2}) \\cdot |\\dot{m_{2}}|. \\label{eq_fiasd} \\end{equation}  \\noindent where $|\\dot{m_{2}}|$ is the rate of change of the mass whose MS lifetime is $\\tau$, and $n(m_{2})$ is the distribution function of the secondary masses {\\em in those systems which produce SNIa events}. Fig. \\ref{fig1} shows the two factors in Eq.(\\ref{eq_fiasd}) (left panel), and their product (right panel). The shape of the $n(m_2)$ function reflects the limitations imposed on the components masses in order to give rise to a SNIa. It is assumed that $m_1$ varies between 2 and 8 $M_\\odot$, since most CO WD come from this range. The lower boundary originates a relative lack of systems once $m_2$ drops below 2 $M_\\odot$. The last steep decrease, instead, reflects the need of building a WD up to the Ch limit: as $m_2$ decreases, so does its envelope, that is the donated mass. Then, at late epochs, only the most massive WD will end up as SNIa, which implies a lower limit of $m_1 > 2 M_\\odot$. Below $m_2=2 M_\\odot$, $n(m_2)$ is very sensitive to $\\alpha, \\gamma, \\epsilon$ (defined in the caption), which determine the volume in the parameter space of SNIa progenitors. The resulting $f_{\\rm Ia}$ function in shown in the right panel: notice the dramatic drop at late epochs which stems from the requirement of building up to the Ch mass. S-Ch exploders (dot dashed line) do not show this drop, the only requirement being that the CO WD is more massive than 0.7 $M_\\odot$. If additional limitations are imposed on $m_2$, the $f_{\\rm Ia}$ function is non zero only in the range of $\\tau$ equal to the MS lifetimes of the secondaries with appropriate mass. For example, if systems with $m_2 < 2 M_\\odot$ evolve into a common envelope, $f_{\\rm Ia}^{\\rm SD}=0$ for $\\tau > 1$ Gyr. On the other hand, in order to explain the occurrence of SNIa in Es with the SD model, low mass secondaries must be allowed for.  In the DD model, the delay time is equal to the MS lifetime of the secondary ($\\tau_{\\rm n}$), plus the gravitational delay ($\\tau_{\\rm gw}$). Considering only CO DDs, $m_2$ is (in most cases) heavier than 2 $M_\\odot$, which implies $\\tau_{\\rm n} \\leq 1$ Gyr; $\\tau_{\\rm gw}$ instead can span a very wide range, being extremely sensitive to the separation $A$ of the DD. For example, for a (0.7+0.7) $M_\\odot$ system, $\\tau_{\\rm gw}$ goes from 14 Myr to 18 Gyr when $A$ increases from 0.5 to 3 $R_\\odot$. Since the separation of the progenitors of the CO DDs ranges from some tens to some hundreds $R_\\odot$, a great degree of shrinkage must occur at the RLO phases in order to produce explosions within a Hubble time. $f_{\\rm Ia}^{\\rm DD}$ can be derived as a modification of $f_{\\rm Ia}^{\\rm SD}$, with early explosions from systems with short $\\tau_{\\rm n}$ {\\em AND} $\\tau_{\\rm gw}$, and late explosions from systems with long $\\tau_{\\rm gw}$. The distribution of $\\tau_{\\rm gw}$ will be very sensitive to the distribution of $A$ within the range of a few $R_\\odot$. Motivated by literature results, Greggio (2004) considers two possibilities: (i) the WIDE DDs, in which the first RLO corresponds to a modest shrink of the system (as, e.g., in Nelemans et al. 2001); (ii) the CLOSE DDs, similar to the classical recepy (as, e.g., in Han 1998) . The two channels imply different characterizations of the distribution of the gravitational delays. Fig. \\ref{fig2} shows the resulting $f_{\\rm Ia}^{\\rm DD}$ functions for some choice of the  major parameters: $\\tau_{\\rm n,x}$, the MS lifetime of the least massive secondary still progenitor of a SNIa; $\\beta_{\\rm a}$ and $\\beta_{\\rm g}$ which are the exponents of the two power laws assumed to represent the distribution of $A$ for the WIDE DDs, and of $\\tau_{\\rm gw}$ for the CLOSE DDs \\footnote{ If SNIa originate from DD systems with approximately the same mass, $\\beta_{\\rm g} \\simeq -0.75+0.25\\beta_{\\rm a}$}.  The $f_{Ia}^{DD}$ function shows an early rise, a wide maximum and a late epoch decline. The width of the maximum is $\\simeq \\tau_{\\rm n,x}$. Relative to the CLOSE DD, the WIDE DD scheme of evolution leads to flatter $f_{\\rm Ia}^{\\rm DD}$ distributions, and the more systems with short $\\tau_{\\rm gw}$ (i.e. steeper $\\beta_{\\rm a}$ or $\\beta_{\\rm g}$), the steeper the early rise as well as the late epoch decline. ",
        "conclusions": ""
    },
    "0410/astro-ph0410323_arXiv.txt": {
        "abstract": "The majority of the ultraluminous X-ray sources (ULXs) in external galaxies are believed to be accreting black holes in binary systems; some of the black holes could be as massive as $\\sim 100-1000\\,\\ms$. We have performed evolution calculations for intermediate-mass black hole X-ray binaries, assuming they are formed in dense star clusters via tidal capture. The results are compared with those for stellar-mass black holes X-ray binaries. We find that these two types of black holes may have similar companion stars and binary orbits if observed as ULXs. However, intermediate-mass black holes seem to be favored in explaining the most luminous ULXs. We also discuss the possibilities of transient behavior and beamed emission in the evolution of these binary systems. ",
        "introduction": "Ultraluminous X-ray sources (ULXs) are point-like extra-nuclear X-ray sources found in nearby galaxies, with (isotropic) X-ray luminosities in excess of $10^{39}$ ergs$^{-1}$, roughly corresponding to the Eddington limit accretion luminosity for a 10 $\\ms$ star \\citep{f89}. The majority of ULXs are believed to be black holes accreting from their binary companion stars \\citep[see, e.g., reviews by][]{fw04,cm04}, or in few cases, perhaps from the fallback material originating from supernovae that have produced these sources \\citep{li03}. Many works have shown that stellar-mass black holes (SMBHs) can account for the properties of  most ULXs, but for the most luminous ones (with X-ray luminosities $L_{\\rm X}\\gsim 10^{40}\\,\\ergs$), assumptions of either truly super-Eddington X-ray emission \\citep{b02}, or anisotropic (beamed) emission \\citep{k01,kfm02} are required. Alternatively, the latter sources could be intermediate-mass black holes \\citep[IMBHs;][]{cm99}. This suggestion seems to be supported by their relatively low disk color temperatures inferred from X-ray spectral analysis \\citep[e.g.][]{mfm04}.  It is interesting and important to discriminate between these two kinds of black holes observationally. \\citet{k04} suggested that IMBHs in X-ray binaries are more likely to be transient due to the thermal-viscous instability \\citep{kkb96}. The reason is that, both the high masses and wide orbits of IMBHs lead to a larger accretion disk and hence lower temperature in the disk compared to the SMBH case. In this Letter, with different assumptions on IMBH binary formation from \\citet{k04}, we have calculated the evolution of IMBH X-ray binaries with a massive donor star. We describe the initial parameters for the binary evolution and the calculated results in section 2 and 3 respectively. In section 4 we discuss the relation between IMBHs, SMBHs, and ULXs. Our conclusions are summarized in section 5. ",
        "conclusions": "The purpose of this work is to examine whether the hypothesized IMBHs can reproduce the observed properties of (some of) the ULXs. There have been several works on this subject in the literature. In the binary evolution calculations we have used the results in \\citet{hpza04} on tidal capture to determine the initial conditions for the incipient IMBH binaries. But we allowed stellar expansion before Roche-lobe overflow, and got the following results different from theirs: (1) the donor stars have already evolved off main-sequence when mass transfer begins; (2) the X-ray lifetime for IMBHs as ULXs is around $10^6$ yr, much shorter than the main-sequence lifetime of the donor stars. These can present constraints on the donor stars of IMBH-ULXs. For example, the age of the star cluster MGG-11 in the irregular galaxy M82 is $7-12$ Myr. If the ULX M82 X-7 in this cluster \\citep{mt99} is powered by an IMBH of mass $\\sim 1000\\,\\ms$, a $10-15\\,\\ms$ donor star is required to account for the luminosity and age of the ULX. In a more recent paper, \\citet{pzdm04} performed detailed binary evolution calculations of $2-15\\,\\ms$ stars which transfer mass to a $100-2000\\,\\ms$ BH. These authors initialized the binary systems with an (arbitrarily) selected orbital period at birth (ranging from less than a day to a few hundred days). Their calculated mass transfer rates and time with initial orbital period of the order of a day are in general agreement with ours.  Systematic studies on evolution of SMBH X-ray binaries have been carried by \\citet{prh03} and \\citet{rpp04}. The results in Fig.~2 can be regarded as individual cases of their population synthesized results. We notice that different parameterization for estimating the X-ray luminosities has been adopted here and in their works. \\citet{rpp04} assumed an energy conversion coefficient depending on the BH spin and let $L_{\\rm X}$ be as high as 10 times of the Eddington luminosity, as suggested by \\citet{b02}. We have used the results of \\citet{o02}, who calculated the emission from the supercritical accretion disk, considering the photon trapping effect. This difference causes the mass transfer rate to be about 10 times higher in our work (i.e., $\\dot{M}\\sim 10^{-5}\\,\\my$) than in \\citet{rpp04}, for $L_{\\rm X}$ to reach $10^{40}\\,\\ergs$, for a $10\\,\\ms$ BH. The duration of the high-luminosity X-ray emission phase is correspondingly $\\sim 10$ times shorter to be $\\sim 10^6$ yrs for a massive donor star.  From evolution calculations we have shown that, as ULXs, both IMBH and SMBH X-ray binaries actually have similar characteristics in their companions stars and binary orbits, which make it difficult to discriminate them observationally. However, if the donor star is more massive than $\\sim 10\\,\\ms$, the IMBH binaries are most likely to reside deeply inside the host star cluster, while SMBH binaries generally have been \"kicked\" outside \\citep{z02}. For an intermediate- or low-mass donor star, the IMBH binaries could be hostless, since the preceding time before mass transfer could be long enough for the cluster to be tidally dispersed, if it is close to the galactic center \\citep{p04}.  For the most luminous ULXs with $L_{\\rm X}\\gsim 10^{40}\\,\\ergs$, a potential problem related to the SMBH model  is that, according to \\citet{o02}, the effective temperatures of supercritical accretion disks with $\\dot{M}\\sim 10^{-5}\\,\\my$ are still as high as $\\sim 1$ keV, inconsistent with spectral analyzed results for these sources \\citep[e.g.][]{mfm04}. It is not known how the temperature profile is in the \\citet{b02} disk model. On the other hand, IMBHs have the advantage of low disk temperatures because of their high masses.  We also find that transient behavior is a common feature in late evolutionary stages in these two types of binaries. It may be difficult to tell IMBHs from SMBHs by transient behavior, as suggested by \\citet{k04}. The differences between our work and \\citet{k04} result from the following facts. (1) We have adopted the BH masses no more than $1000\\,\\ms$ in the calculations. With these values, one can actually draw from Fig.~1 in \\citet{k04} the similar results as ours. (2) We assumed tidal capture as the main formation channel of IMBH binaries, from which we got the initial orbital periods of $\\sim 2-4$ days; \\citet{k04} favor orbital periods in excess of $\\sim 100$ days, but the duration of the mass transfer episode would decrease to be very short (less than a few $10^4$ yrs), making them difficult to be observed.  It should be noted that the calculated mass transfer rates are long-term, averaged ones. It is unclear how to relate these secular mass transfer rates to observable instantaneous X-ray luminosities, and to disk instability influenced by (short-term) mass transfer and X-ray irradiation. Moreover, the mechanisms and criterion for the thermal-viscous instability in irradiated accretion disk are not yet fully understood. It is premature to predict the instability occurrence from only calculated mass transfer rates.  Another interesting feature is that, if anisotropic (or beamed) emission is associated with mass transfer rates comparable to the Eddington rate, as suggested before \\citep[e.g.][]{k01}, our calculations indicate that IMBHs may also have anisotropic emission, since the mass transfer rates can be sufficiently high to satisfy the condition above. However, anisotropic X-ray emission is not preferred in our opinion for persistent ULXs, since the most important stabilizing factor for disk instability, the efficiency of X-ray irradiation, will be greatly reduced if the emission is beamed (usually in the direction perpendicular to the disk plane), and the ULXs would become transient sources."
    },
    "0410/astro-ph0410494.txt": {
        "abstract": "We have carried out a study of the early type stars in nearby OB associations spanning an age range of $\\sim$ 3 to 16 Myr, with the aim of determining the fraction of stars which belong to the Herbig Ae/Be class. We studied the B, A, and F stars in the nearby ($\\le 500$ pc) OB associations Upper Scorpius, Perseus OB2, Lacerta OB1, and Orion OB1, with membership determined from Hipparcos data. We also included in our study the early stars in the Trumpler 37 cluster, part of the Cep OB2 association. We obtained spectra for 440 Hipparcos stars in these associations, from which we determined accurate spectral types, visual extinctions, effective temperatures, luminosities and masses, using Hipparcos photometry. Using colors corrected for reddening, we find that the Herbig Ae/Be stars and the Classical Be stars (CBe) occupy clearly different regions in the JHK diagram. Thus, we use the location on the JHK diagram, as well as the presence of emission lines and of strong 12 $\\mu$m flux relative to the visual to identify the Herbig Ae/Be stars in the associations. We find that the Herbig Ae/Be stars constitute a small fraction of the early type stellar population even in the younger associations. Comparing the data from associations with different ages and assuming that the near-infrared excess in the Herbig Ae/Be stars arises from optically thick dusty inner disks, we determined the evolution of the inner disk frequency with age. We find that the inner disk frequency in the age range 3 - 10 Myr in intermediate mass stars is lower than that in the low mass stars ($<$ 1 \\msun); in particular, it is a factor of $\\sim$ 10 lower at $\\sim$ 3 Myr. This indicates that the time-scales for disk evolution are much shorter in the intermediate mass stars, which could be a consequence of more efficient mechanisms of inner disk dispersal (viscous evolution, dust growth and settling toward the midplane). ",
        "introduction": "\\label{sec:int}  The Herbig Ae/Be stars (HAeBe) are young emission line objects with spectral types B, A, and in a few cases F, in most instances spatially  correlated with dark clouds or bright nebulosities \\citep{herbig60,waters98,hernandez04}. The masses of HAeBe range from 2 to 10 {\\msun}. Like their low mass pre-main-sequence (PMS) counterparts, the T Tauri stars (TTS), these objects become optically visible before reaching the main sequence, so their PMS evolution can be studied in some detail. Stars more massive than 10 {\\msun} are expected to spend their whole PMS time as optically obscured objects.  In addition to H$\\alpha$ emission, HAeBe exhibit infrared excesses relative to the photosphere which are attributed to emission from dust in circumstellar accretion disks \\citep{finkenzeller84,lorenzetti83, davies90,hill92,ancker97,malfait98}. Millimetric and submillimetric observations confirm the existence of disks of substantial mass around some of these objects \\citep{blake04,fuente03,mannings00,mannings97,natta00, natta01}. Most HAeBe exhibit a flux excess in the near infrared (NIR) portion of their spectral energy distributions (SEDs), which has been attributed to emission from an optically thick ``wall'' at the dust destruction radius, directly heated by the star \\citep{natta01, maheswar02,dullemond01,dullemond04a,muzerolle04, muzerolle03b}.  Studies of the circumstellar environments of HAeBe suggest that  these objects are progenitors the stars surrounded by debris disks ( e.g, $\\beta$ Pictoris, Vega-like stars), which could be sites of planet formation \\citep{natta00, lagrange00}. However,  long standing questions about how and when the transition between HAeBe  and Vega-like stars occurs have yet to be answered. Observations  of  the numbers of young intermediate mass stars surrounded by optically thick inner disks in samples of different ages could be used as  a first step toward answering these questions. Roughly coeval stellar groups are ideal to study since uncertainties in properties like distances and ages are minimized.  OB associations are defined as stellar groups with significant populations of intermediate mass stars and with stellar mass density less than 0.1 \\msun/$pc^3$ \\citep{brown99b}. Ages derived from Hertzsprung-Russell diagrams (HRD) indicate that they are young \\citep[e.g,][]{blaauw64}. For this reason, OB associations are excellent for studying young stellar objects covering a complete range of mass \\citep[e.g. ][]{preibisch02}. Moreover, by studying OB associations with different ages, we can analyze the temporal evolution of the processes characterizing the early history of their stars.  An important but difficult aspect of studying OB associations is the identification of their members. OB associations have small internal velocity dispersions, so the streaming motion of the stellar group, in combination with the solar motion, is reflected as a motion of their members toward a convergent point on the sky \\citep[e.g.,][]{brown99a,zeeuw99}. Thus, membership can be established by studying the kinematic properties of the group (proper motions, parallaxes, radial velocities). Nearby associations cover large regions in the sky, which in the past limited the astrometric membership determination to the bright stars (V $\\leq$ 6 mag). Photometric studies could add fainter members to the associations, but this membership determinations were less reliable due to several factors as, for instance, undetected duplicity, spread of distances within the association, photometric variability, and interlopers. The publication of positions, proper motions and trigonometric parallaxes for $\\sim$120000 stars in the Hipparcos Catalog \\citep{esa97} has significantly improved the situation, enabling astrometric studies down to its limiting magnitude, V$\\sim$12.  Using  Hipparcos measurements, \\citet{zeeuw99} made a comprehensive census of the stellar content of the OB associations within 1 Kpc from the sun. Members of each association were identified using two methods. The first method is a modification of the classical convergent point method \\citep{brown50} detailed by \\citet{bruijne99}; the second uses the positions, parallaxes and proper motions to define members using probability distributions in the velocity space \\citep{hooger99}. The combination of these methods gives reliable memberships of stars in 12 young stellar groups. However, since some groups are beyond 500 pc, where the Hipparcos parallaxes are no longer useful, or have unfavorable kinematics, 10 associations studied by \\citet{zeeuw99} do not show astrometric evidence for a moving group and their results are inconclusive. One of these associations with unfavorable kinematics is Orion OB1, which has a motion mostly directed radially away from the Sun. Nonetheless, the Hipparcos data of Ori OB1 was analyzed carefully by \\citet{brown99b}, who found a relation based on proper motions which characterizes approximately the members of this association (see \\S 2.2). The resulting set of stars selected by this relation overlaps in 96\\% with the photometric members given by \\citet{brown94}.    In this contribution we explore the frequency of HAeBe in three associations with different ages (Upper Scorpius, Perseus OB2, and Lacerta OB1) with members selected from \\citet{zeeuw99}. In addition, we study the Orion OB1 association,  which was divided in 2 sub-association with different ages and distances. We complement our studies by analyzing the NIR properties of the early stars in Trumpler 37  from \\citet{contreras02}. Details of the selection and the observation of each sample is described in \\S2. In our entire sample, more than 82\\% of the spectral types published by Hipparcos are from the SIMBAD\\footnote{SIMBAD, Centre de Donn\\'{e}es astronomiques de Strasbourg: http://simbad.u-strasbg.fr/sim-fid.pl} data-base, with references labeled as miscellaneous; the other 18\\% are from the Michigan catalogue for the HD stars \\citep{houk82,houk88}. We obtained spectra for all the stars in our sample and reclassified them using the classification scheme described in \\citet{hernandez04},  to obtain an homogeneous spectral typing scheme for all the objects in the associations. Details of the classification scheme are in \\S 3.1, where we also calculate the visual extinction for the objects and describe the identification and measurement of the emission lines present in the spectra.  We analyze the NIR excess of the members of each group in \\S 3.2. We estimate distances from both astrometric and photometric data in \\S 3.3, and a HRD  is given for each association in \\S 3.4. A census of HAeBe, based on observational properties characteristic of the presence of disks (emission at H$\\alpha$, NIR excesses and IRAS fluxes) is described in \\S 4, and the disk frequency is discussed on \\S5. Finally, we summarize the most important aspects of this work in \\S6. ",
        "conclusions": "In this contribution we studied the presence of inner disks around intermediate mass pre-main sequence stars in 5 stellar groups with range of ages from 3 to 16 Myr. The stars were selected from the Hipparcos catalog. The astrometric membership for US, Per OB2 and Lac OB1 associations was determined  from \\citet{zeeuw99}. The stars in Orion OB1 were selected as members using astrometric and photometric constrains \\citep{brown94,brown99b} and divided in two groups, Ori OB1a and  Ori OB1bc, with different ages and distances. We obtained spectra for 440 stars in these OB associations, from which we have determined spectral types, visual extinctions, \\teff, luminosities and the presence of emission at H$\\alpha$. We also included data for the cluster Trumpler 37 in the association Cep OB2, with data from \\citet{contreras02}.  We compared the dereddened J-H and H-K colors of a samples of HAeBe stars \\citep{hernandez04} and CBe \\citep{yudin01} in order to estimate the loci of these objects, showing that the HAeBe and the CBe occupy well defined different regions on this diagram.  We identified the Herbig Ae/Be stars in each association using the following criteria (1) H$\\alpha$ emission, (2) location in the HAeBe region of the JHK diagram, (3) strong IRAS fluxes with \\citep[$\\beta$ index $>$ -2; ][]{vieira03}. We identified 2 HAeBe in Upper Scorpius, 0 HAeBe in Perseus OB2, 0 HAeBe in Lacerta OB1, 2 HAeBe in Orion OB1a, 4 HAeBe in Orion OB1bc and 2 HAeBe in Trumpler 37. From these identifications, we estimate a disk frequency of 3.3$\\pm$1.7\\%, 0$\\pm$4.4\\%, 0$\\pm$1.9\\%, 2.5$\\pm$1.7\\%, 5.1$\\pm$2.0\\%, 4.3$\\pm$1.8\\%,  respectively. We have compared the inner disk frequency in intermediate mass stars (HAeBe star) with that for low mass stars in a number of stellar groups taken from the literature, and estimated by us for Ori OB1 from the sample of \\citet{briceno04}. We find that stars with inner disks are more frequent in low mass stars than in intermediate mass stars. Particularly, for the youngest associations (Ori OB1bc and Tr 37) the inner disk frequency for intermediate mass stars is $\\sim$ ten times lower than the inner disk frequency reported for low mass stars. Although the inner disk frequency in intermediate mass stars tends to decrease slowly with age, we find no evidence of the strong decrease in the inner disk frequency observed in low mass stars at 3-5 Myr. Studies of intermediate mass stars in stellar groups younger than 3 Myr are required to determine in which age range, if any, a similar strong decrease with age occurs. We suggest that the lower inner disk frequency in intermediate mass stars is a result of more rapid mechanisms of inner disk dispersal (accretion, dust growth and settling toward the midplane).  This contribution suggests that most of the stars in the range of mass of the HAeBe disperse their inner disks in a timescale shorter than 3 Myr; in contrast, the inner disk frequency of low mass stars in this age range is larger than 50\\%. Samples of younger stellar groups are required to determined when the most of the intermediate mass stars disperse their inner disks."
    },
    "0410/astro-ph0410579_arXiv.txt": {
        "abstract": "The chemical evolution history of a galaxy hides clues about how it formed and has been changing through time. We have studied the chemical evolution history of the Milky Way (MW) and Andromeda (M31) to find which are common features in the chemical evolution of disc galaxies as well as which are galaxy-dependent. We use a semi-analytic multi-zone chemical evolution model. Such models have succeeded in explaining the mean trends of the observed chemical properties in these two Local Group spiral galaxies with similar mass and morphology. Our results suggest that while the evolution of the MW and M31 shares general similarities, differences in the formation history are required to explain the observations in detail. In particular, we found that the observed higher metallicity in the M31 halo can be explained by either a) a higher halo star formation efficiency or b) a larger reservoir of infalling halo gas with a longer halo formation phase. These two different pictures would lead to a) a higher $[$O/Fe$]$ at low metallicities or b) younger stellar populations in the M31 halo, respectively. Both pictures result in a more massive stellar halo in M31, which suggests a possible correlation between the halo metallicity and its stellar mass. ",
        "introduction": "The chemical properties of galaxies hide clues about their formation and evolution. Semi-analytic chemical evolution models (\\citealt{TA}; \\citealt{Tinsley}) have succeeded in explaining the mean trends of galactic systems by numerically solving a set of equations governing the simplified evolution of the chemical elements as they cycle through gas and stars.  One strength of these models is that they typically have the fewest number of free parameters, making convergence to a smaller set of solutions more likely.  Strong constraints can be placed on chemical evolution models only by contrasting them against a comprehensive set of observed properties. Since the most detailed observational data are generally available for the Milky Way (e.g. \\citealt{FBH}), successful agreement between model predictions and these observed properties has been obtained by several studies in the past, which help us to understand the formation history of the Milky Way (MW). Some models have focused on the evolution of the chemical abundances both in the solar neighbourhood and in the whole disc, adopting a framework in which the MW has been built-up inside-out by means of a single accretion event (e.g. \\citealt{MF}; \\citealt{PT}). Others have used an early infall of gas to explain the thick disc formation, followed by a slower infall to form the thin disc (e.g. Chiappini, Matteucci \\& Gratton 1997). Several studies have paid particular attention to the chemical evolution of a larger range of elements (e.g. Timmes, Woosley \\& Weaver 1995; \\citealt{GP}; Alib\\'es, Labay \\& Canal 2001) or focused on the features of the Metallicity Distribution Function (MDF) in the solar neighbourhood (e.g. \\citealt{FG}). However, an implicit assumption remains unanswered: \\textit{is the MW a typical spiral?}  Andromeda (M31) is the closest spiral to our MW (e.g. \\citealt{vanDenBergh2003}). Previous theoretical studies of the chemical evolution of M31 have been done by \\cite{DT} and Moll\\'a, Ferrini \\& Diaz (1996), both emphasising the evolution of the M31 disc. They concluded that M31 has a formation history and chemical evolution similar to that of the MW. More recently, however, there has been considerable observational progress in the study of both galaxies. A striking difference between M31 and the MW is that the metallicity of the M31 halo ($\\langle$[Fe/H]$\\rangle\\approx-0.5$) \\footnote{Hereafter [X/Fe]$=\\log_{10}(X/Fe)-\\log_{10}(X/Fe)_{\\odot}$.} is significantly higher than its MW analogue ($\\langle$[Fe/H]$\\rangle\\approx -1.8$, e.g. \\citealt{RN}), as revealed by many recent studies (e.g. Holland, Fahlman \\& Richer 1996; Durrell, Harris, \\& Pritchet 2001; \\citealt{SvD}; \\citealt{FJ}; \\citealt{RG}; \\citealt{FergusonETal}; \\citealt{WE}; \\citealt{BellazziniETal}; \\citealt{BrownETal}; \\citealt{ReitzelETal}; \\citealt{RichETal}).  It is therefore timely to attempt the construction of a chemical evolution model for both the MW and M31, using the same framework. Such an attempt may be helpful in highlighting the common features in the chemical evolution of spiral galaxies (at least in these two spirals) and those which remain galaxy-dependent.  In Section~2 we describe our multi-zone chemical evolution model, and in Section~3 we present the results of our MW model. In Section~4 we show the results for M31. Finally, our results are discussed in Section~5. ",
        "conclusions": "The results in the previous section have shown that the main trends of the chemical evolution history of both the MW and M31 can be described using the same framework. We have also found that to explain the observations in more detail, different sets of parameter values for the formation history are required for the MW and M31, respectively. In this section, we discuss how these different parameters values relate to different galaxy formation histories.  The previous sections have shown that the main differences between the MW and M31 include the hydrogen surface density in the disc and, more significantly, the halo metallicity. The smaller hydrogen surface density in M31 can be explained by a combination of the shorter disc infall time-scale and the higher disc SFE.  The most striking difference is that the mean metallicity observed in the M31 halo is about ten times higher than that observed in the MW halo, though both galaxies have similar mass and morphology. \\textit{Which is a ``typical'' halo?} Metal-rich halos seem more common than metal-poor ones, as pointed out by \\cite{ZibettiETal} who analysed a sample of about 1000 edge-on disc galaxies within the SDSS. \\cite{HH} have also shown that the NGC5128 halo MDF closely resembles that in the M31 halo. This similarity is perhaps suggestive of a common history in the halo formation of both galaxies, despite their different Hubble-type.  A straightforward way to obtain a more metal-rich halo would be by requiring homogeneous pre-enrichment of the infalling halo gas with a metallicity $Z\\approx 0.1~Z_{\\odot}$. However, this metallicity seems too high, if such pre-enrichment occurred homogeneously in the whole universe. For example, although this value is close to the mean metallicity of Damped Lyman $\\alpha$ systems (DLAs), many DLAs have much lower metallicity (\\citealt{Pettini}; \\citealt{ProchaskaETal}). If DLAs are more evolved systems than the infalling halo gas, their metallicity should be higher than that induced by pre-enrichment. Therefore, a pre-enrichment of $Z\\approx 0.1~Z_{\\odot}$ may not be tenable.  We found that the observed metal-rich halo of M31 could be explained by two scenarios without any pre-enrichment: a) a higher halo SFE; b) a larger reservoir of infalling halo gas with a longer halo phase. Scenario~a) might be explained in a hierarchical clustering regime, where the halo is formed by accretion of building-blocks. Theory predicts that gas removal by supernova-driven winds should be more effective in lower mass systems, leading to a consequent suppression of the SFR and a low SFE (e.g. \\citealt{DS}; \\citealt{E}). In addition, observations based on the SDSS also suggest that the SFE decreases with decreasing stellar mass in low mass galaxies with a stellar mass M$_{\\star}< 3\\times10^{10}$~M$_{\\odot}$ \\citep{KauffmannETal}. Therefore, higher SFE in the M31 halo phase would be explained if the building-blocks of the M31 halo are systematically higher mass systems than those of the MW halo (see also \\citealt{HH}). This result again supports the notion of diversity in the formation of spiral galaxies which are apparently similar in mass and morphology and belong to the same Local Group. Admittedly, our model does not adhere to the hierarchical-clustering scenario in a self-consistent manner, however, our chemical evolution models can be interpreted as the ``mean'' SF and chemical evolution history of the stars which end up at each radius.  As an observational consequence of our models, we found that both scenarios a) and b) predict a more massive stellar halo in M31, respectively about 6 and 9 times more massive than in the MW, whose stellar halo mass is $\\approx 10^{9}$~M$_{\\odot}$. This result suggests that there might be a correlation between the halo metallicity and its stellar mass. Using dynamical simulations, Bekki, Harris \\& Harris (2003) show that the stellar halo comes from the outer part of the progenitor discs when the bulge is formed by a major merger of two spiral galaxies. Based on this, they also predict the correlation of the metallicity of the stellar halos and the mass of the bulges which were formed by major mergers, since larger bulges have the larger progenitors, and progenitor spiral galaxies should follow the observed mass-metallicity relation. Although they do not mention a correlation between the masses of the bulge and halo, it would naturally be expected. Thus, a major merger scenario could explain our conclusion. Future observational surveys will better quantify the correlation between the halo metallicity and its stellar mass (e.g. \\citealt{MouhcineETal}).  It is possible to distinguish scenarios~a) and b) through observation. First, due to a longer halo phase, scenario~b) predicts intermediate-age population in the M31 halo. This picture could explain the recent evidence of intermediate-age population by deep imaging of the M31 halo \\citep{BrownETal}.  It is worth mentioning that the metallicity gradient in the stellar halo is sensitive to the assumed stellar halo density profile especially in scenario~b); a steeper density profile leads to a steeper metallicity gradient. \\cite{PvdB1994} suggested that the outer halo of M31 can be modeled by a power law surface brightness profile of $I\\propto R^{-4}$, which is much steeper than what we assumed ($\\propto R^{-2}$). We found that such a steep profile rules out scenario~b) to reproduce the flat metallicity gradient observed in the M31 halo. However, it is still difficult to accurately measure the halo surface brightness of M31 (e.g. \\citealt{ZibettiETal}). Thus, more observational estimates of the M31 halo surface brightness profile would be an important test for this scenario.  We also found that the higher halo SFE in scenario~a) leads to about 0.2~dex higher [O/Fe] when [Fe/H]$<$~-1~dex, due to intense halo SF. Although measuring [O/Fe] is a hard challenge for the current available instruments, this task could be accomplished by the next-generation large-aperture optical telescope.   \\subsection{Prospect}  The framework we have used can explain the main trends in the chemical properties of both the MW and M31. However, recent observations of stellar streams both in the MW (e.g. \\citealt{HWdZ99}; \\citealt{CB2000}; \\citealt{ILIC02}; \\citealt{BrookETal}; \\citealt{NHF04}; \\citealt{MKL04}) and M31 (\\citealt{IbataETal}; \\citealt{McConnachieETal}; \\citealt{MerrettETal}; \\citealt{ZuckerETal}; \\citealt{LewisETal}) clearly identify inhomogeneities in the chemical and dynamical history of both galaxies, which could suggest that a significant fraction of the halo stars results from later accretion of satellite galaxies. In the light of these recent observations, it will be important to study both the MW and M31 in more detail by employing a self-consistent chemo-dynamical model (e.g. \\citealt{BrookETal2004}) to trace their interrelated chemical and dynamical histories."
    },
    "0410/astro-ph0410053_arXiv.txt": {
        "abstract": "We argue that the recent thermal model of jet launching by young stellar objects, when applied to system containing disk-accreting white dwarfs naturally explain the otherwise astonishing absence of jets in cataclysmic variable stars. Thermal launching is possible when the accreted material is strongly shocked due to large gradients of physical quantities in the boundary layer (or at the inner boundary of a truncated disk) and then cools on a time scale longer than its ejection time from the disk. In our framework the magnetic fields are weak, and serve only to recollimate the outflow at large distances from the source, or to initiate the shock, but not as a jet-driving agent. Using criteria for shock formation and mass ejection, we find the mass accretion rate above which jets can be launched from boundary layers around accreting white dwarfs to be $\\dot M_{\\rm WD} \\ga 10^{-6} \\textrm{M}_\\odot \\textrm{y}^{-1}$, which explains the absence of jets in cataclysmic variable stars and their presence in other white-dwarf accreting systems such as super-soft X-ray sources, symbiotic stars and classical novae. ",
        "introduction": "It is widely believed that accretion with angular momentum leads to ejection of jets. This belief is based on observations of jets in Young Stellar Object (YSOs), Low-Mass and High-Mass X-ray Binaries (LMXBs \\& HMXBs) and Active Galactic Nuclei. Cataclysmic Variables (CVs), however, are a blatant exception from the presumed universality of the accretion -- jet connection. However, CVs are not an exception from a more general, supposedly universal, relation between accretion and ``ejection\". Indeed, although no jets have ever been observed in CVs, some of them emit winds. For example, P Cygni profiles in resonant UV lines are observed in some very luminous CVs such as the nova-like stars and dwarf novae at outburst maximum. These winds are too cold to be ejected by a thermal mechanism and are most probably driven by radiative line pressure with some help of magnetic fields (see e.g. Proga, these proceedings). But they are winds, even if slightly collimated, but not jets.  The presence of a white dwarf in the center of an accretion flow cannot be considered to be responsible for this jet-blowing impotence since jets are observed in other accreting white-dwarf systems such as Super Soft X-ray Source (SSXS), symbiotic stars and novae. Fast, $\\sim 1000-5000 \\rm \\ km\\ s^{-1}$, collimated outflows have been observed in some SSXSs, RX J0513.9-6951, RX J0019.8+2156 and RX J0925.7-4758 \\citep[see references in][hereafter SL04]{SL04} and symbiotic systems are also known to blow jets \\citep[][and references therein]{sok1,Bro}. Also the fast nova V1494 Aql produced high velocity $\\sim 2800 \\rm \\ km\\ s^{-1}$ jets during its decline from outburst maximum \\citep{Iijima,retter04}. SSXSs, symbiotic systems and classical novae differ from CVs by much higher accretion rates and the presence of an additional source of energy produced by thermonuclear reactions. Although novae are CVs undergoing a thermonuclear runaway for the purpose of the present investigation we will separate them from the other members of this class of binaries. \\citet{Livio00} speculated that this latter difference might account for lack of jets in CVs, however, without providing a physical explanation.  Recently we (SL04) have shown that this difference is naturally explained if jets in accreting white-dwarf systems are produced by a thermal mechanism. As discussed below, magnetic fields would play a role in jet formation but this would be rather auxiliary, the main thrust being provided by thermal pressure. ",
        "conclusions": "The thermal model elucidates why there exists a critical accretion rate necessary for jet launching. This critical rate is well above the maximum rate encountered in CVs and so solves the mystery of their jet quietness. One could argue that this is just a coincidence. The critical value could, for example, correspond to the appearance of a large-scale (poloidal) magnetic field necessary for a \\citeauthor{BP82} -- type mechanism to work, as proposed by \\citet{Livio03}. According to these authors, at a critical rate the accretion disc would switch from a standard radiative disc to a state where most of the accretion energy is released in the form of a bulk flow. For the time being no MHD simulation is capable to follow such a process (see the article by Balbus in these proceedings), and the value of the critical accretion rate (if any) can be only matter of speculation. \\citet{Livio03} suggested that the inner parts of the nova-like star discs are underluminous because of the transition into an outflow phase. This would imply a critical accretion rate of $\\sim 10^{-8}\\rm M_\\odot \\rm yr^{-1}$. However, this rate corresponds rather to launching of winds, not of a jets, and the mystery of their absence in CVs would be still with us. It should be also noted that the alleged luminosity deficit in the inner disc of nova-like stars could be (at least in part) just an artefact of the disc model used \\citep{Smak94}. In any case the existence of a critical accretion rate for the presence of large-scale poloidal field would not be necessarily in contradiction with our criterion Eq. (\\ref{accwd2}) which is only a necessary condition for jet launching. On the other hand, if the poloidal field is generated at the cost of local energy dissipation thermal jet launching could be problematic.  Finally, our model has the vocation to be universal despite the fact that we use properties of the boundary layer which would not exist when accretion occurs onto black holes or strongly magnetized stars. However, the boundary layer is important only because it is where strong shocks can be produced. The required strong gradients could be also produced in accretion discs by ``magnetospheric\" MHD even when the central object is a black hole \\citep[e.g.][and references therein]{Linar}. In such a case MHD instabilities, turbulence, or other disturbances may lead to strong shocks; the high post-shock pressure may accelerate gas and form jets and/or winds, e.g., as was shown for non-radiative accretion around a black hole by \\citet{Devil03}.  Scaled to the case of an accreting black hole Eq. ({\\ref{acc01}}) becomes \\begin{equation} \\dot m \\equiv \\frac{\\dot M}{\\dot M_{\\rm Edd}} \\ga 0.02 \\left( \\frac{\\alpha}{0.1} \\right) \\left( \\frac{0.4 \\rm \\ g \\ cm^{-2}} {\\kappa}\\right) \\frac{R_j}{R_{\\rm G}}, \\label{comp} \\end{equation} where $\\dot M_{\\rm Edd}= L_{\\rm Edd}/0.1c^2= 2.3\\times 10^{-8} \\left(M/{\\rm M}_{\\odot}\\right) \\rm M_\\odot y^{-1}$ is the Eddington accretion rate. In \\citet{SL04} we mistakenly claimed that this formula might be relevant to the appearance of steady jets (we Friedrich Meyer for pointing this out at the present conference). It remains to be seen how and when the thermal jet-launching model applies to systems with black hole, but we expect \\citetext{in preparation} that in the case of microquasars it would be rather relevant to the launching of powerful, high Lorentz-factor jets \\citep[e.g.][]{fbg}."
    },
    "0410/astro-ph0410265_arXiv.txt": {
        "abstract": "\\label{abs} A quark-cluster state, rather than the color-super-conductivity state, may appear in matter with low-temperature but high density, since the phase transition of chiral symmetry broken and that of color-confinement could not occur simultaneously. Such quark clusters might be stable in strange quark matter. Quark stars would have temperatures to be low enough to freeze by residual color interaction between the clusters, and become then solid ones. The charge numbers and the sizes of quark clusters, as well as the residual interactions in-between are investigated. The solid state properties of quark stars are constrained by observations. ",
        "introduction": "\\label{intro}  Since E. Witten suggested that strange quark matter may be the true ground state of color interaction system \\citep{wit84}, it has been of great interest to find such matter in reality. There are hitherto many kinds of attempts, which can be classified into two ways. One is to detect directly strange quark matter on earth or in space (e.g., the ground and the AMS-II observations of strangelets in cosmic rays), the other is to study distant objects by means of astrophysics (e.g., to identify quark stars with strangeness). As for the nature of quark matter of strange stars, current researchers are focusing on quark condensation in momentum space (e.g., 2SC, CFL and LOFF states). Recently it has been proposed alternatively that quarks in strange quark matter with high density and low temperature may condense in position space (i.e., forming quark clusters which are {\\em not} color-confined), and may be in a solid state when the temperature is much low \\citep{xu03,xu04}. This kind of solid quark stars could be applied to explain well the discrepancy between the observed glitches and free-precessions of pulsars \\citep{xu04,zhou04}.  In the following sections, based on the idea of quark clustering, observational constrains on the nature of quark matter with quark clusters are derived. ",
        "conclusions": "We have obtained the upper limits of the quark number per cluster, $n$, from featureless spectra and pulsar glitches. The maximum value of $n$, $n_{\\rm max }$, constrained from featureless spectra is $\\sim 10^{4-8}$ for parameters of $50<m_s/{\\rm MeV}<200$, $0<\\alpha_s<0.6$, and reasonable bag constants $B$. The upper limit of $n$ for generating pulsar glitches is $10^5$ if bond-energy between clusters is 10MeV. We conclude therefore that the upper limit of quark number per cluster could be $\\sim 10^{4-5}$ in order for bare strange stars to reproduce both featureless X-ray spectra and pulsar glitches.  What have been done in this paper are based on the solid quark star model. This model is only a phenomenological model, into which it is very necessary to research from first principles. Nevertheless, the conclusion of our paper may provide helpful information about the detail processes of quark clustering and matter solidifying.  One problem in our paper is that we had used an equation of state derived from the traditional concept of free-quark matter, not clustered-quark matter. This problem may affect the result significantly if clustering changes the equation of state too much. Although this issue needs to be treated seriously, the deviation of equation of state may be slight, since color-confinement does not occur yet in the quark-clustering matter and the property of the vacuum could thus be not much different from that described by \\citet{alc86}."
    },
    "0410/astro-ph0410503_arXiv.txt": {
        "abstract": "We discuss an alternative approach to quintessence modifying the usual equation of state of the cosmological fluid in order to see if going further than the approximation of perfect fluid allows to better reproduce the available data. We consider a cosmological model comprising only two fluids, namely baryons (modelled as dust) and dark matter with a Van der Waals equation of state. First, the general features of the model are presented and then the evolution of the energy density, the Hubble parameter and the scale factor are determined showing that it is possible to obtain accelerated expansion choosing suitably the model parameters. We use the the data on the dimensionless coordinate distances to Type Ia supernovae and distant radio galaxies to see whether Van der Waals quintessence is viable to explain dark energy and to constrain its parameters. We then compare the model predictions with the estimated age of the universe and the position of the first three peaks of the anisotropy spectrum of the cosmic microwave background radiation. ",
        "introduction": "In the last few years an increasing bulk of data have been accumulated favouring the scenario of a spatially flat universe dominated by some form of dark energy. A first strong evidence came from the Hubble diagram of type Ia supernovae (hereafter SNeIa) that turned out to be best fitted by spatially flat accelerating cosmological models including a non trivial component \\cite{SNeIa}. On the other hand, the results from the observed first and second peak in the cosmic microwave background radiation (hereafter CMBR) spectrum strongly suggested that the geometry of the universe is spatially flat \\cite{Boom,CMBR}. When combined with the data on the matter density parameter $\\Omega_M$, these results lead to the conclusion that the contribution $\\Omega_X$ of the dark energy is the dominant one, being $(\\Omega_M, \\Omega_X) \\simeq (0.3, 0.7)$. This picture of the universe has been further strengthened by an increasing sample of high redshift SNeIa \\cite{NewSNeIa,Riess04} and most precise and extended measurements of the CMBR spectrum \\cite{WMAP,VSA}.  After the discovery of these evidences of a spatially flat and accelerating universe, an overwhelming flood of papers, presenting a great variety of models for the explanation of this phenomenon, has appeared. The simplest explanation is the cosmological constant \\cite{Lambda} which is able to fit the SNeIa data with good confidence, but it is also plagued by many problems on different scales. This situation has strongly encouraged the search for alternative approaches which now ranges from minimal coupled scalar fields, to strings and anthropic principle (see, e.g., \\cite{QuintRev} and references therein).  In a recent work \\cite{CdMF}, some of us have introduced a new approach to the problem modifying the usual equation of state of the cosmological fluid in order to see whether going further than the standard  approximation of cosmological perfect fluid allows to  reproduce the available data. The authors have considered a standard cosmology with a Van der Waals equation of state for matter without any other kind of energy source. This approach is extremely \"natural\" and \"obvious\" since it starts from the consideration that the universe, in its evolution, is not always well described by a perfect fluid in the forms of radiation or non\\,-\\,interacting dust. In our opinion, before adding an exotic and mysterious dark energy into the cosmic pie, it is worth wondering whether the observed acceleration of the universe can be implemented using the minimal number of fluids\\footnote{Following Newton {\\it ``Hypotheses non fingo''}.}. The price to pay is taking into account a more complicated equation of state and motivate it physically. From elementary thermodynamics, we know that a real fluid is never perfect \\cite{Rowlinson}. Moreover, it is also well known that the perfect fluid equation of state $p = \\gamma \\rho$ with $\\dot{\\gamma} = 0$ is just a rough approximation of cosmic epochs capable of describing stationary situations where phase transitions (e.g., from radiation dominated to dust dominated regions) are not considered \\cite{E93+EvE98}. On the other hand, the only thing we know about dark energy is that it gives rise to an accelerated expansion, but there are no hints about its nature. Due to this situation, it is worth asking whether dark energy is indeed needed or, on the contrary, the observed acceleration is driven by standard dark matter provided that its equation of state is more realistically treated.  A Van der Waals fluid could be a first step toward the goal to get a whole dynamics where {\\it 1.)} only observed fluids are taken into account, {\\it 2.)} phase transitions occur in the framework of the same evolution, {\\it 3.)} accelerated and decelerated periods depend on the relative values of the parameters of the state equation with respect to the pressure and matter energy density which are functions of time. It is worth noting that a similar approach has yet been explored in \\cite{Kr03} where the author considers a mixture of two fluids, using the perfect gas equation of state for the matter and the Van der Waals one for the dark energy. However, our approach is radically different since the model we consider is made out of matter only. Is is worth noting that the term {\\it matter} usually refers to both baryons and dark matter and both these substances are described by the same equation of state. Actually, we have a direct knowledge of the properties of baryons only and indeed their equation of state is well described (on cosmological scales) by the dust approximation. On the opposite, the nature of the dark matter is still completely unknown so that, a priori, nothing prevents us from exploring the possibility that its properties call for a more general equation of state such as the Van der Waals one. The aim of the present paper is to explore further this approach in order to see whether a cosmological fluid with Van der Waals state equation can be reconciled with observations. We have thus constrained the effective parameters of the theory looking for cosmological models which: $i)$ can admit a nowaday accelerating universe; $ii)$ satisfy the constraints on the estimated age of the Universe, $iii)$ are able to fit the data on the dimensionless coordinate distance to SNeIa and radio galaxies.  The paper is organized as follows. In Sect.\\,2, we briefly summarize the main features of Van der Waals equation of state and choose the set of parameters which are best suited to assign the model. Dynamics of the model is described in Sect.\\,3 where we determine the evolution of the energy density, the Hubble parameter and the scale factor. Matching with observations is performed in Sect.\\,4 where the data on the dimensionless coordinate distance to SNeIa and radio galaxies are used to select among the models. In Sect.\\,5, we compare the model prediction for the age of the universe with the estimated one, while, in Sect.\\,6, we evaluate the position of the first three peaks of the CMBR anisotropy spectrum and qualitatively speculate on how structure formation could take place in the Van der Waals quintessence scenario. Sect.\\,7 is then devoted to the discussion of the results and conclusions. ",
        "conclusions": "The increasing bulk of astrophysical data accumulated in recent years has delineated a new standard cosmological paradigm. According to this picture, the universe is spatially flat and driven by an unknown form of dark energy leading to an accelerated expansion. Soon after the establishment of such a scenario, the hunt for candidates to the dark energy throne has started leading to the proposal of a plethora of mechanisms ranging from the old cosmological constant to various scalar field quintessence, to modification of Friedmann equations (motivated by extradimensions and braneworld theories) and higher order geometrical terms in the gravity Lagrangian \\cite{fR1,fR2,fR3}. Although being completely different in their dynamical properties and underlying physics, they all share the ability of well fitting the same set of data.  Furthermore, it is worth noting that the dark side of the universe is not only populated by dark energy, but also by dark matter whose nature is still far to be understood. It is thus tempting to ask whether these two ingredients are indeed different substances or two aspects of a single fluid whose properties are different from what we usually expect. To this regard, it makes sense wondering if the dust approximation used for matter is the right one to describe the equation of state of the dark matter. In a tentative to explain the observable quantities of our universe with the minimal number of ingredients, here we have investigated the possibility that dark matter and dark energy are actually a single fluid whose equation of state is that of a Van der Waals gas. In Van der Waals quintessence scenario, there is a single fluid whose equation of state comes directly from classical thermodynamics since the perfect gas approximation cannot be used to describe phase transitions which occur during the evolution of the universe. Although their contribution to the energy budget is nowaday subdominant, we also include baryons in the model since, because of the nonlinear character of the dynamical equations, they play an important role to determine how the main quantities evolve with the redshift.  Any theory, as elegant and motivated it can be, is meaningless if it is unable to give a coherent description of the universe as it is observed. That is why we have tested Van der Waals quintessence against the dimensionless coordinate distance to SNeIa and radio galaxies. This allows to narrow the parameter space of the model considering the present day values of the deceleration parameter and of the ratio between the critical density of the universe and the Van der Waals critical density. We find\\,:  \\begin{eqnarray} -0.52 \\le q_0 \\le -0.33 \\ \\ , \\ \\ -0.42 \\le \\log{\\eta_0} \\le -0.11 \\ \\ & \\ \\ {\\rm at \\ 68\\% \\ CL} \\ \\ , \\nonumber \\\\ ~ \\nonumber \\\\ -0.58 \\le q_0 \\le -0.28 \\ \\ , \\ \\ -1.38 \\le \\log{\\eta_0} \\le -0.05 \\ \\ & \\ \\ {\\rm at \\ 95\\% \\ CL} \\ \\ , \\nonumber \\end{eqnarray} with $(q_0, \\log{\\eta_0}) = (-0.38, -0.21)$ as best fit parameters. The constraints on $\\log{\\eta_0}$ are weak because of a serious degeneracy among the model parameters. Indeed, different models of this class predicts values of $y(z)$ which are in agreement with each other within the observational errors. However, lowering the uncertainties on $y(z)$ in the redshift range $(1.0, 2.0)$ or extending observations to higher redshifts (up to $z \\sim 3 - 4$) will allow to break this degeneracy. Both these possibilities are likely to become realistic in a near future thanks to the next\\,-\\,to\\,-\\,come satellite and ground based experiments.  Motivated by this encouraging result, we have compared the predicted age of the universe with that estimated from globular clusters in a model independent way. Unfortunately, it turns out that only a quite small region of the parameter space delimited by the above constraints is allowed by this test. Moreover, $t_0$ turns out to be systematically high so that one could argue against the proposed model. Actually, this conclusion is premature since our ignorance of the thermodynamical properties of the dark fluid makes it impossible to determine the redshift range over which Eq.(\\ref{eq: vdwour}) holds. Indeed, during the evolution of the universe, the dark fluid may undergo phase transitions thus changing its equation of state in a unpredictable way. Since $t_0$ depends on the full evolutionary history of the universe, using Eq.(\\ref{eq: vdwour}) at all $z$ may induce a systematic error if the relation $p = p(\\rho)$ is altered. Given these theoretical uncertainties, overemphasizing the result of the age test should be avoided.  As a further test, we have also evaluated the position of the first three peaks of the CMBR anisotropy spectrum using the analytical approximation developed in Refs.\\,\\cite{DorLil01,DorLil02}. Comparing with the measured $l_m$, we have individuated a region of the parameter space $(q_0, \\log{\\eta_0})$ which is consistent with both the distance fit and the peaks position. Although very encouraging, this result should be confirmed by evaluating the full spectrum and not the peaks position only. To this aim, one has to first investigate in more detail how perturbations grow in the Van der Waals quintessence scenario. Because of the negative adiabatic sound speed, perturbations go unstable and one has to take into account also the role of anisotropic stresses and entropy production.  We would like to conclude with a general comment. Van der Waals quintessence has turned out to be an interesting scenario for describing the late universe and seems able to solve the puzzle of dark energy without adding exotic fluids or arbitrary modifications of Friedmann equations. On the other hand, classical thermodynamics tells us that the Van der Waals equation of state is only an approximated description of a realistic fluid. In our opinion, before invoking the help of {\\it new physics}, it is worth wondering whether {\\it classical physics} could still suggest the way to shed light on the dark side of the universe.  \\ack  It is a pleasure to thank our friend Monica Capone for the fruitful discussions on this topic and Ester Piedipalumbo and Crescenzo Tortora for help with the position of the peaks. We are also grateful to an unknown referee for his stimulating report."
    },
    "0410/astro-ph0410029_arXiv.txt": {
        "abstract": "{A serious candidate for dark matter in spiral galaxies are cold molecular hydrogen globules with a condensed central core and a disc-like space distribution probably similar to that of neutral hydrogen. This paper shows that the H$_2$ cores are sufficiently compact and massive to be detected by microlensing in the outer Galactic disc and that the Maffei~1 elliptical galaxy, at a distance of 3~Mpc and Galactic latitude $b=-0.6^{\\circ}$, offers an ideal target for such an experiment. The microlensing optical depth of H$_2$ cores along the line of sight to this galaxy is estimated to $\\tau\\sim 0.7\\cdot 10^{-6}$ if most of the dark mass in the Milky Way resides in such cores, and the typical event timescale to $\\la 1$~day. Detection rates are computed both in the classical and pixel lensing approaches in the $I$- and $K$-bands, and for a representative selection of existing observing facilities. In the more efficient pixel lensing case, two 10-hour observing runs, separated in time by at least several days, should yield of order 10~positive detections at the $5\\sigma$ level using ground-based 8m-class telescopes in the $K$-band or the Hubble Space Telescope~ACS camera in the $I$-band, and the corresponding fraction of events with timescale measurable to an accuracy better than $50\\%$ amounts to about $9\\%$ and $4\\%$ respectively for these observing alternatives. ",
        "introduction": "The nature of dark matter (DM) in the universe is one of the main persistent problem in modern astrophysics. In spiral galaxies, dark matter is highlighted by flat HI rotation curves extending far beyond the optical disc, where one would expect a Keplerian fall-off (e.g. Rogstad \\& Shostak~\\cite{RS}). Many DM candidates have been proposed in the literature, including massive compact objects, neutrinos and other more exotic particles. Compact objects have been and are still searched for in the local group using classical or pixel microlensing of stellar sources located in the Galactic bulge, the Magellanic Clouds and Andromeda (e.g. the OGLE, MACHO, EROS, MOA, AGAPEROS, MEGA, POINT-AGAPE, WeCAPP collaborations\\footnote{See respectively Wozniak et al.~\\cite{OGLE}, Alcock et al.~\\cite{MACHO}, Afonso et al.~\\cite{EROS}, Bond et al.~\\cite{MOA}, Melchior et al.~\\cite{AGAP}, de Jong et al.~\\cite{MEGA}, Paulin-Henriksson et al.~\\cite{POINT}, and Riffeser et al.~\\cite{WECAPP} for selected references.}). All these experiments are designed to find dark objects in the inner Galaxy or in the more or less spherical halos around the Milky Way and Andromeda galaxies. \\par An attractive possible constituent of galactic DM is H$_2$ in the form of cold ``globules'', made of a condensed solid or liquid central core surrounded by an extended atmosphere (Pfenniger~\\cite{P}; Pfenniger \\& Combes~\\cite{PC}; see also White~\\cite{RSW}). These H$_2$ cores would have radii of up to $2\\cdot 10^4$~km and masses up to the Earth mass. A number of constraints, like the \\hbox{HI--DM} relation (Bosma~\\cite{BA1}), the disc-halo conspiracy, the maximum disc property implying hollow spherical halos, and the large angular momentum of accreted high velocity clouds, do support a space distribution of H$_2$ globules similar to the HI distribution, i.e. as an outer disc around spiral galaxies. Evidence for large amounts of molecular gas in such regions is also provided by the observed star formation in the extreme outskirts of the HI disc in M31 (Cuillandre et al.~\\cite{CLA}). Moreover, \\hbox{$N$-body} simulations show that massive collisionless DM discs are prone to bending instabilities leading to long lasting warps (Revaz \\& Pfenniger 2004), consistent with the observed high frequency of these structures. \\par This paper shows that if H$_2$ cores contain most of the Milky~Way's DM and have a disc-like distribution, they should be quite easily detectable through a short microlensing experiment targeting the low Galactic latitude Maffei~1 galaxy. \\par The structure of the paper is as follow. Section~\\ref{maf} summarises the main observational characteristics of Maffei~1. Section~\\ref{basic} briefly outlines the microlensing aspects useful to our investigation and evaluates the basic properties resulting from the lensing of extragalactic sources by H$_2$ cores. Section~\\ref{lum} introduces the concept of fluctuation magnitude and attempts to derive a realistic stellar luminosity function for Maffei~1. Section~\\ref{rate} then computes the expected event detection rates for a representative set of available observational facilities, both in the classical and pixel lensing workframes. Finally, Sect.~\\ref{constr} argues that a massive Galactic H$_2$-globule disc is not excluded by current microlensing constraints, and Sect.~\\ref{concl} concludes the paper. The problem of contamination by variable stars is beyond the scope of this paper and will not be treated, and the case of more diffuse H$_2$ clouds is analysed in Rafikov \\& Draine~(\\cite{RD}). Much of the notations and theory elements are based on Gould~(\\cite{G96}; hereafter G96). \\begin{figure}[t!] \\centerline{\\includegraphics[width=7cm]{fig1.ps}} \\caption{\\small $JHK$ composite image of Maffei~1 from the 2MASS Large Galaxy Atlas (Jarrett et al.~\\cite{JCC}). The field of view is $11.9'\\!\\times\\! 11.9'$.} \\label{obs} \\end{figure} ",
        "conclusions": "\\label{concl}  The theory of cold H$_2$ globules with condensed central cores to account for dark matter in spiral galaxies can be readily tested by microlensing. If such globules exist in the Milky Way, the former microlensing experiment have looked at directions which do not optimise their detection rates. These globules are indeed expected to follow the HI distribution and should be concentrated in the outer Galactic disc, where the contribution of invisible mass is highest. The large galaxies in the Maffei group provide excellent microlensing targets to probe the H$_2$-core content of this region. In particular, the Maffei~1 elliptical is at a distance of only 3~Mpc, very close to the Galactic plane and $\\sim\\! 44^{\\circ}$ away from the Galactic anti-centre, and offers a huge reservoir of microlensing sources. \\par The condensed H$_2$-globule cores have mass and size comparable to the Earth or less. The Earth-mass Galactic H$_2$ core--Maffei~1 lensing geometry implies an Einstein radius of $\\sim\\! 3\\;R_{\\odot}$, large enough to consider the source stars as point-like over a large range of impact parameters. Assuming that the H$_2$ cores are distributed with twice the HI mass density, the microlensing optical depth towards Maffei~1 of these cores is $\\tau\\approx 0.7\\cdot 10^{-6}$, and in the circular orbit approximation, the typical timescale of the lensing events is at most $\\sim\\! 1$~day. \\par In the classical lensing approach, due to the fluctuation noise, the surface density of resolvable stars in Maffei~1 as a function of apparent major-axis radius $a$ peaks away from the centre, at $a\\approx 230$~arcsec in the $I$-band and at $a\\approx 90$~arcsec in the $K$-band, and high resolution is crucial to maximise the amplitude of this peak. Because of the smaller Maffei~1 area with high density of resolvable stars and the restricted detector size in the $K$-band, and of the poor seeing in ground-based optical observations, the ACS camera on HST is currently the only useful instrument. In the $I$-band and monitoring several fields at the same time, it would take $\\sim\\! 5$ days of ACS observations to detect one event. \\par The pixel lensing approach, however, reveals over one order of magnitude more efficient. In this case, the detection rate is maximum in the central region ($a=0$) and it is best to rely on a single field centred at this position. Since pixel lensing involves the subtraction of reference images to the images with the ongoing lensing events, at least two observing runs are required, spaced by a time interval large relative to the lensing timescale, to get the two sets of images. The roles of the reference and event images can fortunately be interchanged to double the number of detectable events. We find that in the $K$-band and resorting to adaptive optics, the ground-based 8m-class telescopes could detect as much as 10 events at the $5\\sigma$ level in two nights, and that in the $I$-band, the ACS camera could detect a similar number of events at the same threshold in two 10-hours observing periods. \\par About $9\\%$ of the events detected by 8m-class telescopes in $K$ should have timescales measurable at a precision better than $50\\%$, representing about one such timescale per couple of nights, while this fraction reduces to $\\sim\\! 4\\%$ for the ACS observations in $I$. A potentially powerful way to constrain the typical lens mass $m_{\\rm L}$ could be to measure the event detection rate as a function of the radius~$a$. Indeed, the average impact parameter of detectable events decreases with increasing galaxy surface brightness, and the finite radius of the source stars imposes a lower limit $\\beta_{\\rm min}\\propto m_{\\rm L}^{-1/2}$ below which the true flux amplification becomes much less than the point-source approximation. As a consequence, the detection rate curve will display a maximum at a radius related to the lens mass, with precise location depending on the signal-to-noise threshold. \\par The advent of instruments like WEBCam on JWST in the near-IR, with a field of view of $2.16'\\times 2.16'$, will boost the detection numbers derived here. Although a positive detection will not yet prove that the intervening lens is an H$_2$ core, it will certainly represent a significant step in understanding the origin of dark matter, whereas a no-detection result may refute the cored H$_2$ globules as a representative mass constituent of the Galaxy."
    },
    "0410/astro-ph0410735_arXiv.txt": {
        "abstract": "Different interaction models treating hadronic collisions below $E_{\\rm lab}$$<$$100$ GeV in air shower simulations are compared within the frame of the CORSIKA program. Their behaviour is studied in isolated collisions of protons or pions with light-mass targets in comparison with experimental data. It is discussed how model properties influence air-shower parameters like the longitudinal profile or the lateral distribution of particles arriving at ground. \\vspace{1pc} ",
        "introduction": "In extensive air showers (EAS) induced by cosmic hadronic particles low-energy secondary hadrons collide with the atmospheric nuclei thus forming the final branches of the hadronic shower backbone. Recently Ref.~\\cite{drescher} reported that in simulations the lateral particle distribution at core distances $>$1 km shows a strong dependence on the used low-energy hadronic interaction model. Therefore several codes describing low-energy interactions have been examined with the EAS simulation program CORSIKA \\cite{corsika} in the present study. Preliminary results are given in \\cite{heck_tsukuba}. The used interaction codes are presented in Sect.~\\ref{sec-models}. In Sect.~\\ref{sec-sing_coll} their predictions are compared with experimental results of single hadronic collisions with light targets close in mass to the atmospheric constituents. How the behaviour in single collisions transforms into measurable properties of EAS is discussed in Sect.~\\ref{sec-shower_sim}. The consequences to simulations are exemplified in Sect.~\\ref{sec-implicat}. ",
        "conclusions": "\\label{sec-implicat}  The CPU-time requirements to treat single collisions rsp.~complete showers is important for EAS-simulations. In Table \\ref{tab-cpu} CPU-times are collected.  The HSA-code as the fastest has an unsatisfactory performance as demonstrated in the preceding sections and lacks predictive power. A revision of its parameters is needed.  As shown in previous sections the kinematics of GHEISHA 2002 is still not satisfying. \\begin{table}[b] \\vspace{-12pt} \\begin{center} \\caption{CPU times for DEC-alpha 1000XP (500 MHz).} \\label{tab-cpu} \\begin{tabular}{l|r|r|r} \\hline &\\multicolumn{2}{c|}{$10^5$ collisions}& \\multicolumn{1}{c}{1 shower $^1$} \\\\[-1pt] & \\multicolumn{2}{c|}{at 10 GeV}& \\multicolumn{1}{c}{$10^{19}$ eV}\\\\[-1pt] $~~~$Model& p-air        & $\\pi$-air & \\multicolumn{1}{c}{$\\varepsilon=10^{-6}$}\\\\ \\hline FLUKA        &        181   &   164           & 63300 \\\\ GHEISHA$~$2002&       108   &   102           & 29100 \\\\ HSA (AIRES)  &         64   &    64           & 18000 \\\\ UrQMD 1.3 $^2$ &    12200   & 11400           &($\\approx$800000)\\\\ \\hline {\\sc neXus 3}&       6173   &  5861           &           \\\\ QGSJET$~$01  &         88   &    87           &            \\\\ \\hline \\multicolumn{4}{l}{\\scriptsize $^1$ QGSJET 01, $\\theta = 0^\\circ$, $E_{\\rm h}>300$MeV, $E_{\\rm \\mu}>100$MeV,} \\\\[-2pt] \\multicolumn{4}{l}{\\scriptsize $~~$ $E_{\\rm em}>250$keV, W$_{\\rm em}<10^4$, W$_{\\rm h}<10^2$.}\\\\[-1pt] \\multicolumn{4}{l}{\\scriptsize $^2$ H.J. Drescher has accelerated UrQMD 1.3c by a factor 15.} \\\\[-2pt] \\end{tabular} \\end{center} \\vspace*{-21pt} \\end{table}  The microscopic UrQMD model gives reliable results at the expense of extreme long CPU-times. In this model no parametrized cross sections are used, rather the projectile hits on a `disk' chosen large enough to cover the maximum collision parameter. Several trials (with their time-consuming microscopic calculation) without interaction are  usually needed before in an inelastic collision secondary particles are produced. Therefore this model is less suited for EAS simulations.\\\\[3pt] The FLUKA model shows the best performance and it is recommended for EAS-simulations despite the longer CPU-times. Unfortunately its source code is not yet publicly available, only object codes for 5 types of machine/operating systems are distributed at present.  Finalizing one can state that measurable shower parameters are predominantly influenced by the high-energy interaction programs \\cite{heck_cern} and only to some minor extend by the low-energy models."
    },
    "0410/gr-qc0410128_arXiv.txt": {
        "abstract": "In this paper, a braneworld black hole is studied as a gravitational lens, using the strong field limit to obtain the positions and magnifications of the relativistic images. Standard lensing and retrolensing situations are analyzed in a unified setting, and the results are compared with those corresponding to the Schwarzschild black hole lens. The possibility of observing the strong field images is discussed. ",
        "introduction": "Gravitational lensing by ordinary stars and galaxies can be analyzed in the weak field approximation, i.e., only keeping the first non null term in the expansion of the deflection angle \\cite{schne}. If the lens is a black hole, this approximation is only valid for photons with large impact parameter, so a full strong field treatment is needed instead. In the general case where the lens is a compact object with a photon sphere, besides the primary and secondary weak field images, two infinite sets of faint relativistic images are formed by photons that make complete turns (in both directions of rotation) around the black hole before reaching the observer. In the last few years, several works studying different strong field lensing scenarios appeared in the literature. Virbhadra and Ellis \\cite{virbha1} made a numerical analysis, using an asymptotically flat metric, of the case where the lens is a Schwarzschild black hole situated in the center of the Galaxy, and in another paper \\cite{virbha2} they investigated numerically the lensing by naked singularities. Fritelli, Kling and Newman \\cite{fritelli} found an exact lens equation without any reference to a background metric and compared their results with those of Virbhadra and Ellis. In the Schwarzschild geometry, several authors \\cite{prev} used a logarithmic approximation of the deflection angle as a function of the impact parameter, for light rays passing very close to the photon sphere, to treat strong field situations. This asymptotic approximation is the starting point of an analytical method for strong field lensing, called the strong field limit \\cite {bozza1}, which gives the lensing observables in a straightforward way. Eiroa, Romero and Torres \\cite{eiroto} extended this method to Reissner-Nordstr\\\"{o}m geometry, and Bozza \\cite{bozza2} showed that it can be applied to any static spherically symmetric lens. It was subsequently used by Bhadra \\cite{bhadra} to study a charged black hole lens of string theory, and by Petters \\cite{petters} to analyze the relativistic corrections to microlensing effects produced by the Galactic black hole. Bozza and Mancini \\cite{bozman1} applied the strong field limit to study the time delay between different relativistic images, showing that different types of black holes are characterized by different time delays, and  Bozza \\cite{bozza3} extended the strong field limit to analyze the case of quasi-equatorial lensing by rotating black holes.\\\\ In standard lensing situations, the lens is placed between the source and the observer. When the lens has a photon sphere and the observer is placed between the source and the lens, or the source is situated between the lens and the observer, two infinite sequences of images with deflection angles closer to odd multiples of $\\pi$ are obtained, a situation called retrolensing. Holtz and Wheeler \\cite{holtz} recently analyzed the two stronger images for a black hole situated in the galactic bulge with the sun as source, and suggested retrolensing as a new mechanism for searching black holes. De Paolis et al. \\cite{depaolis1} considered the retrolensing scenario of the bright star S2 orbiting around the massive black hole at the galactic center. Eiroa and Torres \\cite{eitor} studied the case of a spherically symmetric retrolens, using the strong field limit to obtain the positions and magnifications of all images.  De Paolis et al. \\cite{depaolis2} extended, using the strong field limit, the work of Holtz and Wheeler to slowly spinning Kerr black hole, restricting their treatment to the black hole equatorial plane. Bozza and Mancini \\cite{bozman2} analyzed, in the strong field limit, standard lensing, retrolensing and intermediate situations under a unified formalism, and studied in detail the case of the star S2, suggesting the possibility of observing the relativistic images in the year 2018. For other works that considered related topics on strong field lensing, see Ref. \\cite{other}.\\\\  Braneworld cosmologies, where the ordinary matter is on a three dimensional space called the brane, embedded in a larger space called the bulk in which only gravity can propagate, became popular in the last few years \\cite{bwrev}. These models, proposed to solve the hierarchy problem, i.e., the difficulty in explaining why the gravity scale is 16 orders of magnitude greater than the electro-weak scale, have motivation in recent developments of string theory, known as M-theory. The presence of extra dimensions would affect the characteristics of black holes \\cite{kanti}. The possibility of the existence of primordial black holes in the simplest of braneworld scenarios, the Randall-Sundrum \\cite{rsII} models (a positive tension brane in a bulk with a negative cosmological constant) with one extra dimension, have been studied by Clancy, Guedens and Liddle \\cite{cgl}. They showed that black holes formed in the high energy epoch of this theory have a longer lifetime, due to a different evaporation law. Majumdar \\cite {majumdar} found that the primordial black holes could have a growth of their mass through accretion of surrounding radiation during the high energy phase, increasing their lifetime. These black holes could have survived up to present times and have an induced four dimensional metric on the brane distinct from Schwarzschild metric. They also may be created in high energy collisions in particle accelerators and in cosmic rays \\cite{kanti}. The possibility that these primordial black holes could act as gravitational lenses was analyzed by Majumdar and Mukherjee \\cite{majumuk}. They only considered the case of photons with small deflection angles in the standard lensing configuration. In another braneworld model Frolov, Snajdr and Stojkovic \\cite{frolov} calculated, using a weak field approximation, the deflection of light propagating in the brane produced by a small size black hole in the bulk.\\\\  In this paper, the strong field limit is applied to study the relativistic images produced by the braneworld black hole analyzed by Majumdar and Mukherjee \\cite{majumuk} only for the weak field situation. Standard lensing and retrolensing are considered for the case of high alignment between the source, the lens and the observer, which is the case where the images are more prominent. In Sec. \\ref{s2}, the deflection angle is calculated for the braneworld black hole. In Sec. \\ref{s3}, the positions and magnifications of the relativistic images are obtained, and in Sec. \\ref{s4} the results are compared with those corresponding to a four dimensional Schwarzschild black hole. Finally, in Sec. \\ref{s5}, some conclusions are shown. Units such that $c=\\hbar =1$ are used throughout this work. ",
        "conclusions": "\\label{s5}  In this work, the positions and magnifications of the relativistic images were calculated, using the strong field limit, for a braneworld black hole. These black holes were previously studied as primordial black holes \\cite{cgl,majumdar}, and can be relics of a past high energy phase of the Universe. They have a larger size, are colder, and live longer than their four dimensional Schwarzschild counterpart of the same mass \\cite{kanti}. Current experiments that test the gravitational inverse square law, lead to an upper limit for the extra dimension given by $l\\le 0.1$ mm \\cite{long}. The condition that the event horizon radius $r_{H}$ is much smaller than the Ads radius $l$ for the braneworld black holes studied in this paper, means that these black holes, if they exist, are very small. The astronomical observation of the relativistic images is a difficult task that is beyond current technologies, and it will be a challenge for the next generation of instruments in the case of more massive black holes \\cite{bozman2}. The observation seems even more difficult in the case of the small size braneworld black holes analyzed in this paper. But there could be another route for seeing the strong field lensing effects for these low mass black holes. The presence of the extra dimension dramatically decreases the energy necessary to produce black holes by particle collisions \\cite{kanti,lab}. Only energies of about $1$ TeV are needed instead of energy scales about $10^{16}$ TeV required if no extra dimensions are present. These small size black holes could be created in the next generation particle accelerators or detected in cosmic rays \\cite{kanti,lab}. With these black holes acting as gravitational lenses, the order of magnitude of the distances involved would be meters or less, instead of kilo-parsecs, so that the angular positions and magnifications obtained in Sec. \\ref{s3} will be considerably larger than in the astronomical case corresponding to black holes with the same mass \\footnote{For example, let us consider, in the standard lensing configuration, a black hole lens with $r_{H}=10^{-3}$ mm placed halfway between a point source and the observer in two cases: a laboratory (L) situation with $D_{os}^{(L)}=1$ m and an astronomical (A) situation with $D_{os}^{(A)}=10$ kpc. Then, the quotients between the first Einstein angular radii and the total magnifications are, respectively, $\\theta^{E\\,(L)}_{1}/\\theta^{E\\,(A)}_{1}=3.1\\times 10^{20}$ and $\\mu^{(L)}/\\mu^{(A)}=9.5\\times 10^{40}$.}. If the extra dimensions hypothesis is correct, it will open up the exciting possibility that the interesting phenomena of strong field lensing can be observed in the laboratory in future times."
    },
    "0410/astro-ph0410445_arXiv.txt": {
        "abstract": "{ We use some of the recently released observational data to test the viability of two classes of minimally coupled scalar field models of quintessence with exponential potentials for which exact solutions of the Einstein equations are known. These models are very sturdy, depending on only one parameter - the Hubble constant. To compare predictions of our models with observations we concentrate on the following data: the power spectrum of the CMBR anisotropy as measured by WMAP, the publicly available data on type Ia supernovae, and the parameters of large scale structure determined by the 2-degree Field Galaxy Redshift Survey (2dFGRS). We use the WMAP data on the age of the universe and the Hubble constant to fix the free parameters in  our models. We then show that the predictions of our models are consistent with the observed positions and relative heights of the first 3 peaks in the CMB power spectrum, with the energy density of dark energy as deduced from observations of distant type Ia supernovae, and {\\ bf with} parameters of the large scale structure as determined by 2dFGRS, in particular with the average density of dark matter. Our models are also consistent with the results of the Sloan Digital Sky Survey (SDSS). Moreover, we investigate the evolution of matter density perturbations in our quintessential models, solve exactly the evolution equation for the density perturbations, and obtain an analytical expression for the growth index $f$. We verify that the approximate relation $f\\simeq \\Omega_{{\\rm M}}^{\\alpha}$ also holds in our models. ",
        "introduction": "Recent observations of the type Ia supernovae and CMB anisotropy strongly indicate that the total matter-energy density of the universe is now dominated by some kind of dark energy or the cosmological constant $\\Lambda$ (\\cite{rie+al98,Riess00,per+al99,Riess04}). The origin and nature of this  dark energy remains unknown~(\\cite{zel67,weinberg2,car}).\\\\  In the last several years a new class of cosmological models has  been proposed. In these models the standard cosmological constant  $\\Lambda$-term is replaced by a dynamical, time-dependent  component - quintessence or dark energy - that is added to  baryons, cold dark matter (CDM), photons and neutrinos. The equation of state of the dark energy is given by $w_Q \\equiv \\rho_Q /p_Q$, $\\rho_Q$ and $p_Q$ being, respectively, the pressure and energy density, and $-1 \\leq w_Q <0$, which implies a negative contribution to the total pressure of the cosmic fluid.  When $w_Q =-1$, we recover a constant $\\Lambda$-term. One of the possible physical realizations of quintessence is a cosmic scalar field, minimally coupled to the usual matter action  (\\cite{pe+ra88,cal+al98}). Such a field induces dynamically a repulsive  gravitational force, causing  an accelerated expansion  of the Universe, as recently discovered by observations of distant  type Ia supernovae (SNIa) (\\cite{per+al99,rie+al98, Riess04}) and confirmed by WMAP observations (\\cite{spergel}). Accelerated expansion together with the strong observational evidence that the Universe is spatially flat ( \\cite{deb+al00, spergel}) calls for an additional component and quintessence could be responsible for the missing energy in a flat Universe with a subcritical matter density. Quintessence drives the cosmological expansion at late times and also influences the growth of structure arising from gravitational instability. Dark energy could cluster gravitationally only on very large scales ($\\geq 100$ Mpc), leaving an imprint on the microwave background anisotropy (\\cite{cal+al98}); on small scales, fluctuations in the dark energy are damped and they do not influence the evolution of perturbations in the pressureless matter ( \\cite{ma+al99}). On the scales we are considering in the following, we assume that quintessence behaves as a smooth component, so that in our analysis the formation of clusters is due only to matter condensation, while the quintessence alters only the background cosmic evolution. The leading candidates  for the dark energy as suggested by fundamental physics include vacuum energy, a rolling scalar field, and a network of slight topological defects~(\\cite{tu00,cl01}).  Moreover, an eternally accelerating universe seems to be at odds with some formulations of the string theory~(\\cite{fischler01}). However, this is still controversial. Actually in the last few years it has been suggested that string theory could be compatible with the presently considered cosmological models (see for example, ~\\cite{hell01,Townsend03,gibbons01} and references therein). This has stimulated a revival of interest in the exponential scalar field quintessence. In different scenarios, such exponential potentials can reproduce the present accelerated expansion of the universe and some predict future deceleration. Moreover,  despite the criticism that exponential potentials require fine tuning, recently several authors  have pointed out that the degree of fine tuning needed in this case is not greater than in other scenarios~(\\cite{cl01,ru+sc01,card1}). \\\\    In this work we show that the two models of quintessence for which general exact solutions of the Einstein equation are known are compatible with recent  observational data, in particular with the power spectrum of the temperature anisotropy of the CMBR, the observations of type Ia  supernovae, and the parameters of the large scale structure of matter distribution. As a first step we introduce a new parametrization of quintessence models considered by Rubano et al. and Rubano \\& Scudellaro ~(\\cite{mc,ru+sc01}), which avoids the problem of branching of solutions highlighted in~Cardenas \\& al., (2002). \\\\ Since the physical features of these models will be extensively discussed in a forthcoming paper, in Sect. 2, we simply present the basic equations of the quintessence models used in this paper. In  Sect. 3 the linear perturbation equation is solved for the two potentials and the growth of density perturbations is discussed; in  Sect. 4 we show that predictions of our models are compatible with the observationally established power spectrum of CMB anisotropy, the SNIa data, and the results of estimates of the average mass density of the universe from galaxy redshift surveys. In  Sect. 5 we discuss the possibility of constraining the equation of state of dark energy by age estimates of the universe. Sect. 6 is devoted to final conclusions. ",
        "conclusions": "We tested the viability of  two classes of cosmological models with scalar field models of quintessence with  exponential potentials. To compare predictions of our models with observational  data we used  three different types of observations: namely, the power spectrum of the CMB temperature anisotropy and in particular positions and heights of the peaks, the high-redshift SNIa data  compiled in Riess \\& al. (2004) (gold data set), and the large survey of galaxies by the 2dFGRS team (\\cite{haw02}) and in particular their estimate of the average density of dark matter.  We showed that predictions  of our models  \\begin{table*}[ht] \\begin{center} \\caption{The basic cosmological parameters derived from our models are compared with those presented by the  WMAP team. The position of the third peak $l_3$ has been determined by the Boomerang team. The last set of  errors in the case of the double exponential potential quantifies the effect of the parameter $\\omega$, when it changes from $\\omega=0.1$ to $\\omega=0.9$. }  \\label{table-1_CMB} \\vspace{0.5cm} \\begin{tabular}{|c|c|c|c|} \\hline & $  \\exp\\left\\{ -\\sqrt{3\\over 2}\\varphi\\right\\}$ &  $ \\left( A\\exp{({1\\over 2}\\sqrt{{3\\over2}\\varphi})}-B \\exp{(-{1\\over 2}\\sqrt{{3\\over 2}}\\varphi)}\\right)^2$ & WMAP\\\\ \\hline $ \\Omega_{M_0}h^2$& $0.13\\pm 0.03$& $0.13\\pm 0.03\\pm 0.04$&$0.135\\pm 0.008$\\\\ \\hline $ \\Omega_{\\varphi_0}$& $0.75\\pm 0.08$& $0.76\\pm 0.07\\pm0.04$& $0.73\\pm 0.04$\\\\ \\hline $ w_{0}$& $-0.76\\pm 0.05$ &$ -0.74\\pm 0.05\\pm0.03$ &$w < -0.78$ \\\\ \\hline $ \\Gamma$& $0.19^{+0.07}_{-0.07}$ &$ 0.19^{+0.07+0.02}_{-0.05-0.02}$ &$ 0.19\\pm 0.02$ \\\\ \\hline $R$& $1.71^{+0.09}_{-0.09}$ &$1.71^{+0.09}_{-0.09}$ &$1.71\\pm 0.137$ \\\\ \\hline $l_1$& $225^{+ 11}_{- 12}$ &$ 224^{+ 11}_{- 6}$ &$220.1^{+ 0.8}_{- 0.8}$ \\\\ \\hline $l_2$& $546^{+ 27}_{- 24}$ &$ 543^{+ 24}_{- 17}$ &$546^{+ 10}_{- 10}$ \\\\ \\hline $l_3$& $830^{+ 42}_{- 29}$ &$ 825^{+ 28}_{- 34}$ &$851^{+ 31}_{- 31}$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  are fully compatible with the  recent observational data. However the models considered here, despite  being fully compatible with the present day observational data have very different behavior on long time scales. The model with the single exponential potential will eternally accelerate, and though for sufficiently large t its evolution is dominated by dark energy with $\\Omega_{\\varphi}\\rightarrow 1$ and $\\Omega_{M}\\rightarrow 0$ the scale factor increases only as a power law $a(t)\\sim t^{4/3}$ and not exponentially. In this model there is a particle horizon. The model with the double exponential potential has a different asymptotic $(t\\rightarrow \\infty)$ behavior. First of all, though $a(t)\\rightarrow \\infty$ as $t \\rightarrow \\infty$, $\\Omega_{M}$ does not decrease to zero, so for large time matter and dark energy coexist. For sufficiently large t the scale factor grows like in a FRW  matter-dominated flat model $a(t)\\sim t^{2/3}$. In this model, the particle horizon does not form and asymptotically ( for large $t$) the expansion of the universe is decelerating. It also turns out that the available observational data cannot discriminate between the quintessence models considered here and the $\\Lambda$CDM model. New data coming from high redshifts observations could remove this degeneracy."
    },
    "0410/astro-ph0410390_arXiv.txt": {
        "abstract": "{We compute polarization maps for molecular cloud cores modeled as magnetized singular isothermal toroids, under the assumption that the emitting dust grains are aspherical and aligned with the large-scale magnetic field. We show that, depending on the inclination of the toroid with the line-of-sight, the bending of the magnetic field lines resulting from the need to counteract the inward pull of gravity naturally produces a depolarization effect toward the centre of the map. We compute the decrease of polarization degree with increasing intensity for different viewing angles and frequencies, and we show that an outward increasing temperature gradient, as expected in starless cores heated by the external radiation field, enhances the decrease of polarization.  We compare our results with recent observations, and we conclude that this geometrical effect, together with other mechanisms of depolarization, may significantly contribute to the decrease of polarization degrees with intensity observed in the majority of molecular cloud cores. Finally, we consider the dependence of the polarization degree on the dust temperature gradient predicted for externally heated clouds, and we briefly comment on the limits of the Chandrasekhar-Fermi formula to estimate the magnetic field strength in molecular cloud cores. \\keywords ISM: clouds; dust, extinction; magnetic fields} ",
        "introduction": "Mapping the polarization of the thermal emission of dust at millimitre or submillimetre wavelengths (usually $\\lambda=850$~$\\mu$m or 1.3~mm) is the principal means of probing the magnetic field geometry in molecular cloud cores. A frequent characteristic of these observations is the decrease of polarization degree $p$ as a function of the total observed intensity $I$.  Usually $p$ decreases with increasing intensity $I$ with a power-law behavior, from a maximum value of $\\sim 15$\\% to about the observable limit of $\\sim 1$\\%.  This depolarization effect (sometimes referred to as ``polarization hole'', or ``polarization limb brightening'') has been observed in many dense cores and filamentary clouds.  Recent examples include the OMC-3 region of the Orion A filamentary molecular cloud where $p\\propto I^{-0.7}$ (Matthews \\& Wilson~2000), several dense cores in the dark cloud Barnard~1, where $p\\propto I^{-0.8}$ (Matthews \\& Wilson~2002), Bok globules mapped by Henning et al.~(2001) where $p\\propto I^{-0.6}$, and dense cores mapped interferometrically by Lai et al.~(2002), where $p\\propto I^{-0.8}$. In some cases, the decrease of polarization with intensity is quite steep: for example, Crutcher et al.~(2004) find $p\\propto I^{-1.2}$ in the dark cloud L183, implying that not only the polarization degree but also the polarized intensity $I_p=pI$ decreases toward the centre of this cloud. A similar situation is apparently found also in the starless core L1544 (Ward-Thompson et al.~2000).  The dark cloud Barnard 1 is an excellent example of the type of object of interest to us. This cloud, mapped with SCUBA by Matthews \\& Wilson~(2002), has dimensions of roughly $0.2 \\times 0.4$~pc, densities of a few times $10^4$~cm$^{-3}$, and temperatures of order $10$~K. These appear unfavorable characteristics for aligning grains (see e.g. Lazarian et al.~1997). Nevertheless, B1 exhibits extremely uniform polarization at a level of around 5\\% and with an extremely uniform position angle. This last point in our opinion is a strong argument in favor of an essentially uniform and likely dominant magnetic field over most of the core. However, there are exceptions to the rule in the shape of four high density (of order $10^6$~cm$^{-3}$) prestellar or protostellar ``inclusions'' of dimensions $\\sim 0.03$~pc (B1-a,b,c,d) apparently embedded within the region of the ``general core'' referred to above.  These extremely high density cores also surprisingly show evidence for polarization, albeit at a lower level and with different position angles than the general core.  In particular, in the core B1-c the polarization degree is very low and almost uniform over the whole observed intensity range (Matthews \\& Wilson~2002) \\footnote{ We note however parenthetically that the observations of polarization at high densities ($n\\sim 10^6$~cm$^{-3}$) where radiation fields can hardly be of importance represent a serious challenge to grain alignment theorists. In fact, recent work shows that radiative torques due to starlight are required to drive grains to the suprathermal rotation rates necessary to minimize the disaligning effect of random collisions (Draine \\& Weingartner~1996, 1997).}.  For illustration purposes, the $p$--$I$ relations measured in three well studied cores (L183, L1544, and B1-c) are collected in Fig.~\\ref{observations}, together with power-law fits of the original data.  \\begin{figure}[ht] \\resizebox{10cm}{!}{\\includegraphics{observations.ps}} \\caption{Distribution of percentage of polarization $p$ vs. intensity $I/I_{\\rm max}$ at $\\lambda=850$~$\\mu$m, normalized to its maximum value, for the starless cores L183 ({\\it filled triangles}, from Crutcher et al.~2004), L1544 ({\\it empty triangles}, from Ward-Thompson et al.~2000), and the B1-c core ({\\it circles}, from Matthews \\& Wilson~2002). The {\\it dashed lines} are power-law fits, with slopes $-1.6$ for L1544, $-1.2$ for L183, and $-0.4$ for B1-c.} \\label{observations} \\end{figure}  The observed decrease of polarization with increasing intensity has been attributed to changes in the dust grain optical properties or shapes in the cold cores interiors due to grain-growth (Vrba et al.~1993, or Hildebrand et al.~1999): not only are bigger grains more difficult to align than smaller grains, but also the agglomeration process may make grains more spherical, thus further reducing the alignment efficiency.  All these effects are difficult to quantify. A larger density tends in general to dealign grains (Lazarian et al.~1997), but the enhanced magnetic field strength should have the opposite effect.  On the basis of the Davis \\& Greenstein~(1951) alignment model by paramagnetic relaxation, Vrba et al.~(1981) showed that the ratio $p/A_V$ (equivalent to $p/I$ at millimetre wavelengths) should scale as $B^2 a^{-1} n^{-1}$, where $B$ is the intensity of the magnetic field, $a$ the grain size, and $n$ the ambient density. If the magnetic field strength scales as $n^{1/2}$ in molecular clouds, as empirically observed (Crutcher~1999), the effect of the increased density and magnetic field cancel, and one is left with the dependence of $p/A_V$ on the grain's size (and shape).  Fiege \\& Pudritz~(2000) showed that a helical magnetic field geometry in a cylindrical cloud with uniform grain properties could naturally produce a depolarization effect, in agreement with the submm polarization observations of the OMC-3 region of the Orion A filamentary molecular cloud (Matthews \\& Wilson~2000; Matthews, Wilson, \\& Fiege~2001). Padoan et al. (2001) were able to produce a decreasing $p$ vs. $I$ relation in three-dimensional MHD simulations of supersonic and super-Alfv\\'enic turbulence assuming that grains are not aligned above $A_V\\simeq 3$~mag, but the dynamic range of the simulation did not extend much beyond this value of extinction.  They noticed, however, that even assuming a uniform grain alignment efficiency, a decrease of $p$ with $I$ could be reproduced for particular orientations of the core magnetic field relative to the line of sight, an effect that will be further explored in the present paper.  It is possible in principle that the observed depolarization may be accounted for (at least in part) by beam smearing over tangled, small-scale field structures: Rao et al.~(1998), for example, found that the decrease in polarization toward the Kleinman-Low nebula previously seen with single-dish observations was a result of subresolution-scale variations in the magnetic field that are averaged out by larger beams. However, given the relatively high levels of polarization detected at the core peaks, this effect should not be dominant.  We will try in this paper to further analyze the influence of the magnetic field geometry on the observed depolarization effect, studying cloud models dominated by a large scale magnetic field. In particular, we show that the pinching of the magnetic field expected in dense, self-gravitating molecular cloud cores naturally produces a decrease of the polarization degree toward the centre of the core, for a large range of viewing angles. In addition, we find that this geometrical depolarization effect is further enhanced by a dust temperature gradient increasing outward, as expected in externally heated starless cores (Evans et al.~2001, Zucconi et al.~2001; Stamatellos \\& Whitworth~2003; Gon\\c {c}alves et al.~2004, hereafter GGW). However, while these effects can contribute substantially to the observed depolarization, the measured $p$--$I$ relations are sometimes steeper than our predictions. Thus, we do not suggest that the field geometry and the dust temperature distribution are the sole means by which these observations could be explained: variations in the grain optical and/or geometrical polarization properties in the densest parts of a cloud are still required to account for the full range of observed polarization values.  The structure of the paper is the following: in Sect.~2, we describe the model adopted for magnetized molecular cloud cores and define the relevant physical and geometrical parameters of the problem; in Sect.~3, we show synthetic polarization maps computed with these models, and illustrate the effect of geometrical depolarization as well as its dependence on inclination, wavelength, and temperature distribution. In Sect.~4 we summarize our conclusions. ",
        "conclusions": "In this paper we have computed the polarization of the radiation emitted by dust grains in molecular cloud cores represented as magnetically supported equilibrium configurations. To this end, we have adopted the magnetostatic models of Li \\& Shu~(1996) and the radiative transfer method developed by GGW. The Stokes parameters (and therefore the polarization degree) have been computed at two wavelengths ($\\lambda=850$~$\\mu$m and $\\lambda=450$~$\\mu$m) and for various inclinations of the cloud with respect to the line of sight, assuming that the dust grains are elongated and aligned with the large-scale magnetic field.  The main result of this paper is the demonstration that a significant depolarization effect, with characteristics very similar to those observed in actual cloud cores, can arise only because of geometrical effects, if the large-scale magnetic field has the equatorially pinched morphology predicted by magnetically dominated models.  However, we do not claim that the field geometry is the sole means by which such an effect could be produced. There may also be contributions due to grain growth in the densest parts of a cloud, and turbulence.  We also note that the assumption of axial symmetry of the density and field distribution implies that the model polarization vectors (after a $90^\\circ$ rotation), are symmetrically distributed with respect to the cloud's minor axis, at variance with observational evidence for some cores (see e.g. Ward-Thompson et al.~2000). A toroidal component of the magnetic field, ignored in the present analysis, can in principle account for the observed misalignment between the core apparent elongation and the average polarization position angle (see e.g. Vall\\'ee, Greaves \\& Fiege~2003)."
    },
    "0410/astro-ph0410359_arXiv.txt": {
        "abstract": "A TeV gamma-ray signal from the direction of the Galactic center (GC) has been detected by the H.E.S.S. experiment. Here, we investigate whether Kaluza-Klein (KK) dark matter annihilations near the GC can be the explanation. Including the contributions from internal bremsstrahlung as well as subsequent decays of quarks and $\\tau$ leptons, we find a very flat gamma-ray spectrum which drops abruptly at the dark matter particle mass. For a KK mass of about 1 TeV, this gives a good fit to the H.E.S.S.~data below 1 TeV. A similar model, with gauge coupling roughly three times as large and a particle mass of about 10 TeV, would give both the correct relic density and a photon spectrum that fits the complete range of data. ",
        "introduction": "Currently there is great interest in the very high-energy gamma-ray signal from the center of the galaxy, where recently the CANGAROO \\cite{cang}, VERITAS \\cite{verit} and H.E.S.S. \\cite{hess} collaborations have reported results. The three experiments find different spectra, something that may depend on different angular acceptances, energy calibrations, etc. The H.E.S.S. result superficially looks more detailed and accurate and is therefore the one we will focus on here, although of course the three groups will have to find a conclusive reason for the reported discrepancies. There have already been speculations of a TeV-scale dark matter (DM) particle giving, through annihilations, a steady source of these photons \\cite{sarkar,hess,horns}. Of course, there are other possibilities for emission of these gamma rays related to the physics around the Galactic massive black hole \\cite{aharonian}.  The prototype DM particle is the neutralino, the lightest supersymmetric particle. The phenomenology of this candidate has been worked out in detail over the last two decades (for a summary, see \\cite{neut}). However, these models give a rather soft gamma-ray spectrum, so one has to push masses up to the tens of TeV range to match the results of H.E.S.S., which will be our main concern in this Letter. We focus instead on an interesting alternative which has received much attention in recent years; that of universal extra dimensions (UED) \\cite{app,ser,che}. In these models, all standard model fields propagate in the higher dimensional bulk. For the effective four-dimensional theory, this means that all particles are accompanied by a tower of increasingly more massive Kaluza-Klein (KK) states. Momentum conservation in the extra dimensions, which corresponds to KK mode number conservation, is broken by an orbifold compactification which projects out unwanted degrees of freedom at the zero mode level. If the boundary terms introduced at the orbifold fixed points are identical, then a remnant of KK mode number conservation is left in the form of KK parity, and the lightest KK particle (LKP) is stable \\cite{che}. This is analogous to conserved R-parity in supersymmetric models which ensures the stability of the lightest supersymmetric particle. If the LKP is also neutral and non-baryonic, it is a potential dark matter candidate.  We consider the simplest, five-dimensional model with one compactified UED on an $S^1/Z_2$ orbifold of radius $R$. At tree level, the $n$th KK mode mass is then given by \\begin{equation} m^{(n)} = \\sqrt{(n/R)^2 + m_\\text{EW}^2}, \\end{equation} where $m_\\text{EW}$ is the zero mode mass. However, the identification of the LKP is nontrivial because radiative corrections to the mass spectrum of the first KK level are typically larger than the corresponding electroweak mass shifts given by $m_\\text{EW}$. A one-loop calculation \\cite{che} shows that the LKP is well approximated by the first KK mode of the hypercharge gauge boson $\\B$. The $\\B$ relic density was determined in \\cite{ser}. Depending on the exact form of the mass spectrum and the resulting coannihilation channels, the limit from the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{spe} of $\\Omega_\\text{CDM} h^2 = 0.12 \\pm 0.02$ corresponds to $0.5\\text{~TeV} \\lesssim m_{\\B} \\lesssim 1\\text{~TeV}$. Here $\\Omega_\\text{CDM}$ is the ratio of DM to critical density and $h$ is the Hubble constant in units of $100\\text{~km} \\text{\\,s}^{-1} \\text{\\,Mpc}^{-1}$. Collider measurements of electroweak observables give a current constraint of $R^{-1} \\gtrsim 0.3\\text{~TeV}$ \\cite{app,aga}, whereas LHC should probe compactification radii up to 1.5 TeV \\cite{cheb}.  The prospects of KK DM detection have been studied in some detail \\cite{chec,hoo,maj,serb,ber,hoo2}. Secondary gamma rays from cascading quark decays due to $\\B\\B$ annihilation in the GC produce relatively soft spectra \\cite{chec,ber}. However, unlike the supersymmetric case, charged lepton production is not helicity suppressed, which, e.g., results in a striking peak signal in the positron spectrum \\cite{chec,hoo2} for masses $\\lesssim$ 0.5 TeV. This has led us to also study the decays of $\\tau$ leptons and the spectrum of \\emph{primary} continuum gamma rays from the production and subsequent radiative photon emission of high-energy charged leptons. ",
        "conclusions": "In this Letter we have presented the somewhat surprising result that gamma rays radiated from charged leptons in $\\B\\B$ annihilations (e.g., near the GC, but also in other overdense regions of the halo, and maybe even in the diffuse extragalactic gamma-ray spectrum \\cite{BEU_01}) could be the most promising way to detect the effects of KK DM and to possibly differentiate it from other candidates. We have seen that the signature is a surprisingly hard spectrum, at high energies several orders of magnitude larger than the previously studied case of quark fragmentation alone, which extends all the way to the mass of the DM particle and then drops rapidly. Even if the present gamma-ray signal detected by H.E.S.S. will turn out to have another cause, the sharp drop at $E_\\gamma=m_{\\B}$ is a nice signature for that and other experiments to look for.  L.B. is grateful to the Swedish Science Research Council (VR) for support. We acknowledge useful discussions with J.~Edsj\\\"o and L.~Pieri.  \\vspace{-0.36cm}"
    },
    "0410/hep-th0410072_arXiv.txt": {
        "abstract": "s{The enormous red-shifting of the modes during the inflationary epoch suggests that physics at the very high energy scales may modify the primordial perturbation spectrum.  Therefore, the measurements of the anisotropies in the Cosmic Microwave Background (CMB) could provide us with clues to understanding physics beyond the Planck scale.  In this {proceeding}, we study the Planck scale effects on the primordial spectrum in the power-law inflation using a model which preserves local Lorentz invariance.  While our model reproduces the standard spectrum on small scales, it {\\it naturally}\\/ predicts a suppression of power on the large scales---a feature that seems to be necessary to explain deficit of power in the lower multipoles of the CMB.} ",
        "introduction": "In the inflationary scenario\\cite{Linde:1990-bk}, the perturbations corresponding to comoving length scales of cosmological interest today would have emerged from quantum fluctuations at the beginning of inflation with physical wavelengths smaller than the Planck length. Hence, in principle, quantum gravitational effects should have left their signatures on the primordial spectrum. This opens up the interesting possibility of probing trans-Planckian (TP) physics using the CMB~\\cite{Brandenberger-Mart:2000}.  The first year results of WMAP\\cite{Peiris-WMAP:2003} data show that the power in the quadrupole and the octopole moments of the CMB are {\\it lower than}\\/ as expected in the best fitting $\\Lambda CDM$ models. The deficit of power in the lower multipoles can not be explained within the context of the standard inflationary models (unless these models are fine-tuned\\cite{Contaldi-Pelo:2003}) and suggests a possible signature of TP physics.%  Most of the earlier efforts in incorporating TP physics into the standard field theory have involved models which {\\it break}\\/ local Lorentz invariance\\cite{Brandenberger-Mart:2000}. However, theoretically, there exists no apriori reason to believe that Lorentz invariance may be broken at the scales of inflation. More importantly, recent observations of % synchrotron emission % from the Crab nebula seem to suggest that Lorentz invariance may be preserved to very high energies\\cite{Jacobson-Libe:2002}. In such a situation, in order to study the TP effects on the primordial perturbation spectrum, it becomes important that we also consider models which preserve Lorentz invariance {\\it even as they contain a fundamental scale}.\\/ In this proceeding, we consider one such model, evaluate the resulting spectrum of density perturbations in power-law inflation and also discuss its implications for the CMB angular power spectrum. ",
        "conclusions": "In this proceeding, we have studied the TP effects on the spectrum of primordial perturbations in power-law inflation using an approach that preserves local Lorentz invariance. We assumed that the TP effects modify the standard propagator in a particular manner. We find that the resulting modified spectrum remains scale invariant at the ultra-violet end, but, interestingly, it exhibits a suppression of power at the infra-red end---a feature that seems to be necessary to account for the deficit of power in the lower multipoles of the CMB\\cite{Peiris-WMAP:2003,Contaldi-Pelo:2003}. However, the amount of suppression predicted by our model in power-law inflation turns out to be {\\it far less}\\/ than that seems to be required to fit the WMAP data. Nevertheless, the loss of power at large scales suggests that the power spectrum we have obtained may fit the WMAP data better than the standard $\\Lambda$CDM model. It will be interesting to analyze the implications of our model for WMAP data in the context of slow-roll inflation."
    },
    "0410/astro-ph0410384_arXiv.txt": {
        "abstract": "Recent observational and theoretical studies have raised the possibility that the collimated outflows in gamma-ray burst (GRB) sources have two distinct components: a narrow (opening half-angle $\\theta_{\\rmj,\\rmn}$), highly relativistic (initial Lorentz factor $\\eta_\\rmn \\gtrsim 10^2$) outflow, from which the $\\gamma$-ray emission originates, and a wider ($\\theta_{\\rmj,\\rmw} \\lesssim 3\\, \\theta_{\\rmj,\\rmn}$), moderately relativistic ($\\eta_\\rmw\\sim 10$) surrounding flow. Using a simple synchrotron emission model, we calculate the R-band afterglow lightcurves expected in this scenario and derive algebraic expressions for the flux ratios of the emission from the two jet components at the main transition times in the lightcurve.  For viewing angles $\\theta_\\rmobs < \\theta_{\\rmj,\\rmn}$ we find that the contribution of the wide component to the optical afterglow is negligible if its kinetic energy $E_\\rmw$ is significantly smaller than that of the narrow component, $E_\\rmn$, as expected for the jet core and cocoon outflow components in the collapsar jet-breakout model. However, if $E_\\rmw/E_\\rmn>1$ [but the isotropic-equivalent energy ratio $E_{\\rmiso,\\rmw}/E_{\\rmiso,\\rmn}=(E_\\rmw/E_\\rmn)(\\theta_{\\rmj,\\rmn}/ \\theta_{\\rmj,\\rmw}) ^2$ remains $<1$], as expected for the decoupled neutron and proton components, respectively, in an initially neutron-rich, hydromagnetically accelerated jet model, then the narrow component only dominates the early afterglow and the wide component takes over after its nominal deceleration time $t_{\\rmdec,\\rmw}$ (typically $\\sim 0.1-1\\ {\\rm days}$). Given that $t_{\\rmdec,\\rmw}$ is comparable to the jet-break time $t_{\\rmjet,\\rmn}$ of the narrow component for characteristic parameter values, the emergence of the wide component at $t_{\\rmdec,\\rmw}$ may mask the jet break in the narrow component at $t_{\\rmjet,\\rmn}$, which in turn may lead to an overestimate of the $\\gamma$-ray energy emitted by the source and hence of the required $\\gamma$-ray emission efficiency. We apply this scheme also to X-ray flash sources, which we interpret as GRB jets viewed at an angle $\\theta_\\rmobs > \\theta_{\\rmj,\\rmn}$. Finally, we argue that a neutron-rich hydromagnetic outflow may naturally give rise to repeated brightening episodes in the afterglow lightcurve as observed in GRB 021004 and GRB 030329. ",
        "introduction": "\\label{introduction}  Gamma-ray bursts (GRBs) and their afterglows are commonly interpreted in terms of a relativistic outflow that emanates from the vicinity of a solar-mass neutron star or black hole \\citep[e.g.,][]{P99,M02}. In this picture, the prompt gamma-ray emission is attributed to a highly relativistic ejecta (with an initial Lorentz factor $\\gamma \\gtrsim 10^2$), whereas the subsequent afterglow emission in the X-ray, optical, and radio (over hours, days, and weeks, respectively, after the GRB) arises from the shock that is driven into the ambient medium as the ejecta sweeps up the external medium and decelerates. Most afterglow observations to date have been carried out hours to days after the GRB event, by which time the Lorentz factor of the afterglow shock has decreased to $\\lesssim 10$. These observations have revealed the presence of achromatic breaks in the afterglow lightcurves of many sources, which strongly indicate that GRB outflows are collimated into narrow jets \\citep{R99,SPH99}.  Recently, however, the possibility that at least some GRB outflows consist of two distinct components has been raised in the literature. On the observational side, this possibility was first invoked by \\citet{P98}, who suggested that the afterglow from GRB 970508 could be explained in terms of a narrow jet surrounded by an isotropic outflow. \\citet{F00} subsequently proposed that the $\\gamma$-rays and early (shorter-wavelength) afterglow emission in GRB 991216 could be attributed to a narrow, ultrarelativistic outflow component, and that the late (longer-wavelength) afterglow emission in this source originates in a separate wide component that is only mildly relativistic. A similar picture was proposed for GRB 030329 by \\citet{B03} and \\citet[][]{S03}. A two-component model was also suggested as an explanation of the observed rebrightening of the X-ray flash (XRF) source XRF 030723 \\citep[][]{H04} as well as of the apparent peak-energy distribution of GRBs and XRFs \\citep[][]{LD04} and of the origin of the blueshifted optical absorption features in the spectrum of the GRB 021004 afterglow \\citep[][]{S05}.  The possibility of a two-component outflow in GRB sources has been independently indicated by theoretical considerations. One can broadly divide the physical models that give rise to such an outflow into two classes: models in which the separation into two components is an intrinsic property of the outflow, and those in which a narrow relativistic jet gives rise to a wider and slower component as it propagates through (and interacts with) the envelope of a progenitor massive star. One example of a model of the first type was worked out by \\citet{LE93} and \\citet{VPL03}: it consists of (1) a relativistic, baryon-poor jet driven electromagnetically along disk-anchored magnetic field lines that thread the horizon of a rotating black hole, and (2) a subrelativistic, baryon-rich wind that is driven thermally from that disk. Another example is provided by hydromagnetically driven jets that originate from a neutron star or a neutron-rich accretion disk that form in the collapse of a massive star \\citep{VPK03}. In this case the neutrons decouple at a moderate Lorentz factor while the protons continue to be accelerated and collimated by the electromagnetic forces, giving rise to a narrow, highly relativistic proton component and a wider and slower neutron component (which, after decoupling, is transformed into a moderately relativistic proton component through neutron decay). Examples of models of the second type include jet-induced core-collapse supernovae, wherein collimated high-velocity jets cause the envelope of the progenitor massive star through which they propagate to be ejected with subrelativistic speeds and an oblate geometry \\citep{K99}, and the collapsar model, in which the outflow resulting from the jet breakout through the progenitor star's envelope is predicted to consist of a highly relativistic jet core and a moderately relativistic surrounding cocoon (\\citealt{ZWH04}; see also \\citealt{RCR02}).  In this paper we focus on two-component GRB outflows in which both components initially move with relativistic speeds and therefore end up contributing to the optical afterglow emission. Accordingly, we adopt as representative examples the hydromagnetically accelerated, initially neutron-rich jet model of \\citet{VPK03} and the collapsar jet-breakout model of \\citet{ZWH04}. According to the numerical simulations of the latter authors, the narrow component in the collapsar model has a Lorentz factor $\\eta_{\\rmn} \\gtrsim 100$ and an opening half-angle $\\theta_{\\rmj,\\rmn} \\sim 3^\\circ-5^\\circ$, whereas the corresponding quantities for the wide component are $\\eta_{\\rmw} \\sim 15$ and $\\theta_{\\rmj,\\rmw} \\sim 10^\\circ$, respectively. The characteristic Lorentz factors in this scenario are very similar to those ($\\eta_{\\rmn} \\sim 200$ and $\\eta_{\\rmw} \\sim 15$) in the representative neutron-rich hydromagnetic model of \\citet{VPK03}.\\footnote{In the simplified model used by \\citet{VPK03}, the value of $\\theta_{\\rmj,\\rmw}$ could not be calculated exactly; in this paper we assume that it can be as large as $\\sim 3~\\theta_{\\rmj,\\rmn}$.} However, {\\em in contrast} with the collapsar model, in which the highly relativistic jet component is in general more energetic (typically by about an order of magnitude) than the cocoon material, the asymptotic kinetic energy $E_{\\rmw}$ of the wide component in the hydromagnetic jet model typically exceeds the corresponding energy of the narrow component ($E_{\\rmw} \\approx 2\\, E_{\\rmn}$ in the fiducial model of \\citealt{VPK03}).  Our goal in the present work is to examine some of the observational properties of two-component GRB outflows. In particular, we calculate (\\S~\\ref{AG_LC}) the approximate optical afterglow lightcurves that are produced by the shocks that the two jet components would drive into the ambient medium. We then argue (\\S~\\ref{apply}) that outflows of this type may have significant general implications to our understanding of GRB and XRF sources. Our conclusions are given in \\S~\\ref{conclude}. ",
        "conclusions": "\\label{conclude}  In this paper we study the optical afterglow lightcurves produced by GRB sources that have two distinct outflow components. The possibility that GRB jets have a narrow/fast core and a wider/slower outer component has been indicated by observations of both the $\\gamma$-ray and the afterglow emission (including the afterglows of GRB 970508, GRB 991216, and GRB 030329) and independently by theoretical considerations. Since we are interested in the distinct afterglow signatures of the two components, we focus on models in which the Lorentz factors of both outflow components are initially $\\gg 1$. We choose the hydromagnetically driven, initially neutron-rich jet model of \\citet{VPK03} and the collapsar jet-breakout model of \\citet{ZWH04} as being representative of scenarios in which the dual nature of the outflow reflects initial conditions and propagation effects, respectively. In both of these models the characteristic initial Lorentz factor and opening half-angles are $\\eta_\\rmn \\gtrsim 10^2$ and $\\theta_{\\rmj,\\rmn} \\sim 0.05$ for the narrow component and $\\eta_\\rmw \\sim 10$, $\\theta_{\\rmj,\\rmw} \\lesssim 3\\, \\theta_{\\rmj,\\rmn}$ for the wide one, and the $\\gamma$-ray emission originates in the narrow component. They are distinguished, however, by the ratio of the kinetic energy injected into the two components: $E_\\rmw/E_\\rmn \\sim 0.1$ for the collapsar model and $\\gtrsim 2$ for the neutron-rich hydromagnetic model (with $E_\\rmw + E_\\rmn$ inferred to be $\\sim 10^{51}\\ {\\rm ergs}$).  Using a simple synchrotron emission model, we calculate the afterglow emission produced by the shocks that the two components drive into the ambient medium. We derive useful algebraic expressions for the component flux ratios at the main transition times in the light curve (in particular, the wide component's deceleration time $t_{\\rmdec,\\rmw}$ and the jet-break times $t_{\\rmjet,\\rmw}$ and $t_{\\rmjet,\\rmn}$) for the cases where the observation angle $\\theta_\\rmobs$ satisfies $\\theta_\\rmobs < 1.5\\,\\theta_{\\rmj,\\rmn}$ and $1.5\\,\\theta_{\\rmj,\\rmn}<\\theta_\\rmobs<1.5\\,\\theta_{\\rmj,\\rmw}$, respectively (where in the latter case $t_{\\theta,\\rmn}$ is the relevant break time for the narrow component). We study the behavior of the optical lightcurves for different values of $E_\\rmw/E_\\rmn$ and find that, for the adopted characteristic parameters, the contribution of the narrow component dominates at all times if this ratio is $\\ll 1$ (as in the collapsar jet-breakout model), but that the contribution of the wide component becomes dominant for $t \\gtrsim t_{\\rmdec,\\rmw}$ if $E_\\rmw \\gtrsim 2\\, E_\\rmn$ (as in the neutron-rich hydromagnetic model). The emergence of the wide component may be related to the pronounced brightening detected in the lightcurves of several afterglows $\\sim 0.1-1\\ {\\rm days}$ after the GRB (see Fig.~\\ref{fig3}).  For typical parameter values $t_{\\rmdec,\\rmw}$ is found to be comparable to $t_{\\rmjet,\\rmn}$. It follows that, if $E_\\rmw > E_\\rmn$, then the steepening of the narrow component's lightcurve at $t \\gtrsim t_{\\rmjet,\\rmn}$ could be masked by the emergence (and subsequent dominance) of the wide component (see Fig.~\\ref{fig1}). Under these circumstances, the only clearly discernible jet break in the optical lightcurve would occur at $t_{\\rmjet,\\rmw}$. We suggest that this may have led to an overestimate of the emitted $\\gamma$-ray energy in many GRBs because the wide component's opening half-angle $\\theta_{\\rmj,\\rmw}$ --- rather than the narrow component's angle $\\theta_{\\rmj,\\rmn}$ --- was used in converting the measured $E_{\\gamma,\\rmiso}$ into the true energy $E_\\gamma$. This, in turn, would have led to an overestimate [by a factor $\\sim (E_\\rmn/E_\\rmw)(\\theta_{\\rmj,\\rmw}/\\theta_{\\rmj,\\rmn})^2 =E_{\\rmiso,\\rmn}/E_{\\rmiso,\\rmw}$] of the ratio $E_\\gamma/E_\\rmn$ that determines the kinetic-to-radiative energy conversion efficiency of the outflow. Factoring in this overestimate [which can be done when the component kinetic energies satisfy $1< E_\\rmw/E_\\rmn< (\\theta_{\\rmj,\\rmw}/\\theta_{\\rmj,\\rmn})^2$] would alleviate the need to account for conversion efficiencies $O(1)$ in internal shock models of GRBs. Dense monitoring of the afterglow lightcurve during the time interval ($\\sim 0.1-1\\ {\\rm days}$) that encompasses $t_{\\rmjet,\\rmn}$ and $t_{\\rmdec,\\rmw}$ could provide a test of this suggestion. If $t_{\\rmjet,\\rmn} \\lesssim t_{\\rmdec,\\rmw}$ then the lightcurve should exhibit a convex bump during this time interval, whereas if this inequality is reversed a concave flattening would be expected (see Fig.~\\ref{fig2}). The above considerations also apply to jets observed at $\\theta_\\rmobs\\gtrsim 1.5\\, \\theta_{\\rmj,\\rmn}$, which would be perceived as X-ray flash sources. For $E_\\rmw > E_\\rmn$, the afterglow emission from such sources would be dominated by the wide outflow component, although the narrow component might give rise to a bump in the lightcurve at $t \\approx t_{\\theta,\\rmn}$.  The well-monitored afterglow lightcurves of GRB 021004 and GRB 030329 exhibited a significant early brightening that was followed by several less pronounced rebrightenings on a time scale of days. These episodes can be satisfactorily interpreted as refreshed shocks. We point out that the initially neutron-rich hydromagnetic jet model, in which the decoupled protons and neutrons give rise to a narrow/fast and wide/slow outflow components, respectively, could naturally account for the appearance of refreshed shocks following (or even coincident with) the wide component's emergence in the afterglow lightcurve. Future high-quality observations should be able to determine whether the nonsteady behavior found in these two objects is a common trait of GRB sources."
    },
    "0410/astro-ph0410667_arXiv.txt": {
        "abstract": "We present the optical identifications of a 95 $\\mu$m ISOPHOT sample in the Lockman Hole over an area of about half square degree. The catalogue (Rodighiero et al. 2003) includes 36 sources, making up a complete flux-limited sample for $S_{95 \\mu m} \\geq 100$ mJy. Reliable sources were detected, with decreasing but well-controlled completeness, down to $S_{95 \\mu m} \\simeq 20$ mJy. We have combined mid-IR and radio catalogues in this area to identify the potential optical counterparts of the far-IR sources. We found 14 radio and 13 15 $\\mu$m associations, 10 of which have both associations. For the 11 sources with spectroscopic redshift, we have performed a spectrophotometric analysis of the observed Spectral Energy Distributions. Four of these 95 $\\mu$m sources have been classified as faint IR galaxies ($L_{FIR}<1.e11 L_{\\odot}$) , six as LIRGs and only one ULIRG. We have discussed the redshift distribution of these objects, comparing our results with evolutionary model predictions 95 and 175 $\\mu$m. Given their moderate distances (the bulk of the closest spectroscopically identified objects lying at $z<0.2$), their luminosities and star formation rates (median value $\\sim$10M$_{\\odot}$/yr), the sources unveiled by ISOPHOT at 95 $\\mu$m seem to correspond to the low redshift ($z<0.3$) FIRBACK 175 $\\mu$m population, composed of dusty, star-forming galaxies with moderate star formation rates. We computed and compared different SFR estimators, and found that the SF derived from the bolometric IR luminosity is well correlated with that computed from the radio and mid-IR fluxes. ",
        "introduction": "The cosmic infrared/sub-millimeter background (CIRB), detected by COBE (Puget et al. 1996; Fixsen et al. 1998) and peaking around 140 $\\mu$m with an energy comparable to that of the optical/UV background, has been interpreted as the integrated emission by dust present in distant and primeval galaxies. The emission of the CIRB represents more than half of the overall cosmic background energy density (Gispert et al. 2000) while only approximatively one-third of the bolometric luminosity of local galaxies ($z < 0.1$) is processed by dust into the infrared (Soifer \\& Neugebauer 1991).  This implies that the universe at $z>0.1$ is even more active in the infrared (IR) than the local one investigated by the Infrared Astronomical Satellite (IRAS, Ashby et al. 1996). The Infrared Space Observatory (ISO, Kessler et al. 1996), with its improved sensitivity and angular resolution compared to IRAS, made possible deeper IR surveys, thus allowing to detect faint more distant galaxies, both in the mid and in the far infrared. In the last few years, different studies tried to address the question about the nature of those sources contributing to the CIRB. Elbaz et al. (2002) found that the galaxies unveiled by ISOCAM surveys at 15 $\\mu$m are responsible for the bulk of the CIRB ($\\sim$60\\%). At longer wavelengths, the sources resolved by ISOPHOT account for less than 10\\% of the CIRB at 175 $\\mu$m (Dole et al., 2001). A similar result was recently confirmed by us (Rodighiero et al. 2003): we evaluated that from 10\\% to 20\\% of the CIRB has been resolved into sources at 95 $\\mu$m.  The models (Lagache et al. 2003, Franceschini et al. 2001) predict a bimodal redshift distribution for the 175 $\\mu$m FIRBACK (Dole et al., 2001) population: Lagache et al. (2003) forsee that $\\sim$60\\% of the 175 $\\mu$m sources with fluxes $S_{175}>180$ mJy should lie at redshift below 0.25, the rest being mostly at redshift between 0.8 and 1.2. This trend is confirmed by Kakazu et al. (2002), who published the first results from optical spectroscopy of 175 $\\mu$m Lockman Hole sources. From sub-millimeter and near-IR photometric observations of the FIRBACK sources, Sajina et al. (2002) found that normal star-forming galaxies lie at $z\\simeq$0, while a much more luminous population at $z\\sim$0.4-0.9 are ultraluminous IR galaxies (ULIRGs). These results have been more recently strengthened by Patris et al. (2003), who performed optical spectroscopy of the brighest 175 $\\mu$m sources. They found that the fraction of AGNs is low (about 15\\%) and that most sources are nearby ($z<0.3$), dusty, star forming galaxies, with moderate star formation rates (a few 10 $M_{\\odot}/yr$), essentially undistinguishable from the faint IRAS sources. A preliminary spectroscopic analysis of the optically faint FIRBACK sources (Chapman et al. 2002) indicates that the far-infrared background must include a substantial population of cold and luminous galaxies. These galaxies should represent an intermediate population between the local luminous IR galaxies and the high-redshift sub-millimeter sources.  The above mentioned analyses, aimed to discerne the nature of the far-IR ISO population, are essentially based on ISOPHOT 175 $\\mu$m selected samples. For the first time, in the present paper we try to address this question by directly looking at those sources selected in the ISOPHOT 95 $\\mu$m channel. This filter samples the dust emission peak in the Spectral Energy Distribution (SED) of star-forming galaxies around 60 to 100 $\\mu$m. For luminous infrared galaxies, emitting more than 80\\% of the bolometric luminosity in the far-IR, this far-IR peak is the best measure of the bolometric luminosity, and the best estimator of the star formation rate.  The paper is organized as follows. In Section 2 we present the 95 $\\mu$m sample and the available observations, covering the SED from the optical to the radio. In Section 3 we discuss the optical identifications of the ISOPHOT sources. Section 4 is devoted to the spectrophotometric analysis of the SEDs, and to the discussion of the physical properties of the sample galaxies. In Section 5 different SFR estimators are compared. We then report in Section 6 the summary of our results.  We assume throughout this paper $\\Omega_M$=0.3,  $\\Omega_{\\Lambda}$=0.7 and $H_0$=65 $km~ s^{-1} Mpc^{-1}$. ",
        "conclusions": "We have studied the optical identifications of a 95 $\\mu$m ISOPHOT sample within the Lockman Hole. Exploiting the quality of the data reduction (Rodighiero et al. 2003) and the depth of the survey (that is almost complete around 100mJy), we have combined mid-IR and radio catalogues in this area to identify the potential counterparts of the far-IR sources. We found 14 radio and 13 15 $\\mu$m associations, 10 of which have both associations. For these sources we have been able to detect the optical and the near-IR (where available) counterparts. 19 sources do not present evident identifications at other wavelengths. Deeper radio observations are needed to improve the statistics.  We have performed a preliminary spectrophotometric analysis of the observed SEDs for the 11 sources with an associated spectroscopic redshift. This classification, based on a multiwavelength fitting procedure, led to the identification of four IR faint galaxies, six LIRGs and only one ULIRG as counterpart of the ISO 95 $\\mu$m sources.  We have discussed the redshift distribution of these objects, comparing our results with evolutionary model predictions at 95 and 175 $\\mu$m. The sources unveiled by ISOPHOT at 95 $\\mu$m might correspond to the FIRBACK 175 $\\mu$m population at low redshift ($z<0.3$), that is composed of dusty, star-forming galaxies with moderate star formation rates. From the spectral fit we derived SFRs in the range $\\sim$1-400M$_{\\odot}$/yr with a median value of $\\sim$10M$_{\\odot}$/yr. Only one galaxy is found at $z>0.5$.  Finally, we compared the mid- and far-IR fluxes and argued that the 95 $\\mu$m should be a better tracer of the stellar formation, given that the mid-IR emission seems to underestimate the bolometric luminosity of IR selected sources (by a median percentage of $\\sim 30\\%$). We then argue that spectral templates fitting and empirical equations are not globally consistent yet.  However, the SF derived from the IR luminosity is not inconsistent with that computed from the radio, apart the observed scatter in the distribution. However, wider samples are required to constrain the results presented in this paper. Moreover, there is still many controversy about how to translate monochromatic fluxes to total FIR luminosities, and these to SFRs.  The Spitzer Space Observatory (Fazio et al., 1999) has recently observed the same region of the Lockman Hole in complementary wavebands at 24, 70 and 160 $\\mu$m with MIPS and in the near-IR with IRAC. This will provide additional information on the SEDs of ISO sources and will improve their identifications."
    },
    "0410/astro-ph0410451_arXiv.txt": {
        "abstract": "\\begin{center} {\\large\\bf Abstract} \\end{center} The relation between the results of cosmological N-body simulations, and the continuum theoretical models they simulate, is currently not understood in a way which allows a quantification of N dependent effects. In this first of a series of papers on this issue, we consider the quantification of such effects in the initial conditions of such simulations. A general formalism developed in \\cite{andrea} allows us to write down an exact expression for the power spectrum of the point distributions generated by the standard algorithm for generating such initial conditions. Expanded perturbatively in the amplitude of the input (i.e. theoretical, continuum) power spectrum, we  obtain at linear order the input power spectrum, plus two terms which arise from discreteness and contribute at large wavenumbers. For cosmological type power spectra, one obtains as expected, the input spectrum for wavenumbers $k$ smaller than that characteristic of the discreteness. The comparison of real space correlation properties is more subtle because the discreteness corrections are not as strongly localised in real space. For cosmological type  spectra the theoretical mass variance in spheres and two point correlation function are well approximated  {\\it above} a finite distance. For typical initial amplitudes this distance is a few times the inter-particle distance, but it diverges as this amplitude (or, equivalently, the initial red-shift of the cosmological simulation) goes to zero, at fixed particle density. We discuss briefly the physical significance of these discreteness terms in the initial conditions, in particular with respect to the definition of the continuum limit of N-body simulations. ",
        "introduction": "The goal of dissipationless cosmological N-body simulations (NBS) is to trace the evolution of the clustering of matter under its self-gravity, starting from a cosmological time at which perturbations to homogeneity are small at the physically relevant scales (for reviews, see e.g. \\cite{efstathiou_init, couchman, bertschinger_98}). A fundamental question about such simulations concerns the degree to which they reproduce the evolution of the simulated models. The problem arises because the N-body technique, which simulates a large number of particles evolving under their self-gravity (with some small scale regularization), is not a direct discretization of the theoretical models.  The N-body approach is taken because it is not numerically feasible to simulate, at a useful level of resolution, the continuum Vlasov-Poisson equations describing the evolution theoretically.  Since the Vlasov-Poisson system corresponds to an appropriately defined N $\\rightarrow \\infty$ of this particle dynamics \\cite{braun+hepp, spohn}, NBS can be considered as related to the models in this limit. The problem of discreteness is thus that of the relation of the results obtained from these simulations, for typical statistical quantities characterising clustering, with those which would be obtained with such a simulation done with N $\\rightarrow \\infty$ particles. The existing studies of this ``convergence'' problem in the literature (e.g. \\cite{efstathiou_init, splinter, melott_all, discreteness-hamana, diemandetal_2body, diemandetal_convergence}) are almost exclusively numerical, and consider the stability of different measured quantities as a function of N. In the absence, however, of any analytic understanding of the possible N dependence of the results, such studies, which extend over a very modest range of N, cannot be conclusive. Different groups of authors have in fact drawn very different conclusions about the correctness of results for standard quantities at smaller scales. Some \\cite{splinter, melott_alone} even place in question the validity of results for clustering amplitudes below the initial interparticle distance, while such results are widely interpreted as physical in almost all current simulations. Further it is not specified in such studies how precisely the limit of large N should be taken, i.e., which other parameters (e.g. box-size, force softening, initial red-shift)  should be kept fixed or varied. These questions are becoming of ever greater practical importance as the quantification of the precision of results from simulations is essential in order to confront cosmological models with a rich host of observations (see, e.g., \\cite{Huterer:2004tr}).  In this paper we address only one simple aspect of this problem: the relation between the discretized initial conditions (IC) of an NBS and the IC of the corresponding theoretical model. More specifically we study and quantify analytically the differences between the two for the two point correlation properties, in real space and reciprocal space, in the infinite volume limit (at a fixed particle density). The discreteness effects, i.e., the differences between the continuous theoretical IC and the discrete IC of the actual NBS are then terms which depend on the particle density. We study these terms and their relative importance for different theoretical IC. We also study and check our analytic results with the aid of numerical simulations, performed in one dimension because of the modest numerical cost of calculating the ensemble average using a large number of realisations. We underline that our analytic results apply to the infinite volume limit, i.e., to the case that the size of the simulation box goes to infinity at fixed particle density. They thus do not include effects associated with the finite size of the box. Such effects, which are quite distinct from those studied here, have been studied extensively elsewhere (see e.g. \\cite{gelb-bertschinger, colombietal1, colombietal2, colombietal3, pen, sirko}), both analytically and numerically.  The direct motivation for this work on IC in NBS came originally from some {\\it numerical} studies of these issues by two groups \\cite{Baertschiger:2001eu, alvaro_knebe, Baertschiger:2003jt, alvaro_knebe2}, who have drawn  quite different conclusions about the accuracy with which the IC produced by the canonically used algorithm in NBS represent the input IC \\footnote{A more detailed account is given in the conclusions section below.}. The results in this paper, which are essentially analytic, clarify the issues underlying the discussion in and differences between these numerical studies. Our conclusions are consistent with the findings of both sets of authors, and explain the differences. In short the authors of \\cite{alvaro_knebe, alvaro_knebe2} are correct that certain real space properties, notably the  mass variance in spheres, are in fact reasonably well represented for typical IC in NBS. The authors of \\cite{Baertschiger:2001eu}, however, are correct in diagnosing the important systematic differences between the actual and theoretical correlation properties in real space.  Indeed one of our main findings is that there is a very non-trivial difference between the two spaces: while the discreteness of the underlying particle distribution is strictly localized in reciprocal space, this is not the case in real space. The result is that, in the  limit of low amplitude initial density fluctuations --- or, equivalently, high initial redshift for the simulation --- the correlation properties of the input theoretical model are approximated well only in reciprocal space. Taking instead the limit that the particle density goes to infinity at fixed amplitude, the theoretical correlation properties are recovered in both real and reciprocal space, provided an appropriate cut-off is imposed at large wavenumber in the input PS.  The implications of these results for what concerns the agreement between an {\\it evolved} NBS and the evolved theoretical model of which it is the discretization is beyond the scope of this paper. In a forthcoming paper \\cite{discreteness2_mjbm}, treating discreteness effects up to shell crossing, we will see that the evolution of an NBS deviates arbitrarily from its continuum counterpart as the initial red-shift increases at fixed particle density, while keeping the amplitude fixed at increasing particle one approximates increasingly well the continuum (fluid limit) evolution. Thus the results we find for the initial conditions here do indeed turn out to have physical significance for the question of discreteness effects in the evolved simulations.  The algorithm used to generate IC in NBS which we analyze is in fact well defined, in the infinite volume limit, only for a certain range of asymptotic behaviors of the input theoretical PS. We specify here carefully these limits. While it turns out that they are not particularly relevant to current cosmological models, they are of interest in other contexts in which this algorithm may be used, notably in the study of gravitational clustering from other classes of initial conditions (e.g. \\cite{melott+shandarin_k4, bagla+padmanabhan_97}). These properties are also of interest in the context of statistical physics, where the problem of ``realizability'' of point processes is studied (see e.g. \\cite{costin+lebowitz, crawford+torquato+stillinger, ucheetal_2006}). Specifically we find that the algorithm has interesting limits for the case of very ``blue'' input PS: for the case of spectra with a small $k$ behavior proportional to $k^{n}$, and $n>1$, the real space variance is never that of the input model, while for $n>4$ the reciprocal space representation is never faithful either.  The paper is organized as follows. In the section which follows we briefly review the standard method for setting up IC for cosmological simulations, using the Zeldovich approximation. This also sets conventions for notation in the rest of the paper. In Sect. \\ref{IC_exact_kspace}  we analyze the PS of the configurations of points generated in this way, comparing it with the PS of the input theoretical model. To obtain these results we use a very general exact result derived in \\cite{andrea}, which gives the two point properties of a point distribution generated by superimposing an arbitrary correlated displacement field on an arbitrary initial stochastic point distribution. In the following section we consider how these properties described in $k$-space translate into real space. Specifically we present a general qualitative analysis of the relation between the two point correlation properties of the IC and those of the input models. We treat specifically  the mass variance in spheres, and the reduced two point correlation function. For the latter case the comparison of the theoretical model and full IC is more difficult, because of the complicated non-monotonic form of the correlation function of the underlying point distributions. In the following section we illustrate, and verify our results quantitatively, using one dimensional numerical simulations. We choose to work in one dimension for numerical economy, and because all the pertinent questions can be posed equally well and answered in this case\\footnote{The same is evidently not true when considering discreteness effect in the dynamics.}. In the final section we summarize our results, discuss what conclusions can be drawn concerning the papers mentioned above which motivated the present study, and finally briefly comment on the physical significance of our results, which will be further developed in the companion paper \\cite{discreteness2_mjbm}. Several technical details in the paper, notably concerning the perturbative expansion of the exact expression for the PS of the generated IC, are given in three appendices. ",
        "conclusions": "We first summarize our findings on the accuracy and limitations of the standard algorithm for generating IC for cosmological simulations. We then discuss the conclusions we can draw, in the light of our analysis, about the some numerical work on IC \\cite{Baertschiger:2001eu, alvaro_knebe2} which partly motivated our study. Finally we turn to the relevance of our results to the problem of understanding discreteness effects in the evolution of cosmological simulations.  \\subsection{Results on generation algorithm}  We have investigated systematically the algorithm used to generate IC of N-body simulations in cosmology, for {\\it any} given input PS. More specifically we have focussed on the comparison of the two point correlation properties, in real and reciprocal space, of the IC with those of the input theoretical models. We consider input PS which are a simple power-law $P(k) \\propto k^n$, but the corresponding results for more complicated cases may be easily inferred. Our main results are:  \\begin{enumerate} \\item Applied on a grid with appropriate sharp cut-off at the Nyquist frequency $k_N$, the point distribution produced by the algorithm has PS exactly equal to the input one, below $k_N$, to linear order in the amplitude and for $-d<n<4$. For $k>k_N$ we have also given the exact expression for the PS, which is thus the leading discreteness correction in this space. It is a term of high amplitude, with a damped oscillating form with maxima at the Bragg peaks of the underlying lattice.  \\item Applied to a `glass' pre-initial configuration, the result is almost the same, except that the discreteness correction has a small $k$ tail proportional to $k^2$. Thus the range of ``faithful representation'' of the PS is $-d<n<2$. This latter restriction is not of relevance to current cosmological models, for which the effective exponent at all $k$ is within this range.  \\item  The algorithm does not produce IC representing faithfully an input PS with $n>4$ for arbitrarily small $k$. There is the case because there is a term proportional to $k^4$ in the PS of the generated PS, at second order in the amplitude of the input PS.  \\item For the case $n<-d$ the algorithm is not well defined in the infinite volume limit, and we have verified that results in a finite volume depend strongly on the volume.  \\item The transposition of these results to real space is more subtle than one might have anticipated, due to the fact that the mass variance and two point correlation of the underlying `pre-initial' point distribution are delocalised in this space.  \\item For models with $-d<n<1$ the real space variance in spheres can be well represented by the generated configurations starting from a finite scale $R_{min}$ proportional to the inter-particle spacing. For typical chosen input amplitudes it is a few times this distance, but we note that it diverges as the amplitude of the input model goes to zero.  \\item For models with $n>1$, the real space variance is always dominated, at linear order in the amplitude, by the ``pre-initial'' variance of the lattice or glass.  \\item The conclusions concerning the representation of the reduced two point correlation function are quite similar to those for the mass variance: the theoretical properties are recovered above a finite scale proportional to the inter-particle distance, which diverges as the amplitude goes to zero.  In practice there is a further difference with respect to the mass variance, in that the value of this scale depends also on the smoothing is necessarily introduced in estimation of the correlation function. For a sufficiently narrow smoothing the correlation function will always show at a given scale, for sufficiently low amplitude of the input model, the underlying structure of the lattice or glass configuration.  \\end{enumerate}  \\subsection{Comments on precedent literature }  Let us now consider, in the light of these results, the articles \\cite{Baertschiger:2001eu, alvaro_knebe, Baertschiger:2003jt, alvaro_knebe2} which have partly motivated this work. These two collaborations draw, on the basis of numerical studies, very different conclusions about the measured mass variance in spheres and two-point correlation function of the IC of cosmological NBS.  In cosmology the IC of NBS are invariably studied only in reciprocal space, simply because it is the natural one for the description of cosmological models at early times.  In the first of these papers \\cite{Baertschiger:2001eu} the authors examined instead IC in real space, through a numerical study of the IC of some large cosmological simulations performed by the Virgo consortium \\cite{virgo}. Their finding was, very surprisingly, that the measured and theoretical values of both the mass variance in spheres and the two point correlation function did not match. In \\cite{alvaro_knebe} the same analysis was repeated by a different set of authors, and an error in the normalization in \\cite{Baertschiger:2001eu} of the theoretical variance was identified. Correcting for this error the authors concluded that the agreement between the measured and theoretical properties was good for the variance, while the authors of \\cite{Baertschiger:2001eu}, in a reply \\cite{Baertschiger:2003jt}, argued that the agreement was still very poor. For the two point correlation function the results of both sets of authors agreed, showing an estimated correlation function qualitatively and  quantitatively different to the expected one. The two sets of authors gave a quite different interpretation to this discrepancy: in \\cite{Baertschiger:2001eu} it was attributed to a probable systematic difference between the two quantities due to the underlying correlation in the ``pre-initial'' configuration, while \\cite{alvaro_knebe} argued that it was more likely simply due to statistical noise of the estimator. In a further article \\cite{alvaro_knebe2} the second authors analyzed these same quantities in the IC of another set  of cosmological simulations, and arrive at the same conclusions as in \\cite{alvaro_knebe} concerning both quantities.  For what concerns the mass variance we have seen that the degree of agreement between the theoretical and measured variance depends on the normalization of the model, i.e., on the initial red-shift of the simulation. Neither collaboration has studied the dependence of their conclusions on this crucial parameter, nor identified it as relevant. Thus the conclusions of \\cite{alvaro_knebe2, alvaro_knebe} about the reliability in general of the representation of the input mass variance by the IC are, strictly, incorrect. However, their conclusion that the representation of this quantity is good for the specific set of IC considered --- normalized at an amplitude which is typical in practice in cosmological simulations -- is correct. That is the agreement they observe in a modest range (see e.g. the figure 3 in \\cite{alvaro_knebe2}), from a few times the interparticle distance to a scale approaching the box size, at which finite size effects start to play a role, is real (rather than purely accidental as is implicitly suggested by \\cite{Baertschiger:2001eu, Baertschiger:2003jt}). However the dominant lattice term can clearly be identified at smaller scales, and it is evident in view of our discussion that the range of agreement will decrease (and ultimately disappear) if one considers the same model with a lower normalization.  For the two point correlation function, we have seen that the degree of agreement depends not only on the amplitude of the input model, but also on the details of the spatial smoothing in the estimator. Again neither collaboration has pinpointed explicitly the importance of this consideration in evaluating the faithfulness of the representation. The authors of \\cite{Baertschiger:2001eu, Baertschiger:2003jt}), however, are correct when they argue that the difference observed is a systematic one, and that the oscillating behavior observed is due to the correlations in the underlying pre-initial (lattice or glass) configuration. In attributing the difference to ``noise'' the  other group is incorrect, insofar as such noise would be a finite size effect which should disappear in the ensemble average. However, noise can of course play a crucial role in a finite sample in hiding the underlying signal in the regime in which it may, in principle, approximate well the theoretical model.  \\subsection{Physical relevance of results on IC}  We have considered in this paper solely the question of the accuracy with which the standard algorithm for generating IC for cosmological NBS represents the theoretical correlation properties. This question is essentially interesting insofar as it is relevant to the question addressed by the series of articles of which this is the first: the quantification of the differences between the results of {\\it evolved} N body simulations and the corresponding theoretical quantities. This question will be addressed fully in the subsequent papers, and we limit ourselves now to a partial discussion of the physical relevance of our findings.  The most important result from a practical point of view is that, at linear order in the theoretical density perturbations, there is a contribution to the PS of the IC additional to the theoretical PS. This is a source for gravitational structure formation through the Poisson equation, which in a given simulation cannot be separated from the theoretical term. Indeed we note that the linearity of this term in the amplitude of the relative displacements means that, if the early time evolution follows the Zeldovich approximation, this term is amplified linearly, just like the theoretical term. On the other hand, it contributes significantly only above the Nyquist frequency, and therefore, given that gravity tends to transfer power very efficiently from large to small scales (see, e.g., \\cite{Little+weinberg+park_1991}), one would expect its effects to be washed out over time. However if one wishes to quantify precisely discreteness effects, our quantification of this leading discreteness contribution in the IC is an important first step.  In quantifying such effects it is important also to first understand the recovery of the continuum limit. Our results here, as we will now discuss, actually are quite informative in this respect. Let us consider the limit in which one recovers {\\it exactly} the properties of the theoretical continuum model. Given an input theoretical model for a cosmological NBS, we introduce two parameters with the standard method of discretization we have discussed here~\\footnote{In reality there is of course also the box size $L$, which we have taken in our study to be infinite. The finite particle number $N$ is given by $(L/\\ell)^3$.}: $\\ell$, the lattice spacing in physical units, and the initial red-shift $z_i$ (which fixes the amplitude $A$ of the input PS, with $A \\rightarrow 0$ as $z_{\\rm i} \\rightarrow \\infty$).  The continuum limit should evidently correspond to taking $\\ell \\rightarrow 0$ (and thus $k_N \\rightarrow \\infty$). Let us consider first taking $\\ell \\rightarrow 0$ at fixed $z_{\\rm i}$. This corresponds in our analysis above to working at fixed amplitude of the PS. Our results above tell us that the representation of the PS is good provided we satisfy the condition Eq.~(\\ref{validity-criterion-2}) for the validity of the perturbative expansion. This quantity in fact converges to zero for $k_N \\gg k_c$, and so the criterion for good agreement in $k$ space for all $k$ is simply $\\Delta_{\\rm th}^2 (k_c) \\ll 1$. This agreement becomes arbitrarily good as we take $z_{\\rm i} \\rightarrow \\infty$ (i.e. $z_{\\rm i} \\rightarrow 0$ ). Likewise in real space, it follows from Eq.~(\\ref{condition-variance-n<1}) that we converge towards an arbitrarily good representation of the mass variance when the limit is taken in this way. The same is true of the two point correlation function. Thus the correlation properties of the discretized IC converge exactly to the continuum IC.  Our results concerning the differences in real space quantities concern the limit $z_{\\rm i} \\rightarrow \\infty$, at fixed $\\ell$. We have seen that there is, in this case, no convergence towards the continuum model. Thus, in the IC, the order of the limits in $\\ell$ and $z_{\\rm i}$ cannot be interchanged. It will be shown in the companion paper \\cite{discreteness2_mjbm}, that the same non-commutativity of the limits is observed in the evolved systems. This in fact is just a specific example of a well-known fact about the validity of continuum Vlasov dynamics to describe a system with long range interactions \\cite{spohn, braun+hepp}. In this context it is known and well documented in certain systems that the continuum limit is approached as $N \\rightarrow \\infty$ keeping the time of evolution fixed, while taking the time to infinity first one diverges from the collisionless limit (see, e.g., \\cite{yamaguchi_etal_04}). Lowering the initial amplitude of a NBS increases the time of evolution (up to a given time), and thus the behavior we are inferring from the analysis of the IC corresponds to this same one.  These comments on the continuum limit are also of practical relevance, as they tell us how one should study convergence to this limit numerically (in order to understand the precision of results). It follows from what we have just discussed that it is best to keep $z_i$ fixed as the particle density is increased. Further the continuum limit can only be defined clearly in the presence of a cut-off in the input PS, with the continuum limit being approached when the interparticle distance is decreased well below the inverse of this scale. In most of the numerical studies in the literature on discreteness in cosmological NBS these points have not been taken into account~\\footnote{An exception is some of the cited work of Melott et al.. Some sets of simulations are compared in which only the particle density is varied, keeping both the initial amplitude and the cut-off in the input PS fixed in units of the box size.}. Indeed we note that the very widely used standard software package COSMICS for generating IC \\cite{bertschingercode} fixes automatically the initial red-shift of the simulation when the physical particle density is given, rather than leaving it as a free parameter, making such controlled tests difficult. Indeed if no cut-off is imposed in the input PS, the criterion used to fix the red-shift makes it increase with the particle density. These points will be further discussed in forthcoming work.  We recall finally that our results on the limitations of the use of the algorithm for very blue spectra are of relevance to some studies in the literature of gravitational evolution from such spectra. Specifically we note their usefulness in understanding quantitatively results in \\cite{melott+shandarin_k4} and \\cite{bagla+padmanabhan_97}. These studies consider gravitational N body simulations (in two and three dimensions, respectively) starting from IC generated on a lattice using the standard algorithm discussed here, taking input theoretical PS with vanishing initial power in some range of small $k$: in \\cite{melott+shandarin_k4} a top-hat PS is used, while in \\cite{bagla+padmanabhan_97} a gaussian centred on a chosen wavenumber. In both cases our results show that there is a term proportional to $k^4$ induced at small $k$ {\\it already in the IC}, which will dominate at small $k$. The explicit expression for this term, which arises at second order in the expansion of the continuum piece $P_c(k)$ of the full PS, is given in Eq.~(\\ref{pc-second-four}). In \\cite{melott+shandarin_k4} the dominant contribution from the $k^4$ term in the IC at small $k$ is observed numerically,  and indeed the authors relate it (as discussed in Sect.~\\ref{Glass pre-initial conditions} above) to Zeldovich's argument  about ``minimal power''. In \\cite{bagla+padmanabhan_97}, on the other hand, the $k^4$ term is seen (and observed, as expected, to grow with an amplification proportional to the square of the linear growth factor) only after some time. The authors describe in this case the $k^4$ tail as ``generated'' by the dynamical evolution, which is evidently not quite accurate as the term is in fact present (albeit at lower amplitude) already in the IC.    We are indebted to Andrea Gabrielli for extensive discussions and explanations of his results reported in \\cite{andrea}. We thank Thierry Baertschiger and Francesco Sylos Labini for numerous useful conversations, and Alvaro Dominguez and Alexander Knebe for helpful comments on the first version of this paper. We are indebted to an anonymous referee for pointing out an important error in the first version of this paper.  \\appendix"
    },
    "0410/astro-ph0410498_arXiv.txt": {
        "abstract": "We study the temperature structure of the intergalactic medium (IGM) using a large cosmological $N$-body/SPH simulation. We employ a two-temperature model for the thermal evolution of the ionized gas, in which we include explicitly the relaxation process between electrons and ions. In the diffuse, hot IGM, the relaxation time is comparable to the age of the Universe and hence the electron temperature in post-shock regions remains significantly smaller than the ion temperature.  We show that, at the present epoch, a large fraction of the warm/hot intergalactic medium (WHIM) has a well-developed two-temperature structure, with typical temperature differences of order a few.  Consequently, the fraction of metals in various ionization states such as O{\\sc vi}, O{\\sc vii}, and O{\\sc viii}, as well as their line emissivities, can differ locally by more than an order of magnitude from those computed with a single-temperature model, especially in gas with $T \\sim 10^7$ K.  It is thus necessary to follow the evolution of the electron temperature explicitly to determine absorption and emission by the WHIM.  Although equipartition is nearly achieved in the denser intracluster medium (ICM), we find an appreciable systematic deviation between the gas-mass weighted electron temperature and the mean temperature even at half the virial radii of clusters.  There is thus a reservoir of warm ($T_{e}<1$ keV) gas in and around massive clusters. Our results imply that relaxation processes need to be considered in describing and interpreting observational data from existing X-ray telescopes as well as from future missions designed to detect the WHIM, such as the {\\sl Diffuse Intergalactic Oxygen Surveyor} and the {\\sl Missing Baryon Explorer}. ",
        "introduction": "The distribution of baryons in the Universe remains one of the puzzling issues in modern cosmology.  At the present epoch, the total amount of baryons inferred from a census of H{\\sc i} absorption, gas and stars in galaxies, and X-ray emission from hot gas in galaxy clusters is far smaller than that predicted by nucleosynthesis calculations (Persic \\& Salucci 1992; Fukugita, Hogan \\& Peebles 1998) and by measurements of the cosmic microwave background radiation (Spergel et al. 2003).  Hence, it is believed that a large fraction of the baryons lies in an as yet undiscovered {\\it dark} state.  Three dimensional hydrodynamic simulations of cosmic structure formation (e.g., Cen \\& Ostriker 1999; Dav\\'e et al. 2001; Croft et al. 2002) predict that about 30--40\\% of the baryons reside in the so-called warm-hot intergalactic medium (WHIM) at the present epoch. This gas is mostly shock-heated to a temperature of $\\sim 10^5$--$10^7$ K during large scale structure formation, and this relatively low temperature makes its thermal emission difficult to detect with conventional X-ray probes.  A variety of observational approaches have been considered for studying the WHIM, including both ultraviolet (e.g., Furlanetto et al. 2003, 2004) and X-ray lines.  With respect to the latter possibility, there have been several tentative claims that the WHIM has been detected locally in absorption (e.g., Nicastro et al. 2002) and in emission (Finoguenov et al. 2003).  Moreover, Yoshikawa et al. (2004) have proposed using O{\\sc vii}/O{\\sc viii} lines to examine the WHIM in detail with high spectral resolution X-ray detectors.  They argue that, under a variety of assumptions, about half of the WHIM (by mass) can be detected via oxygen line emission (see also Fang et al. 2004). While the proposed missions appear promising, the detectability depends crucially on the assumed metallicity and the relative abundances of particular ions.  Previous theoretical studies used a single-temperature model for the WHIM, under the assumption that equipartition is fully achieved. However, the outskirts of galaxy clusters may actually have a two-temperature structure (e.g., Fox \\& Loeb 1997; Takizawa 1998; Courty \\& Alimi 2004).  Near galaxy clusters, infalling gas shock-heats to roughly the virial temperature as kinetic energy is converted into thermal energy.  Because the infalling ions carry nearly all of the bulk kinetic energy (exceeding that contained in electrons by a factor $\\sim m_p/m_e$), the post-shock electrons can gain thermal energy only via collisions with ions. If collisional relaxation takes a time comparable to or longer than the age of the Universe, the two temperatures, i.e. those of the electrons and ions, remain different.  One can easily show that the relaxation time is rather long for diffuse warm-hot gas around clusters and in large-scale filamentary structures.  Since the relative abundance of metal ions and their emissivities are sensitive to the electron temperature, this has important implications for future observations of the WHIM, particularly for the hotter component.  In this {\\sl Letter}, we study the thermal evolution of the intergalactic medium (IGM) using a large hydrodynamic \\begin{figure*}[t] \\epsscale{2.0} \\plotone{whim.eps} \\caption{The distribution of baryons (left), the warm/hot component with $10^6 < T < 10^7$ K (middle), and gas with $T_{\\rm e}< 0.5T_{\\rm i}$ (right) in a slab of $100\\times 100\\times 20$ ($h^{-1}$Mpc)$^3$ at $z=0$. The vertical bar in the middle panel indicates a `super-cluster' region, for which we compute the soft X-ray intensity (see section 4 and Figure 5). \\label{plot1}} \\end{figure*} \\noindent simulation of structure formation.  We show that a considerable fraction of the WHIM indeed has a two-temperature structure, particularly in and around rich clusters.  In these regions, the electron temperature is typically smaller than the ion temperature by a factor of a few, substantially modifying estimates of the abundances of ions and their emissivities. ",
        "conclusions": ""
    },
    "0410/astro-ph0410721_arXiv.txt": {
        "abstract": "We use the first 25\\% of the DEEP2 Galaxy Redshift Survey spectroscopic data to identify groups and clusters of galaxies in redshift space.  The data set contains 8370 galaxies with confirmed redshifts in the range $0.7\\le z\\le 1.4$, over one square degree on the sky. Groups are identified using an algorithm (the Voronoi-Delaunay Method) that has been shown to accurately reproduce the statistics of groups in simulated DEEP2-like samples. We optimize this algorithm for the DEEP2 survey by applying it to realistic mock galaxy catalogs and assessing the results using a stringent set of criteria for measuring group-finding success, which we develop and describe in detail here. We find in particular that the group-finder can successfully identify $\\sim78\\%$ of real groups and that $\\sim79\\%$ of the galaxies that are true members of groups can be identified as such.  Conversely, we estimate that $\\sim 55\\%$ of the groups we find can be definitively identified with real groups and that $\\sim46\\%$ of the galaxies we place into groups are interloper field galaxies. Most importantly, we find that it is possible to measure the distribution of groups in redshift and velocity dispersion, $n(\\sigma,z)$, to an accuracy limited by cosmic variance,  for dispersions greater than $350\\kms$.  We anticipate that such measurements will allow strong constraints to be placed on the equation of state of the dark energy in the future.  Finally, we present the first DEEP2 group catalog, which assigns 32\\% of the galaxies to 899 distinct groups with two or more members, 153 of which have velocity dispersions above $350\\kms$.  We provide locations, redshifts and properties for this high-dispersion subsample.  This catalog represents the largest sample to date of spectroscopically detected groups at $z\\sim 1$. ",
        "introduction": "\\label{sec:intro}  Groups and clusters of galaxies are the most massive dynamically relaxed objects in the universe; as such, they have long been the subject of intense and fruitful study. More than seventy years ago observations of the Coma cluster gave the first evidence for the existence of dark matter~\\citep{Zwicky}. More recently, studies of gravitational lensing by clusters have yielded intriguing new information about the profiles of dark matter halos~\\citep{Sand}. Identifying and studying galaxies in groups and clusters is essential to understanding the effects of local environment on galaxy formation and evolution \\citep[for a review see][]{BB}. X-ray measurements of the gas mass fraction in clusters have been used to constrain the mass density parameter $\\Omega_M$ and more recently the equation of state of the dark energy, $w$~\\citep[\\emph{e.g.}, ][and references therein]{Allen}. Finally, if we can accurately measure the abundance of groups \\emph{and its evolution with redshift}, we can constrain the growth of large-scale structure, thereby placing significant further constraints on cosmological parameters \\citep{Lilje, ECF, Borgani, HMH, HHM, NMCD}.   A wide array of methods has been used to identify groups and clusters at moderate redshifts: X-ray emission from hot intracluster gas~\\citep[reviewed by][]{RBN}, cosmic shear due to weak gravitational lensing~\\citep[reviewed by][]{fridge}, searches in optical photometric data~\\citep[\\emph{e.g.}, ][]{LCDCS, YG}, the Sunyaev-Zel'dovich (SZ) effect in the Cosmic Microwave Background~\\citep[\\emph{e.g.}, ][]{LaRoque}, and direct reconstruction of three-dimensional objects in galaxy redshift surveys~\\citep[\\emph{e.g.}, ][]{Eke}.  To study the evolution of the group abundance it is necessary to extend observations to more distant objects.  However, most of the methods used for local studies have only limited effectiveness at high redshift.  The apparent surface brightness of X-ray clusters dims as $(1+z)^{-4}$, making only the richest clusters visible at high redshift. The cross-section for gravitational lensing falls rapidly at high redshifts, making weak-lensing detection of distant clusters difficult for all but the most massive objects.  In photometric surveys, the increased depth necessary for high-redshift studies increases the overall number density of objects, thereby increasing the problems of foreground and background contamination and projection effects (though photometric techniques for estimating redshifts can mitigate these difficulties).  The SZ effect is very promising, since it is entirely independent of redshift, but it also suffers from confusion limits and projection effects, and in any case a large survey of SZ clusters is yet to be undertaken.  For the time being, then, one of the few methods that can be applied to large numbers of groups and clusters on similar mass scales at $z\\sim 0$ and $z\\sim 1$ is the direct detection of these structures in the redshift-space distribution of galaxies.   The first sizeable sample of groups detected in redshift space was presented by~\\citet{GH}, who found 176 groups of three or more galaxies in the CfA galaxy redshift survey at redshifts $z\\la 0.03$. Recently, \\citet{Eke} identified groups within the final data release of the Two-degree Field Galaxy Redshift Survey (2dFGRS).  Their catalog extends to $z\\approx 0.25$ and constitutes the largest currently available catalog of galaxy groups, containing $\\sim 3\\times 10^4$ groups with two or more members.  A comprehensive listing of previous studies of local optically selected group samples is also given by these authors.   Work is currently underway to detect groups of galaxies in the spectroscopic data of the Sloan Digital Sky Survey~\\citep{Nichol}. Studies of groups detected in redshift space were extended to intermediate redshifts by~\\citet{Carlberg}, who found more than 200 groups in the CNOC2 redshift survey, with redshifts in the range $0.1\\le z\\le 0.55$. Additionally, \\citet{Cohen} studied a sample of 23 density peaks in redshift space, over the redshift range $0\\le z\\le 1.25$.  Until now, however, no spectroscopic sample has existed with sufficient size, sampling density and redshift accuracy to extend redshift-space studies of large numbers of groups to redshifts $z\\ga 0.5$.    The DEEP2 Galaxy Redshift survey \\citep{DEEP2, Faber} is the first large spectroscopic galaxy catalog at high redshift, with observations planned for $\\sim 5\\times 10^4$ galaxies, most of which will fall in the range $0.7\\le z \\le 1.4$.  The survey is thus a unique dataset for studying high-redshift galaxy groups. With such a broad range in redshift, we expect to observe evolution in the properties of galaxies, groups, and clusters within the DEEP2 sample itself; also, by comparing to local samples from 2dFGRS and the Sloan Digital Sky Survey (SDSS), we expect to observe evolution between $z\\sim 1$ and the present epoch.  DEEP2 is especially well-suited to studies of groups, since its high redshift accuracy allows detailed studies of their internal kinematics.  Repeated observations of some DEEP2 galaxies indicate a velocity accuracy $\\delta v \\approx 25 \\kms$, considerably better than the 2dFGRS value $\\delta v \\approx 85 \\kms$~\\citep{2dFspec} and similar to the $\\delta v \\approx 30 \\kms$ attained in the SDSS~\\citep{EDR}.  In particular, it will be possible to estimate the masses of DEEP2 groups from their velocity dispersions.  We anticipate that, by measuring the evolution of the group velocity function with redshift, it will be possible to constrain cosmological parameters such as the dark energy density parameter $\\Omega_\\Lambda$ and equation of state parameter $w$, as outlined in \\citet{NMCD}.  Before carrying out such studies, however, it will be essential to develop robust methods for detecting groups and clusters within DEEP2.   Identifying groups and clusters in redshift space is well known to be a difficult task.  Most notably, clustering information is smeared out by redshift space distortions like the so-called fingers-of-God effect, in which galaxies in groups and clusters appear highly elongated along the line of sight because of intracluster peculiar motions.   This intermingles group members with other nearby galaxies and causes neighboring groups to overlap in redshift space. A group-finding algorithm that attempts to find all group members will thus necessarily be contaminated by interloper field galaxies and will necessarily merge some distinct groups together into spurious larger structures.  Conversely, a group finder that aims to minimize contamination and over-merging will fragment some larger clusters into smaller groups.  This trade-off in group-finding errors is well known \\citep[\\emph{e.g.}, ][]{NW} and fundamentally unavoidable. It will therefore be essential, before we begin any program of group finding, to construct a suitable definition of group-finding success, identifying in advance which errors we seek to minimize and what sort of errors we are willing to tolerate.  Ultimately, the chosen definition of success will depend on the intended scientific purpose of the group catalog. A major portion of this paper will be devoted to defining appropriate measures of group-finding success for the DEEP2 survey and optimizing our methods using these criteria.  This paper is organized as follows.  In Section~\\ref{sec:DEEP2} we introduce the DEEP2 sample and discuss the unique opportunities and difficulties it presents for group finding.  In the same section we describe the DEEP2 mock galaxy catalogs, which we will use to calibrate our group-finding methods. Then, in Section~\\ref{sec:optimal}, we describe our criteria for group-finding success.  In Section~\\ref{sec:gfchoice}, we give an overview of various group-finding methods that have been used in the literature, and we describe the Voronoi-Delaunay Method (VDM) of ~\\citet{Marinoni}, which we will use in this study.   We then proceed to optimize this method for the DEEP2 sample.  Finally, in Section~\\ref{sec:catalog}, we apply the VDM algorithm to the current DEEP2 observations and present the first DEEP2 group catalog. ",
        "conclusions": "\\label{sec:conclusion}  We have optimized the VDM group-finding algorithm using mock catalogs designed to replicate the DEEP2 survey, and we have applied it to spectroscopic data from three DEEP2 photometric pointings.  In the process of optimization, we have defined a measure of group finding success that focuses on accurately reproducing the overall properties of the group catalog---in particular the distribution of groups in redshift and velocity dispersion, $n(\\sigma, z)$---while paying some attention also to the accurate reconstruction of individual groups.  Tests on DEEP2 mock catalogs show that we are able to accurately reproduce $n(\\sigma, z)$ for $\\sigma \\ge 350\\kms$ and that errors in measuring this quantity in DEEP2 should be smaller than its expected intrinsic cosmic variance.  It should thus be possible to use the test described by \\citet{NMCD} to constrain cosmological parameters, including the dark energy equation of state parameter, $w$.  We find 899 groups with two or more members  within the DEEP2 data considered in this study, roughly $25\\%$ of the expected final sample.  Of these, 153 have velocity dispersions $\\sigma \\ge 350\\kms$.  The distribution of these reconstructed groups with velocity dispersion $n(\\sigma)$ is in good agreement with the distribution for real groups in DEEP2 mock catalogs.   This result provides a useful consistency check for the mock catalogs: assuming our reconstructed $n(\\sigma)$ is accurate (as our tests show that it is for $\\sigma \\ge 350\\kms$), we may be confident that the properties of groups in the mock catalogs are an accurate simulation of real DEEP2 data.  This is especially important in the context of the so-called velocity bias, which is the ratio of the velocity dispersion of galaxies to that of the underlying dark matter halo, $b_v = \\sigma_{\\mathrm{gal}}/\\sigma_{\\mathrm{DM}}$.  Various studies of N-body simulations \\citep[\\emph{e.g.}, ][and references therein]{Diemand} have suggested that $b_v \\ne 1$ at the 15--30\\% level, but no such effects have been included in the DEEP2 mock catalogs.  Our results on $n(\\sigma)$ thus indicate that our data are consistent with $b_v = 1$ within the measurement errors shown in Figure~\\ref{fig:dist_comp}.  This is no surprise, since our current error bars are considerably larger than the expected effect; nevertheless, this result may be viewed as confirmation that no stronger biases exist.   Successful detection of groups and clusters within the DEEP2 redshift survey is an essential first step for a wide variety of planned studies.  By comparing the properties of galaxies in groups to the properties of isolated galaxies at high redshift, we can learn much about galaxy formation and evolution.  This will be discussed further in an upcoming paper~\\citep{properties}.  Also, with the catalog of groups we now have in hand, it is possible to pursue targeted follow-up observations in X-rays or using the Sunyaev-Zeldovich effect to better constrain the gas physics of groups at high redshift; such programs are now being developed, including an upcoming X-ray survey of the extended Groth Strip field with the \\emph{Chandra} space telescope.  Finally, using the groups we find in our current and future spectroscopic data, we expect to put strong new constraints on the formation and evolution of galaxies, groups, and clusters, and to investigate the makeup of the universe, including the nature of the dark energy."
    },
    "0410/astro-ph0410517_arXiv.txt": {
        "abstract": "In this poster we present a detailed chemical analysis of the RV\\,Tauri star of low intrinsic metallicity, V453 Oph ([Fe/H]=$-2.2$), which shows a mild but clear s-process surface overabundance. This result is strengthened by a relative analysis to the RV Tauri star DS Aqr of similar metal deficiency ([Fe/H]=$-1.6$), but without such s-process enrichment. A remarkable result is that the s-process enrichment of V453 Oph is \\emph{not} accompanied by a measurable C enhancement. To explain the surface abundances observed in V453 Oph, two nucleosynthesis models are considered. The first model assumes the partial mixing of protons during the radiative interpulse phase while the second scenario is based on the possible production of s-process elements within the flash-driven convective pulse only. Both models can give a satisfactory explanation for the s-process enrichment of V453\\,Oph, but not the simultaneous low C-abundance observed. ",
        "introduction": "RV\\,Tauri stars are rare supergiants which occupy the high luminosity end of the pop. II Cepheid instability strip. They are identified on the basis of their characteristic light curves, which show alternating deep and shallow minima with periods between 30 and 150 days. Their post-AGB nature is established on the basis of their far-IR excess due to circumstellar dust, which is thought to be a relic of the dusty mass loss at the end of the AGB. The genuine high luminosity nature of this RV\\,Tauri class was established by the detection of some members in Globular Clusters and the LMC.  Post-AGB objects are very useful in constraining the AGB chemical evolutionary models, because the absence of strong molecular veiling allows detection of atomic transitions from CNO up to the heavy s-process elements beyond the Ba-peak. ",
        "conclusions": "This paper reports upon the abundance analysis of a very particular s-process enriched object: V453\\,Oph. The abundances derived for this object are of particular nucleosynthesis interest, since for the first time, the s-process enrichment is not accompanied by a simultaneous C-enrichment. This observation provides us with clear proof that our knowledge about the AGB nucleosynthesis is still unsatisfactory because no current model can fully reproduce the specific history of this object."
    },
    "0410/astro-ph0410271_arXiv.txt": {
        "abstract": "We have undertaken a FUSE survey of Magellanic Cloud supernova remnants, looking primarily for O~VI and C~III emission lines. Work in earlier cycles with FUSE indicates that optical and/or X-ray characteristics of supernova remnants are not always good predictors of the objects that will be bright and detectable in the UV. The goal of our survey is to test this concept by obtaining spectra of a random sample of Magellanic Cloud remnants with a broad range of radio, optical, and X-ray properties. Previously observed objects and remnants with known high extinction (or known high column densities) are the only objects eliminated from consideration. To date, we have clearly detected O~VI emission from 12 of the 33 objects observed, with weak or marginal detections possible in a handful more. ",
        "introduction": "In the study of emission line objects, the far ultraviolet (FUV) is typically viewed as an extension of the optical.  Consequently, FUV observations tend to target objects (or parts of objects) which exhibit strong optical line emission. This criterion has been frequently used in the field of supernova remnant (SNR) research to select suitable targets for ultraviolet observations.  There could be many Magellanic Cloud SNRs detectable at a level useful for analysis in the FUV, but which have never been observed because of selection biases. To avoid the selection bias of targeting bright optical/X-ray remnants, we are currently performing a FUSE Guest Investigator survey of all known LMC and SMC SNRs, excluding only those which have been previously observed or are known to exhibit high interstellar extinction.  Most of these targets have not previously been observed in the UV.  In our observations a standard exposure of 10 ks with the LWRS aperture is requested, allowing us to detect O~VI emission up to 50 times fainter than observed in such bright SNRs as N49 (Blair et al. 2000).  The target coordinates are at the radio or X-ray center of each SNR.  With this observing setup, even non-detections are of significant interest. Most of the objects are from 1\\arcmin\\ to 5\\arcmin\\ in extent, which will fill the LWRS aperture (30\\arcsec $\\times$ 30\\arcsec) with emission. This setup also lessens the impact of relative FUSE channel misalignments and keeps the observations simple.  In addition, Doppler shifting from the front and back sides of the SNR shells ($\\sim\\,$30$-$150 km s$^{-1}$) will move their emission lines out from under the overlying absorption of the host galaxy. ",
        "conclusions": "Thus far we have detected O~VI line emission from over 1/3 of the Magellanic Cloud remnants observed in our survey.  This is an ongoing project, and will yield useful constraints on the kinematics of MC SNRs. A broad assortment of data is available to support our FUSE observations: Narrow band MCELS imaging of many LMC/SMC fields of interest are in hand, while longslit echelle observations cutting across many Magellanic Cloud H~II regions, bubbles, etc., are available to provide an excellent base of general kinematic information on various regions of potential interest.  X-ray observations with {\\it Chandra}, both in the archive and in ongoing programs, will provide invaluable information on shocked ejecta and blast wave emission in these SNRs.  In addition, upcoming {\\it Spitzer} Cycle 1 observations will provide supplemental information on shocked dust in the LMC SNRs."
    },
    "0410/astro-ph0410047_arXiv.txt": {
        "abstract": "The regulation of angular momentum is one of the key processes for our understanding of stellar evolution. The rotational evolution of solar-mass stars is mainly determined by the magnetic interaction with their circumstellar disk and angular momentum loss through stellar winds, and In contrast to solar-mass stars, very low mass (VLM) objects and brown dwarfs are believed to be fully convective. This may lead to major differences of rotation and activity, since fully convective objects may not host a solar-type dynamo.  Here, we report on our observational efforts to understand the rotational evolution of VLM objects. By means of photometric monitoring, we determined 62 rotation periods for targets in three clusters, which form an age sequence from 3 to 125\\,Myr. We find that VLM objects rotate faster than their solar-mass siblings in all evolutionary stages. Their rotational evolution seems to be determined by hydrostatic contraction and exponential angular momentum loss. The photometric amplitudes of the light curves are much lower than for solar-mass stars. This may be explained as a consequence of smaller spot coverage, more symmetric spot distributions, or lower contrast between spots and their environment. Most of these results can be explained with a change of the magnetic field properties with decreasing mass. VLM objects possibly possess only small-scale, turbulent magnetic fields. ",
        "introduction": "Rotation is one of the key parameters for stellar evolution. It is the parameter that -- at least in some well-behaved objects -- can be measured to the highest accuracy. In such objects derivations of rotation periods with a precision of 1 : 10000 are possible.  In solar mass stars the investigation of their rotation has allowed us new insights into their evolution (\\cite{b95}, \\cite{bfa97}, \\cite{st03}). It has become clear that angular momentum regulation is a direct consequence of basic stellar physics: most of the angular momentum of a fragmenting and collapsing molecular cloud is lost in the course of the formation of protostars. The specific angular momentum of protostars is, however, still one or two orders of magnitude higher than that of young main sequence stars. On the other hand, in their T\\,Tauri phase solar-mass stars rotate slowly although they are accreting. The magnetic coupling between the star and its disk, and consequent angular momentum removal in a highly collimated bipolar jet are thought to be responsible for this rotational braking (\\cite{c90}, \\cite{k91}, \\cite{snowrl94}). After the dispersal of the disk, and thus loss of the braking mechanism, the rotation is observed to accelerate as the stars contract towards the zero-age main sequence (ZAMS). On the main sequence (MS) the rotation rates of solar-mass stars decrease again because of angular momentum loss through stellar winds.  Rotation can be investigated either by measuring stellar photospheric lines spectroscopically, or by the determination of rotation periods from photometric time series observations. While the former suffers from projection effects -- the unknown inclination angle of the rotation axis with respect to the line of sight -- the latter can be determined with high precision and independent of inclination angle.  While a large amount of rotation periods are available in the literature for low-mass stars in clusters younger than about 3\\,Myr (ONC: \\cite{herbst01}, \\cite{hbm02}; NGC2264: \\cite{l03}, \\cite{lbm04}), not much is known about the further evolution of very low mass (VLM) objects and brown dwarfs up to the age of a few Gyr, when they are found as field objects in the solar neighbourhood (\\cite{ctc02}). Therefore, we have initiated a programme to obtain the required observations, and to compare them to solar-mass stars and to evolutionary models. For the monitoring programme we decided to follow the photometric time series approach to obtain precise rotation periods.  In this article we will first discuss in Section 2 our targets, observations, and data analysis. In Sect.\\,3 we present our results on rotation and variability of the sources. The observed rotation rates are then in Sect.\\,4 compared to various models of rotational evolution. Finally, Sect.\\,5 contains our conclusions. ",
        "conclusions": "We report results from our ongoing photometric monitoring of VLM objects -- so far in the clusters around \\object{sigma Ori}, \\object{epsilon Ori}, and the \\object{Pleiades}, and first attempts to model their rotational evolution.  The observed periodic variability of many VLM objects is likely caused by magnetically induced cool spots on the surfaces of the objects. In particular in the \\object{Pleiades}, we show that variation amplitudes are lower in VLM objects than in solar-mass stars, indicating either less asymmetric spot distribution, smaller relative spotted area, or lower contrast between spots and average photosphere. VLM objects show shorter rotation periods with decreasing mass, which is observed already at the youngest ages, and hence must have its origin in the earliest phases of their evolution.  Combining the rotation periods for all our objects, we find that their evolution does not follow hydrostatic contraction alone, but some kind of braking mechanism, e.g. wind braking similar to the one observed in solar-mass stars, is required as well. Such a wind braking is intimately connected to stellar activity and magnetic dynamo action (\\cite{schatzman62}). On the other hand, all the investigated VLM objects are thought to be fully convective , and therefore may not be able to sustain a solar-type large-scale dynamo, which is at the heart of the Skumanich type angular momentum loss of solar-mass stars. In fact, our modeling shows that such a Skumanich type wind braking cannot explain our data, while saturated angular momentum loss following an exponential braking law can. This, and the observed small photometric amplitudes may advocate a small-scale magnetic field configuration, and may support turbulent dynamo scenarios. Consistent theoretical models of such dynamos that would permit rigorous testing against the observations are, however, not yet available."
    },
    "0410/astro-ph0410337_arXiv.txt": {
        "abstract": "Projected rotational velocities (\\textit{v sin i}) have been measured for 216 B0--B9 stars in the rich, dense h and $\\chi$ Persei double cluster and compared with the distribution of rotational velocities for a sample of field stars having comparable ages ($t \\sim$ 12-15 Myr) and masses (M $\\sim$ 4-15 $M\\sun$). For stars that are relatively little evolved from their initial locations on the Zero Age Main Sequence (those with masses M $\\sim$ 4-5 $M\\sun$), the mean \\textit{v sin i} measured for the h and $\\chi$ Per sample is slightly more than 2 times larger than the mean determined for field stars of comparable mass, and the cluster and field \\textit{v sin i} distributions differ with a high degree of significance.  For somewhat more evolved stars with masses in the range 5-9 $M\\sun$, the mean \\textit{v sin i} in h and $\\chi$ Per is 1.5 times that of the field; the \\textit{v sin i} distributions differ as well, but with a lower degree of statistical significance.  For stars that have evolved significantly from the ZAMS and are approaching the hydrogen exhaustion phase (those with masses in the range  9-15 $M\\sun$), the cluster and field star means and distributions are only slightly different. We argue that both the higher rotation rates and the pattern of rotation speeds as a function of mass that differentiate main sequence B stars in h and $\\chi$ Per from their field analogs were likely imprinted during the star formation process rather than a result of angular momentum evolution over the 12-15 Myr cluster lifetime. We speculate that these differences may reflect the effects of the higher accretion rates that theory suggests are characteristic of regions that give birth to dense clusters, namely: (a) higher initial rotation speeds; (b) higher initial radii along the stellar birthline, resulting in greater spinup between the birthline and the ZAMS; and (c) a more pronounced maximum in the birthline radius-mass relationship that results in differentially greater spinup for stars that become mid- to late- B stars on the ZAMS. ",
        "introduction": "Much of our current understanding of how stars form derives from the study of nearby star-forming regions such Taurus-Auriga, Ophiuchus, and Chamaeleon. These regions are populated by $\\sim$ 100 young stars having typical masses $M < 1 M\\sun$ contained within irregular molecular cloud complexes that span regions of size $\\sim$ 3-10 pc.  However, the demographics and morphologies of these complexes differ markedly from those thought to produce the majority of stars over the history of the universe: dense stellar clusters containing $10^{4}$ to $10^{6}$ stars having masses ranging from 100 to 0.1 $M\\sun$ formed within regions no more than $\\sim$ 1pc in size. Do the dramatic differences in stellar density between these regions influence measurable properties of individual stars and the statistical properties (e.g. the initial mass function) of the emerging stellar populations? If so, what are the key physical causes? Answering these questions represents an essential first step toward developing a predictive theory of star formation of sufficient power to inform our understanding of how the mix of high and low mass stars populating galaxies  today came to be, and as a consequence how the observed relative abundances of heavy elements were established.  The importance and timeliness of these questions have stimulated several recent theoretical studies aimed at predicting initial protostellar conditions in dense star-forming complexes (e.g. McKee \\& Tan 2003), and their relationship to emergent stellar mass functions (e.g. Elmegreen \\& Shadmehri  2003). Dense stellar clusters form in regions of very high gas surface density characterized as well by close packing of protostars. The turbulent velocity of the gas in these regions is likely to be high, leading to (1) protostars of high initial density; (2) rapid protostellar collapse times; and (3) as a consequence, high time-averaged  accretion rates. The latter may be conducive to the formation of very high mass stars, since the dynamical pressure of accreting material can be high enough to overcome radiation pressure from the forming massive stars (McKee \\& Tan 2003). The combination of protostellar cores characterized by higher turbulent speeds and higher mass accretion rates, combined with collisions between closely packed cores, may in turn produce a 'top-heavy' IMF (Elmegreen \\& Shadmehri 2003).  These theoretical studies thus suggest: (1) higher time-averaged accretion rates; and (2) an initial mass function biased toward higher mass stars in high density stellar clusters. In principle, accretion rates characteristic of different star-forming regions can be diagnosed from the location of the stellar birthline (e.g. Stahler 1988) as determined from spectroscopic and photometric observations of precision sufficient to locate pre-main sequence stars spanning a range of masses in the HR diagram, provided that the target regions are young and accurate age estimates are available. Quantifying initial mass functions requires similar observations.  To date, it has not been possible to determine either birthline locations or IMFs spanning the full range of stellar masses primarily because the best and closest examples of high density clusters (e.g. Arches at the Galactic Center, Stolte et al. 2002; R 136 in the Magellanic  Clouds, Massey \\& Hunter 1998; Sirianni et al. 2000) suffer from extreme crowding, which thus far has limited photometric and spectroscopic observations to main-sequence stars and a few pre-main-sequence stars with masses above $\\sim2 M\\sun$. Such stars are both bright enough to stand out relative to the stellar background and rare enough to avoid overlap with objects of comparable brightness. Next generation adaptive optics systems on current generation large telescopes have the potential to overcome the limitations of crowding and enable determination of stellar luminosities and effective temperatures for stars with masses as small as 0.1 $M\\sun$, a level more than sufficient to confront theoretical predictions of IMF shape and birthline positions. However, until such systems become operational, other and less direct approaches must suffice. We explore here the possibility that the distribution of stellar rotational velocities can provide a surrogate indicator of the differences in initial conditions between low and high density star-forming regions. Our reasoning is as follows.  Current theory suggests that initial stellar angular momenta are established during the primary stellar accretion phase via locking of stellar angular velocity to the angular velocity of the circumstellar accretion disk at or near the radius, r(m), where the stellar magnetosphere links to the disk (e.g. K\\\"{o}nigl 1991; Shu et al. 1994). That radius is set by the balance between the dynamical pressure of accreting material and the magnetic pressure of the magnetic field rooted in the forming star. For a fixed stellar magnetic field strength, the higher the accretion rate through the disk, the smaller r(m), the higher the Keplerian rotation speed of the disk at r(m), and hence the higher the angular rotational speed of the star.  Hence, if accretion rates are higher in high density star-forming regions, the resulting stellar population would be expected to exhibit higher rotation speeds on average.  Supporting observational evidence in the literature is sparse, but highly suggestive. For example, Wolff, Edwards, \\& Preston (1982; hereafter WEP) note that the distribution of rotation speeds among B stars in the relatively dense Orion Nebula Cluster is significantly shifted toward higher values as compared to stars of similar type distributed among the much lower density regions of the Orion star-forming complex. Similarly, Slettebak (1968) argues that stellar rotation speeds among the luminous B giant stars located in the vicinity of the extremely dense h and $\\chi$ Persei double cluster are $\\sim$ 50\\% higher than their field counterparts. Moreover, he reports an unusually high number of rapidly rotating Be stars, possibly indicative of a higher fraction of rapidly rotating stars in h and $\\chi$ Per. However, the Orion study includes only a modest sample of stars, while past discussions of h and $\\chi$ Per rotation properties suffer from concerns regarding  the similarity in age range among field and cluster cohorts combined with  the possibility of evolutionary changes in stellar angular momenta. Confronting the hypothesis that stars born in dense clusters rotate more rapidly than their field counterparts requires sufficiently large samples of cluster and field stars spanning an identical range of ages.  The goal of this contribution is to effect a robust statistical comparison of the distribution of stellar rotational velocities for a sample of B0-B9 stars in h and $\\chi$ Persei (typical stellar density of $10^{4}$ pc$^{-3}$; Slesnick, Hillenbrand, \\& Massey, 2002, hereafter SHM) with those observed among field B stars of comparable age as determined from their location in the Str\\\"{o}mgren  ($\\beta$, c$_{0}$) plane. This latter sample is almost certainly dominated by stars born in much lower density environments. ",
        "conclusions": "Observations of rotational velocities show that B stars in the h and $\\chi$ Per double cluster rotate on average more rapidly than field stars of the same mass and age. This result combined with other similar results in the literature clearly establishes that the rotation rates of B stars \\textit{differ significantly} among stars born in different regions. We have argued that the observed differences are \\textit{likely} built in at the time the stars were formed and not a consequence of subsequent evolutionary processes. Since h and $\\chi$ Per are currently bound clusters, it seems reasonable to assume that stars in these clusters were formed in much denser environments than field B stars of comparable age, which have presumably escaped from the loose clusters, associations, or small groups in which they formed. We have identified two possible explanations for the observed differences in rotation speeds, each related to the density at the time of star formation. One hypothesis is essentially untestable with current facilities: that h and $\\chi$ Per are deficient in wide binaries relative to the field and that when wide binaries are formed, much of the available angular momentum appears as orbital rather than stellar rotational angular momentum. The second possibility is that protostellar accretion is more rapid in higher density regions, and that high accretion rates lead to more rapid rotation---a consequence of the hypothesis that ``disk locking\"\" accounts for initial stellar angular momenta. This hypothesis is attractive because the effects of differences in accretion rate during the stellar assembly phase are predicted to be particularly dramatic among B stars. For stars in the mass range 4-12 $M\\sun$ , the initial radius along the stellar birthline reflects the effects of deuterium shell burning, which causes a substantial expansion of the star; the amount of the expansion and the range of stellar masses over which it occurs both increase with increasing accretion rates. The combination of high initial accretion rate with high initial radius followed by contraction from the birthline to the ZAMS can lead to high ZAMS rotation speeds compared to stars of similar mass formed in regions characterized by lower accretion rates.  We emphasize that our results for h and $\\chi$ Per, while suggestive, do not provide conclusive evidence of a direct relationship between environment and outcome rotational velocities. With the advent of 8- to 10-m telescopes, it should be possible to test the hypothesis that outcome stellar rotation speeds are in fact linked to initial stellar density among stars born in a wide variety of environments by observing a large sample of stars located in young clusters both the in Milky Way and in the Magellanic Clouds. Such observations will provide the basis for establishing robustly both the cosmic dispersion in N(vsini) as well as systematic differences attributable to initial stellar density, chemical composition, or other parameters. For regions of sufficient youth, it may also be possible to search for differences in the location of the birthline. If the predicted correlations between rotation speed, birthline location, and environment can be found, it would be possible for the first time to link the initial conditions under which star formation takes place and to outcome observables---a crucial first step toward a predictive theory of star formation.   Facilities: \\facility{WIYN}."
    },
    "0410/astro-ph0410101_arXiv.txt": {
        "abstract": "{We explore the rotation and activity of very low mass (VLM) objects by means of a photometric variability study. Our targets in the vicinity of $\\epsilon$\\,Ori belong to the Ori\\,OB1b population in the Orion star-forming complex. In this region we selected 143 VLM stars and brown dwarfs (BDs), whose photometry in RIJHK is consistent with membership of the young population. The variability of these objects was investigated using a densely sampled I-band time series covering four consecutive nights with altogether 129 data points per object. Our targets show three types of variability: Thirty objects, including nine BDs, show significant photometric periods, ranging from 4\\,h up to 100\\,h, which we interpret as the rotation periods. Five objects, including two BDs, exhibit variability with high amplitudes up to 1\\,mag which is at least partly irregular. This behaviour is most likely caused by ongoing accretion and confirms that VLM objects undergo a T Tauri phase similar to solar-mass stars. Finally, one VLM star shows a strong flare event of 0.3\\,mag amplitude. The rotation periods show dependence on mass, i.e. the average period decreases with decreasing object mass, consistent with previously found mass-period relationships in younger and older clusters. The period distribution of BDs extends down to the breakup period, where centrifugal and gravitational forces are balanced. Combining our BD periods with literature data, we found that the lower period limit for substellar objects lies between 2\\,h and 4\\,h, more or less independent of age. Contrary to stars, these fast rotating BDs seem to evolve at constant rotation period from ages of 3\\,Myr to 1\\,Gyr, in spite of the contraction process. Thus, they should experience strong rotational braking.\\thanks{Table 2 is only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/} ",
        "introduction": "\\label{intro}  Variability studies are a key tool for investigating the physics of stars. From simple photometric monitoring campaigns alone, it is possible to unveil important properties of stars. One example is the rotation period, which in many cases can be obtained from photometric light curves, if the objects exhibit asymmetrically distributed surface features, e.g. magnetically induced spots (e.g., Bouvier \\& Bertout \\cite{bb89}, Bouvier et al. \\cite{bcf93}, \\cite{bck95}, Herbst et al. \\cite{hrw00}, \\cite{hbm02}). The amplitude of the light curve then contains information about the spots, and thus about the magnetic activity of the targets (e.g., Krishnamurthi et al. \\cite{ktp98}). Other signs of activity can be seen in the light curves as well, in particular rapid brightness eruptions like flares (Stepanov et al. \\cite{sfk95}). Additionally, accretion processes manifest themselves in strong variability, because they often produce hot spots where matter flows from the accretion disk onto the star's surface. These hot spots are again a source of periodic variability, and spot instabilities as well as accretion rate variations can additionally induce irregular photometric variability (Fern\\'andez \\& Eiroa \\cite{fe96}, Herbst et al. \\cite{hmw00}).  This broad output motivated extended photometric monitoring studies of young stars. The advent of wide-field CCD detectors increases the efficiency of such projects enormously, since they allow observers to monitor a large number of objects simultaneously. Many of these studies focus on open clusters, because they deliver coeval target samples for which good estimates for age and distance are available. That way, one can evaluate the evolution and dependence on mass of stellar properties like rotation and activity.  Most monitoring studies in young open clusters, however, are concentrated on solar-mass stars with masses $>0.5\\,M_{\\odot}$. The major outcome of this work is a huge database of more than 1500 rotation periods, mostly for T Tauri stars with ages less than 10\\,Myr (see the recent reviews by Stassun \\& Terndrup \\cite{st03} and Mathieu \\cite{m03}). These periods deliver the crucial constraints for any model of rotational evolution. The state-of-the-art description of the rotational behaviour contains a) magnetic interaction between star and disk, b) angular momentum loss through stellar winds, and c) structural effects caused by contraction and internal angular momentum transport (see, e.g., Barnes \\& Sofia \\cite{bs96}, Krishnamurthi et al. \\cite{kpb97}, Bouvier et al. \\cite{bfa97}).  In the last decade, the low-mass end of the known population of open clusters was shifted well down into the substellar (or even into the planetary) mass regime (see B\\'ejar et al. \\cite{bzr99}, Zapatero Osorio et al. \\cite{zbm00}, Lucas \\& Roche \\cite{lr00}, Muench et al. \\cite{mll02}, \\cite{mll03}). This survey work delivered large samples of very low mass (VLM) objects, an object class herewith loosely defined as 'objects with masses below $0.4\\,M_{\\odot}$', including VLM stars and brown dwarfs. After the detection of these objects, the next logical step is to explore their properties, e.g. rotation and activity. From the extended work of Herbst et al. (\\cite{hbm01}, \\cite{hbm02}) in the Trapezium cluster and Lamm (\\cite{l03}) in NGC2264, first large samples of rotation periods for very young VLM stars (ages around 1\\,Myr) have become available. These samples were complemented by three periods for brown dwarfs in the similarly old Chamaeleon I star-forming region (Joergens et al. \\cite{jfc03}). One important result of these studies is that the mean rotation period decreases with decreasing mass. The total range of periods, however, is very similar to that of more massive stars; they reach from several hours up to two weeks. The fast rotation of young VLM objects has been interpreted mainly as a consequence of imperfect disk-locking (Lamm \\cite{l03}, Lamm et al. \\cite{lmb04}).  For more evolved VLM objects, however, the rotation period database is very sparse. As of the end of 2003, six periods have been published for VLM objects in open clusters with ages between 3\\,Myr and 125\\,Myr (Mart\\'{\\i}n \\& Zapatero Osorio \\cite{mz97}, Terndrup et al. \\cite{tkp99}, Bailer-Jones \\& Mundt \\cite{bm01}, Zapatero Osorio et al. \\cite{zcb03}), complemented by a few periods for ultracool dwarfs in the field (e.g., Bailer-Jones \\& Mundt \\cite{bm01}, Clarke et al. \\cite{ctc02}). All these periods are shorter than one day, and thus give tentative evidence for a lack of slow rotators among VLM objects, which is confirmed by rotational velocity studies (e.g., Terndrup et al. \\cite{tsp00}).  The described lack of rotation periods for VLM objects motivated a long-term project to study their rotational evolution. In the first two papers of this project, we published 23 rotation periods for objects in the $\\sigma$\\,Ori cluster (Scholz \\& Eisl{\\\"o}ffel \\cite{se04a}, hereafter SE1) and nine for VLM Pleiades members (Scholz \\& Eisl{\\\"o}ffel \\cite{se04b}, hereafter SE2). In this paper, we report a variability study of VLM objects near the star $\\epsilon$\\,Ori. This region (and also the $\\sigma$\\,Ori cluster) belongs to the OB1b association of the Orion star forming complex, which harbours a large population of young stars and brown dwarfs, as recently shown by Sherry (\\cite{s03}). The age of the young objects near $\\epsilon$\\,Ori lies between 2 and 10\\,Myr (see Wolk \\cite{w96}), thus the objects are on average probably somewhat older than those in the $\\sigma$\\,Ori cluster (age 3\\,Myr, Zapatero Osorio \\cite{zbp02}). The object density around $\\epsilon$\\,Ori is similarly high as in the $\\sigma$\\,Ori cluster. Therefore, we decided to use a field near $\\epsilon$\\,Ori as target for a monitoring campaign (see Fig. \\ref{field}). The particular aim of the $\\epsilon$\\,Ori project was to enlarge the rotation period database for very young brown dwarfs, and to improve the statistical significance of our previous $\\sigma$\\,Ori results.  \\begin{figure}[h] \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90,width=6.5cm]{1932fig1.ps}} \\caption{Target field of our monitoring campaign. The borders of our field of view are shown as solid lines. All VLM candidates in this region, selected by multi-filter photometry, are shown with large dots. The bright star $\\epsilon$\\,Ori is located close to the lower left corner of the field.} \\label{field} \\end{figure}  The structure of this paper is as follows. We first describe the monitoring campaign and the reduction of the images (Sect. \\ref{obs}), the photometry process and the calibration (Sect. \\ref{photo}). The targets of the variability study were selected with colour-magnitude diagrams, whose construction and evaluation is explained in Sect. \\ref{tar}. The subsequent section concentrates on the time series analysis with the main focus on period search (Sect. \\ref{tsa}). Then, we discuss the origin of the variability and show that the light curve variations are caused by cool, magnetically induced spots, hot spots from accretion flow, and flare activity (Sect. \\ref{ori}). Finally, we concentrate on the periods and discuss the mass-period relationship as well as the rotational evolution of VLM objects (Sect. \\ref{rot}). We give our conclusions in Sect. \\ref{conc}. ",
        "conclusions": "\\label{conc}  We present the results of an extended variability study of young VLM objects near the bright star $\\epsilon$\\,Ori. This region harbours a rich population of young stars and brown dwarfs (called $\\epsilon$\\,Ori cluster), which probably belongs to the Ori\\,OB1b association. The ages of these objects are probably between 2 and 10\\,Myr. In a field of 0.3 sq. deg. we identified 143 young VLM objects by means of colour-magnitude diagrams. For these objects, the photometry in five bands (RIJHK) is consistent with membership of the $\\epsilon$\\,Ori cluster. From simulated star counts for this Galactic direction, we estimate a contamination rate of 16\\% for this sample. We cover a mass range from 0.02 to 0.66$\\,M_{\\odot}$, where most of our candidates have masses below $0.4\\,M_{\\odot}$.  These objects were observed in a photometric time series with the ESO/MPG wide field imager at the 2.2-m telescope on La Silla. We covered four complete consequent nights, and obtained at least 29 I-band images per night for a total of 129 images. The instrumental magnitudes from these time series images were calibrated relative to non-variable reference stars in the field. We reach a mean precision of 5\\,mmag for the brightest targets. A time series analysis procedure was carried out for all candidate light curves. The main focus was on the period search, which was based on periodogram techniques and included several independent reliability checks.  A large fraction of the candidates shows clear signs of variability. We identified three types of variability in our targets:  a) For one VLM star, a large flare event was detected, with a sudden I-band brightness increase of 0.3\\,mag, and a subsequent exponential decline over several hours. From this finding, we estimate an I-band flare rate of VLM objects of about $2 \\cdot 10^{-4}$ per hour. Since the flare intensity usually decreases drastically toward red wavelengths, our single flare event should have been extraordinarily strong in the blue wavelength range. Observations of flares on VLM objects are rare, but the very few detected flares, including the event in $\\epsilon$\\,Ori, are quite strong. Whether this is a characteristic feature of these objects or just a result of poor statistics has to be verified.  b) Five objects, including two brown dwarf candidates, show strong and partly irregular variations, with amplitudes of up to 1\\,mag. For three of these candidates, we nonetheless detected a significant period (see c)), but clearly strong irregular brightness modulations are superimposed on their periodicities. Because this photometric behaviour is comparable to that of classical T Tauri stars, we interpret the variability of these highly variable targets as a consequence of ongoing accretion processes. Thus, this study again provides independent confirmation of the existence of accretion disks around young brown dwarfs and VLM stars. Since the fraction of highly variable objects is only 4\\%, we conclude that most VLM objects lose their disks on timescales of a few Myr.  c) For 30 targets, including nine brown dwarf candidates, we detected highly significant periods, which we interpret as the rotation periods. The periods range from 4\\,h up to 100\\,h, while our analysis is sensitive for periods from 0.2\\,h up to 110\\,h. With the exception of the three highly variable objects discussed under item b), all these periodic light curves have amplitudes below 0.15\\,mag. The most reasonable explanation for their variability is the existence of cool magnetically induced surface spots co-rotating with the objects. The amplitudes are smaller than in similar studies for solar-mass stars. This could be a consequence of a change in the dynamo in the VLM regime or a decrease of the relative spotted area induced by the high resistivities in cool atmospheres.  We interpreted the 30 periods as rotation periods of our objects. This is the largest sample of VLM objects with known rotation periods for ages $>2$\\,Myr. The mass-period relationship was analysed in comparison with the data sets in $\\sigma$\\,Ori (3\\,Myr, SE1), NGC2264 (2\\,Myr, Lamm et al. \\cite{lmb04}), and the Trapezium cluster (1\\,Myr, Herbst et al. \\cite{hbm02}). In agreement with these previous studies, we find that the median period decreases with decreasing mass. In $\\epsilon$\\,Ori, the slope of this relationship is by a factor of 1.3 lower than in the younger Trapezium cluster. In case of angular momentum conservation, we would expect a rotational acceleration by a factor between 1.3 and 7.3, where the exact value depends on the age of the $\\epsilon$\\,Ori targets. Thus, either most of the $\\epsilon$\\,Ori objects have an age of about 2-3\\,Myr, or there is still significant rotational braking, e.g. by magnetic interaction with the accretion disk ('disk-locking').  Our period sample in combination with literature data allows us to investigate the rotational evolution of brown dwarfs. BDs with ages $>$1\\,Myr are very fast rotators with average rotation periods below one day. According to the available data, the lower period limit lies between 2 and 4\\,h and is more or less independent of age. Thus, at ages between 3\\,Myr and 1\\,Gyr a fraction of the BDs rotate with very short but more or less constant rotation periods, in clear contrast to the period evolution of stars. Taking into account the hydrostatical contraction, this implies that these fast rotators experience strong rotational braking, which cannot be done only by disk-locking or stellar winds. For the young objects in $\\epsilon$\\,Ori the lower limit is very close to the breakup period, which might have an influence on their rotational evolution."
    },
    "0410/astro-ph0410146_arXiv.txt": {
        "abstract": "We report on a study of the thermal dust emission of the circumstellar envelopes of a sample of Class\\,0 sources. The physical structure (geometry, radial intensity profile, spatial temperature and spectral energy distribution) and properties (mass, size, bolometric luminosity ($L_{\\rm bol}$) and temperature ($T_{\\rm bol}$), and age) of Class\\,0 sources are derived here in an evolutionary context. This is done by combining SCUBA imaging at 450 and 850~$\\mu$m of the thermal dust emission of envelopes of Class\\,0 sources in the Perseus and Orion molecular cloud complexes with a model of the envelope, with the implementation of techniques like the blackbody fitting and radiative transfer calculations of dusty envelopes, and with the Smith evolutionary model for protostars. The modelling results obtained here confirm the validity of a simple spherical symmetric model envelope, and the assumptions about density and dust distributions following the standard envelope model. The spherically model reproduces reasonably well the observed SEDs and the radial profiles of the sources. The implications of the derived properties for protostellar evolution are illustrated by analysis of the $L_{\\rm bol}$, the $T_{\\rm bol}$, and the power-law index $p$ of the density distribution for a sample of Class\\,0 sources. ",
        "introduction": "The earliest evolutionary phase of star formation, in which a deeply embedded protostar is known to exist, is the so-called Class\\,0 stage (Andr\\'e et al.~1993). Since protostars in this phase (Class\\,0 sources) are highly obscured by extended dusty envelopes, they can be observed mainly in the far-infrared to millimetre wavelength range. Specifically, a Class\\,0 source consists of a central protostellar object surrounded by an infalling dusty envelope and a flattened accretion disk. Many (presumably all) Class\\,0 sources are associated with strong molecular flows, usually more energetic than those from Class\\,1 sources (Davis \\& Eisl\\\"offel~1995).  Since Class\\,0 sources will determine most of the early evolutionary history of the protostars as they evolve toward main-sequence stars, it is crucial to determine the physical conditions within the envelope of Class\\,0 objects. However, because of detection difficulty, constraining physical properties of Class\\,0 sources is a challenging task. Nevertheless, continuum submillimetre observations of these sources let us to detect dust emission of the massive circumstellar envelopes, and provide a powerful tool for constraining the distribution of matter (Adams~1991).  Independent methods of detection of Class\\,0 sources, and techniques and useful ways to help to the interpretation of properties of protostars have been developed and used during the last years (e.g. comparation of submillimetre with near infrared surveys of a star forming region (Stanke et al.~2000), interpretation from the dust emission (Shirley et al.~2000), blackbody fitting (Chini et al.~2001), radiative transfer codes (Wolf et al.~1999), evolutionary models for protostars (Smith~2002)). In this work, we combine them and develop an approach in order to constrain the structure (geometry, radial intensity profile, spatial temperature, and spectral energy distributions (SEDs)), and physical properties (mass, size, bolometric luminosity ($L_{\\rm bol}$) and temperature ($T_{\\rm bol}$), and age) of Class\\,0 sources. Specifically, we combine the information from the dust emission with a physical model of the extended envelope, the implementation of techniques like the blackbody fitting and the 1D continuum self-consistent MC3D radiative transfer code (by Wolf et al.~1999), and the Smith evolutionary model for protostars (by Smith~2002).  To describe and study stellar evolution, the Hertzsprung -Russell Diagram (HRD) has become the main way. However, to study protostellar evolution, the $T_{\\rm bol}$--$L_{\\rm bol}$ diagram (BLTD, bolometric luminosity-temperature diagram) was proposed as direct analogue to the HRD (Myers \\& Ladd~1993; Myers et al.~1998). There, Class\\,0 sources are characterised by the lowest $T_{\\rm bol}$ values of all protostars. Although the BLTD provides a useful way to study the evolution of protostars for all evolutionary phases of relevance for star formation (Myers et al.~1998), nevertheless, the Smith scheme predicts the evolution of $T_{\\rm bol}$ and $L_{\\rm bol}$ as an envelope of initial mass $M_{0}$ and temperature $T$ dissipates on a timescale $t$ and forms a star. ",
        "conclusions": "Detailed physical interpretation of the internal structure of nine observed Class\\,0 sources was carried out by intensive computer modelling. A simple spherically symmetric model envelope, and assumptions about density and dust distributions following the standard envelope model reproduced reasonably well the observed SEDs and the radial profiles of the sources. $T$($r$) $\\propto r^{-0.4}$ is a good approximation for the sample. The radial temperature distribution as function of distance, however, departed significantly from the optically-thin assumption, an observational derivation of a single-power law $q$ of 0.4 for radii $<$ 10~AU. Modelling results indicated a density profile well described by a power-law between $p$=1.5 and 2, which is expected by  \\begin{table*}[tH!] \\caption{Best envelope fit results for the embedded sources. The temperature of the central object T$_{*}$ is set to 3500~K. $\\chi^{2}$ quantifies the agreement between model and data. The age is estimated according to the model of Smith (2000) and given in 10$^3$\\,yrs.} \\label{tab:table} \\begin{center} \\leavevmode \\footnotesize \\begin{tabular}[h]{lcccccccc} \\hline \\\\[-5pt] Object & $T_{\\rm bol}$ [K] & $L_*$ [L$_\\odot$] &  $R_{\\rm sub}$~[AU] & $p$ & $R_{\\rm out}$~[AU] & age [$\\times$ 10$^3$ yr]& $M_{\\rm env}$\\,[M$_{\\odot}$] & $\\chi^2$\\\\[+5pt] \\hline \\\\[-5pt] \\object{L\\,1448\\,NW}              &27    &~6 &5        &2.0            &~4000  &13.8 &~2.8  &~6.4 \\\\ \\object{L\\,1448\\,C}               &40    &11 &5        &1.6            &~3000  &19.5 &~1.7  &~2.6 \\\\ \\object{RNO\\,15\\,FIR}             &45    &~8 &3        &1.6            &~3500  &20.7 &~1.0  &~1.8 \\\\ \\object{NGC\\,1333\\,IRAS\\,1}       &52    &13 &3        &1.5            &~4500  &21.9 &~1.5  &~8.6 \\\\ \\object{NGC\\,1333\\,IRAS\\,2}       &48    &68 &3        &1.5            &~7000  &26.1 &~3.0  &~0.1 \\\\ \\object{HH\\,211-MM}               &30    &~5 &3        &1.9            &~6000  &13.8 &~4.1  &~4.5 \\\\ \\object{L\\,1634}                  &42    &19 &3        &1.6            &~6000  &21.9 &~2.8  &~1.1 \\\\ \\object{L\\,1641\\,N}               &45    &85 &3        &1.5            &~8500  &25.2 &~7.0  &~1.0 \\\\ L\\,1641\\,SMS\\,III        &50    &80 &3        &1.5            &10000  &27.6 &~6.3  &~0.6 \\\\ \\hline \\\\ \\end{tabular} \\end{center} \\end{table*}  \\noindent all of the collapse models and numerical studies. A density structure of $\\rho$($r$) $\\propto r^{-2}$ at younger ages, evolving to $\\rho$($r$) $ \\propto r^{-3/2}$ at later times was found.  SED and radial profile fits constrained physical parameters of the sources such as envelope masses, density distributions, sizes and sublimation radii. Nevertheless, mid-infrared observations, and the inclusion of outflow and disk, and other geometries in the model are desirable to investigate in order to decrease the differences between the model and the data, or prove that data and model are different under several configurations."
    },
    "0410/astro-ph0410693_arXiv.txt": {
        "abstract": "Radio emission from cosmic ray air showers has the potential to become an additional, cost-effective observing technique for cosmic ray research, being largely complementary to the well-established particle detector and air fluorescence techniques. We present Monte Carlo simulatios of radio emission from extensive air showers in the scheme of coherent geosynchrotron radiation from electron-positron pairs gyrating in the earth's magnetic field. Preliminary results of our simulations are the predicted frequency, primary particle energy, shower zenith angle, shower azimuth angle and polarization dependence of the radio emission. These properties can be directly related to data measured by LOPES and other experiments. ",
        "introduction": "Accompanying the experimental efforts of the LOPES project\\cite{Horneffer2004}, we are modeling the radio emission from cosmic ray air showers in the scheme of coherent geosynchrotron radiation\\cite{FalckeGorham2003}, thereby building a theoretical foundation for the interpretation of experimental data. As a follow-up to our analytic analysis of the emission properties and the underlying coherence effects\\cite{HuegeFalcke2003a}, we have carried out sophisticated Monte Carlo simulations based on analytic parametrizations of the air shower properties to achieve a high-precision modeling of the emission including its polarization characteristics\\cite{HuegeFalcke2004a}. Here, we present some of our preliminary results. ",
        "conclusions": "We have developed a powerful Monte Carlo code for the simulation of radio emission from cosmic ray air showers. Our predicted dependence of the radio emission on important air shower parameters provides a necessary prerequisite for cosmic ray research with radio techniques. In particular, the predicted polarization characteristics allow a direct verification of the geomagnetic origin of the radiation with LOPES. A detailed analysis of our results is in preparation."
    },
    "0410/astro-ph0410370_arXiv.txt": {
        "abstract": "{ We use the {\\tt PHOENIX} code package to model the X--ray spectrum of Nova V4743 Sagittarii observed with the LETGS onboard the Chandra satellite on March 2003. Our atmosphere models are 1D spherical, expanding, line blanketed, and in full NLTE. To analyze nova atmospheres and related systems with an underlying nuclear burning envelope at X--ray wavelengths, it was necessary to update the code with new microphysics, as discussed in this paper. We demonstrate that the X--ray emission is dominated by thermal bremsstrahlung and that the hard X--rays are dominated by Fe and N absorption. The best fit to the observation is provided at a temperature of T$_{\\rm eff}$ = $5.8 \\times 10^5$\\,K, with L$_{\\rm bol}$ = 50 000 L$_\\odot$. The models are calculated for solar abundances. It is shown that the models can be used to determine abundances in the nova ejecta.} ",
        "introduction": "Classical novae (CNe) are the third most violent of the stellar explosions that can occur in a galaxy (after $\\gamma$--ray bursts and supernovae). About 5 to 10 Galactic CNe are discovered per year in the Solar neighborhood. CNe participate in the cycle of Galactic chemical evolution, in which they eject grains and metal--enriched gas as a source of heavy elements for the interstellar medium \\cite[supplementing those of supernovae, AGB stars, PNe, and WR stars:][]{gehrz98}. In addition, CNe are related to the Super Soft Binary X--ray Sources (SSS) which are probably progenitors of SN Ia explosions \\cite[]{kahabka97} and which are thought to be the result of thermonuclear explosions in the cores of mass--accreting white dwarfs in close binary systems \\cite[]{hillebrandt00,leibundgut01}. The association of CNe and SSS is important both because of the cosmological implications of SN Ia (currently the best standard candles for determining cosmological distances) and because SN Ia are thought to be responsible for the abundance of iron in the solar system.\\par CN explosions occur in a  cataclysmic variable (CV) binary star system in which a Roche lobe filling secondary supplies hydrogen--rich material that is accreted onto a white dwarf (WD) primary. Theoretical studies show that the accreted layer grows until it is sufficiently hot and dense to initiate a thermonuclear runaway (TNR). The evolution of the TNR depends upon the mass and luminosity of the WD, the rate of mass accretion, the  composition of the accreted material, and the chemical composition in the reacting layers \\cite[]{starrfield89,gehrz98}. Observations imply and theory demands that core material be dredged--up into the accreted material and the gases be thoroughly mixed before being ejected into space \\cite[]{starrfield98,gehrz98}. {\\it Thus, the chemical composition of the ejected material reflects a TNR processed mixture of WD core plus accreted material.}  Abundance studies of the ejecta carried out with the IUE satellite have established that the underlying WDs are either carbon--oxygen (CO) or oxygen, neon, and magnesium (ONeMg) WDs \\cite[]{starrfield98,gehrz98}. CNe are expected to be the major source of $^{15}$N and $^{17}$O in the Galaxy and to contribute to the abundances of other isotopes in this atomic mass range.\\par Previous X--ray studies of CNe in outburst show {\\it that no other wavelength region provides unambiguous information on the evolution and characteristics of the underlying WD} \\cite[]{drake03,krautter96,krautter02,ness03}. For example, the ROSAT studies discovered a ``soft'' component in CNe and placed them at, or near, the bright end of the class of SSS. The ``soft'' component has been identified as emission from the hot nuclear burning photosphere of the WD and theoretical studies show that a determination of the temperature and lifetime is a measure of the mass of the WD. Stellar atmosphere analyses of the photosphere can provide abundances of the material remaining on the WD. CNe at optical maximum are luminous, L$\\ge$ $10^5$L$_\\odot$ and evolve to X--ray maximum where T$_{\\rm eff}$ ranges from $3\\times 10^5$K to 6$\\times 10^5$K or higher \\cite[]{krautter96,balman98,orio02,orio03}.\\par X--ray emission during the CN outburst can be divided into three phases \\cite[]{krautter02b}. First is the early fireball phase, in which the hot atmosphere of the white dwarf is expanding and cooling adiabatically.  The peak temperature of the white dwarf is predicted to exceed $10^6$K \\cite[depending on the mass of the white dwarf:][]{starrfield96} so, if detected in this phase, it would be an extremely hot source. Nevertheless, the expanding shell cools rapidly and becomes more opaque in X--rays in a few hours. Therefore, for the first few hours of the outburst, after the shell has cooled sufficiently, the nova is probably undetectable by an X--ray satellite.  However, after a few days, the expanding shell starts to become ionized and the copious X--rays from the inner layers partially penetrate the shell and X--ray observations of novae during this phase show emission lines from the expanding gas \\cite[]{mukai01,orio02,orio03}. This is the second phase.\\par The third is the constant bolometric luminosity phase when the system again becomes observable as a ``soft'' source. This phase occurs because both observational and theoretical studies of CNe show that not the complete envelope is ejected during the explosive phase of the outburst. Some fraction remains on the WD and rapidly returns to hydrostatic equilibrium. This material provides sufficient fuel so that the enlarged WD remains hot and luminous for months to years \\cite[]{krautter96}. It is predicted that the duration of this phase is an inverse function of the mass of the WD \\cite[]{krautter96} so that a determination of how long a nova is ``on'' in X--rays can provide an estimate of the WD mass. Moreover, since the burning WD is hot, with nuclear burning ongoing in its surface layers, its characteristics closely resemble those of the SSS.\\par In this paper we report on the analysis of an observation of nova V4743 Sgr with the High Resolution Camera and the Low Energy Transmission Grating Spectrograph (HRC--S + LETGS) on the CHANDRA X--ray satellite \\cite[]{ness03}. Because the structure of the rekindled WD in V4743 Sgr consists of a nuclear burning envelope with an very hot stellar atmosphere on top, we have updated the microphysics and numerics in {\\tt PHOENIX} to analyze nova atmospheres at X--ray wavelengths. We describe these improvements in section \\ref{addphys}. {\\tt PHOENIX} is well suited to analyze V4743 Sgr, or any SSS, because it has been developed specifically to treat radiation transport in an expanding, optically thick atmosphere at any temperature. We show that we can now treat the formation and time evolution of the X--ray radiation in a CN in outburst and that it is possible to determine the chemical composition of the CN ejecta. We describe the X--ray observations of V4743 Sgr in section \\ref{obs}, and end with a Summary and Discussion (sections \\ref{res} and \\ref{concl}). ",
        "conclusions": "}  We have modified the {\\tt PHOENIX} code to calculate model atmospheres and synthetic spectra for the late X--ray phase of classical novae. The agreement between the model and the observed spectrum of nova V4743 Sgr is good in the continuum for an effective temperature of $T_{\\rm eff} = 5.8 \\times 10^5$\\,K if the 'new bremsstrahlung' (see section \\ref{addphys}) is used which is physically more reasonable. So far, we only have computed model atmospheres with solar abundances for computing time reasons and the fit to the spectral lines needs to be improved by models with non--solar abundances. The hydrogen column density of $n_h = 4.0 \\times 10^{21}$\\,cm$^{-2}$ providing the best fit is high compared to $n_h = 1.41 \\times 10^{21}$\\,cm$^{-2}$ from \\cite{dickey90}, because of additional absorption by the circumstellar shell.\\par X--ray continuum emission is dominated by thermal bremsstrahlung and its inverse processes, while bound--free emission is much weaker in the nova atmosphere. In the hard spectral range of $\\lambda \\lesssim 29$\\,{\\AA} (E $\\gtrsim 428$\\,eV) iron and nitrogen absorption are dominant. As expected, nova ejecta show strong departures from LTE. They are very extended and several of the highest ionization stages of all elements are simultaneously present.\\par In future work additional physical processes absorbing and emitting X--rays will be implemented in the code. For example there are data for dielectronic recombination and autoionization, and there are better data for free--bound coefficients in the APED and CHIANTI4 databases.\\par Models with non--solar abundances will be calculated to obtain a better fit for the line strengths and the continuum absorption and to analyze the abundances of the nova ejecta. To determine reliable chemical abundances more elements have to be included in the NLTE calculations.\\par The data quality of the atomic databases as well as differences between the CHIANTI4 and APED database will be checked, e.~g. by examining the line profiles. By comparing model and observation, it has to be determined which data of the individual ions reproduce the spectral lines as accurately as possible. This is necessary for an accurate abundance analysis.\\par The evolution of the X--rays should be examined, by modeling the spectra of other observations of novae with the CHANDRA satellite. It is important to test if it is possible to model the spectrum of novae in the emission line phase, when nuclear burning has switched off and the expanding shell is optically thin, like the spectrum of V382 Vel \\cite[]{burwitz02}. For this work more observations with CHANDRA are necessary.\\par"
    },
    "0410/astro-ph0410236_arXiv.txt": {
        "abstract": "Deep $J, H,$ and $K'$ images, recorded with the ALTAIR adaptive optics system and NIRI imager on Gemini North, are used to probe the stellar content of the disk and bulge of the Local Group galaxy M31. With FWHM near 0.08 arcsec in $K$, these are the highest angular resolution near-infrared images yet obtained of this galaxy. One field samples the outer disk of M31 at a galactocentric distance of roughly 62 arcmin along the major axis. The mean metallicity in this field is close to that of the metal-rich globular cluster NGC 6528, and no stars with [Fe/H] $< -0.7$ are detected. Another field, located on the major axis 9 arcmin from the galaxy center, contains a roughly equal mix of disk and bulge stars. The RGB in this field is redder than that of NGC 6528, although it is argued that reddening internal to M31 may be significant in this region of the galaxy. The remaining two fields, located at projected galactocentric distances of 2 and 4 arcmin, are dominated by bulge stars. The RGB-tip occurs between $K = 17.0$ and 17.2, and the color of the RGB in the field closest to the center of M31 is consistent with that of NGC 6528. After accounting for random photometric errors, the upper RGB in each field has a width on the $(K, J-K)$ CMDs that is consistent with a $\\pm 0.5$ dex dispersion in [Fe/H], in rough agreement with what is seen in other disk and spheroid fields in M31. The number of bright AGB and RGB stars also scales with the $r-$band surface brightness in all four fields. Thus, we conclude that the stellar content does not change markedly from field-to-field, and that the photometric properties of the brightest AGB stars in the two innermost fields are not affected significantly by crowding. The brightest star has M$_K = -8.6$ and M$_{bol} = -5.2$, although this may not be a reliable measure of the AGB-tip brightness due to photometric variability.  A population of very bright red stars, which we identify as C stars, are seen in the three fields that are closest to the center of M31. The spatial distribution of these objects suggest that they are well mixed throughout this part of M31, and so likely did not form in a compact region near the galactic nucleus, but more probably formed in the inner disk. We speculate that these C stars may be the most luminous members of the intermediate age population that has been detected previously in studies of the integrated spectrum of the central regions of M31. ",
        "introduction": "In recent years considerable effort has been dedicated to determining when the morphological characteristics of galaxies were set in place. In the case of spiral galaxies a major issue is the relative timescale of bulge and disk formation (e.g. Bouwens, Cayon, \\& Silk 1999). One class of model argues that bulges form early-on, prior to the creation of stable disks, as the result of mergers or a monolithic collapse. Indeed, cosmological models that assume a cold dark matter (CDM) dominated Universe predict that large galaxies formed through the accretion of small gas-rich proto-galactic fragments. These fragments were likely short-lived, as they were consumed in major mergers that altered morphologies through kinematic heating and spurred large-scale episodes of star formation (e.g. Somerville, Primack, \\& Faber 2001). Given the chaotic state of the early Universe, and the difficulty in maintaining stable disks in such environments (e.g. Weil, Eke, \\& Efstathiou 1998), it might be anticipated that bulges were the first long-lived structures to form in the systems that would eventually become spiral galaxies. This is consistent with the relative colors of bulges and disks in the Hubble Deep Field (Abraham et al. 1999), and with the age of the Galactic bulge (e.g. Feltzing \\& Gilmore 2000, Ortolani et al. 1995).  At the other extreme are models in which disks formed first, and bulges are the result of secular processes involving disk material. The motion of gas in the disk can lead to the formation of spheroidal structures, and bars, which may form when gas is channeled towards the central regions of a galaxy, are key elements in this process. The bar transports angular momentum from the central regions of the galaxy to larger radii, and stars that formed in the bar will be scattered out of the disk to form a spheroidal structure when the bar buckles (e.g. Friedli \\& Benz 1995). Stars that form in the inner regions of a galaxy in the absence of a bar may also be scattered into the surrounding bulge by interactions with molecular clouds (Kim \\& Morris 2001). Consequently, the inner bulges of systems that are experiencing, or did experience, nuclear star formation might contain a diffuse population of intermediate age stars that did not form {\\it in situ}.  Evidence that material originally in the disk contributes to the formation of bulges comes from the structural properties of the spheroids in late-type galaxies. Rather than following an R$^{1/4}$ profile that is the expected signature of violent relaxation, the central spheroids of late-type spiral galaxies follow Sersic profiles with exponents that are indicative of a more leisurely formation process (e.g. Balcells et al. 2003; Courteau, de Jong, \\& Broeils 1996; Andredakis, Peletier, \\& Balcells 1995). It might also be anticipated that the stellar contents of bulges will show a morphological dependence if they are built up by secular processes, in the sense that the luminosity-weighted ages of bulges will become younger as one moves to later morphological types. Nevertheless, some exponential bulges have colors that are similar to those of bulges that follow R$^{1/4}$ light profiles (Carollo et al. 2001), suggesting that secular evolution may have occured early-on, but was subsequently curtailed. The development of bulges from disk material may also affect the large scale properties of disks, as the mixing caused by bars will smooth out pre-existing abundance profiles in the disk (e.g. Friedli \\& Benz 1995).  The monolithic collapse and secular evolution models of bulge formation make certain predictions about the properties of bulges and their component stars. If bulges form early-on in violent collapse episodes then they will be made up of old stars. Metallicity gradients that are the signatures of dissipation should also be present. Finally, the chemical contents of stars formed in such a manner will show the signatures of rapid enrichment. On the other hand, bulges that are assembled solely from secular processes will be comprised of young or intermediate age stars. There may be an age gradient, in the sense that the youngest stars are located closest to the disk plane. The chemical compositions of stars in a bulge assembled in such a way will also show the signatures of a slow enrichment history.  As noted by Bouwens et al. (1999), reality likely falls somewhere between the extremes defined by the monolithic collapse and secular evolution models. Indeed, bulges almost certainly do not evolve in isolation, and there are observational indications that the evolution of bulges and disks are coupled, even long after the bulge is in place. There is a high incidence of distinct nuclear sources in the central regions of nearby bulges, which have spectrophotometric properties that differ from those of the surrounding bulge and show evidence of recent star formation (e.g. Boker et al. 1999, 2002; Carollo et al. 2002). Reservoirs of cool gas must be present near the centers of these galaxies, and the surrounding disks are a likely source of this material. This interplay between disks and bulges has been on-going for some time, as central blue knots are seen in many disk galaxies at intermediate redshifts (Ellis, Abraham, \\& Dickinson 2001).  The observational challenges inherent to studies of very distant systems can make it difficult to use them to distinguish between the various processes that affect bulge evolution (e.g. Bouwens et al. 1999). However, the interplay between disks and bulges will have observational consequences in nearby spiral galaxies, and so these systems are key laboratories for testing models of bulge evolution. Studies of the resolved stellar contents in the bulges of nearby galaxies can be used to search for the young and intermediate age stars that are the expected signatures of secular processes. Moreover, if stars on the RGB are resolved then it may be possible to search for the abundance gradients that are the expected signatures of a dissipational collapse.  As the closest external spiral galaxy, M31 is an obvious target for detailed study. There are hints that, like many other nearby systems, the central regions of M31 may have experienced star formation in the not too distant past. The colors of the nucleus differ from those of the surrounding bulge (King, Stanford, \\& Crane 1995; Davidge et al. 1997; Brown et al. 1998; Lauer et al. 1998), suggesting a distinct stellar content, such as might remain after the termination of star formation during the last few Gyr. There are also spectroscopic signatures of a central young and/or intermediate age population (e.g. Bica, Alloin, \\& Schmidt 1990; Davidge 1997; Sil'chenko, Burenkov, \\& Vlasyuk 1998), although a population of resolved intermediate-aged stars has yet to be detected in the inner regions of M31. In fact, the brightest red stars are well-mixed throughout the inner bulge, and so likely formed with the main body of the bulge (Davidge 2001). In addition, the brightnesses and densities of luminous AGB stars in the bulge of M31 are similar to those in systems that are as disparate as M32 and NGC 5128, further supporting the notion that they come from an old population (Davidge 2002). The red giant branch of the M31 bulge has a color that is consistent with a near-solar metallicity (Jablonka et al. 1999; Stephens et al. 2001), although the RGB-tip may be depressed with respect to larger galactocentric distances (Jablonka et al. 1999). This latter result is counter to what would be expected if metallicity decreases with increasing galactocentric distance, but could be indicative of an age gradient.  Additional hints into the star forming history of the central regions of M31 come from studies of stars in the disk of the galaxy. Morrison et al. (2004) find a population of globular clusters in M31 that have disk kinematics and span a range of metallicities, suggesting that a stable disk was in place at very early epochs, so that secular interactions between the disk and bulge may have started early-on. Young and intermediate age stars are seen at moderate to large galactocentric distances (e.g. Richer, Crabtree, \\& Pritchet 1990; Davidge 1993; Brewer, Richer, \\& Crabtree 1995; Cuillandre et al. 2001; Sarajedini \\& Van Duyne 2001; Ferguson \\& Johnson 2001), and there was large scale star formation throughout much of the disk up to 1 Gyr in the past (Williams 2002). However, star formation at present appears to be restricted to the spiral arms, where the star formation rate has not changed during the past Gyr (Williams 2002). The metallicity distribution is remarkably similar over a large range of galactocentric distances, peaking near [Fe/H] $= -0.6$, with only a very modest number of stars with [Fe/H] $< -1.5$, even in fields that sample the halo (e.g. Bellazzini et al. 2003). Finally, there have been interactions between M31 and its companions (e.g. Saito \\& Iye 2000; Choi, Guhathakurta, \\& Johnston 2002; Ferguson et al. 2002), which have probably influenced the star-forming history of M31, and may have triggered flows of gas into the central regions of the galaxy.  Are there spatially resolved young or intermediate age stars in the innermost regions of M31? Do studies of the photometric properties of resolved RGB stars reveal signs of a metallicity gradient in the inner spheroid, as predicted if the bulge experienced a dissipational collapse? In the present study, deep $J, H,$ and $K'$ images obtained with the ALTAIR adaptive optics (AO) system $+$ NIRI imager on Gemini North (GN) are used to probe the age and metallicity distributions in four M31 fields. The near-infrared is an important wavelength region for studies of this nature. The photometric properties of very red stars at visible wavelengths are affected by line blanketing, with the result that the upper RGB of stars with near-solar metallicities slumps over on visible wavelength CMDs. This makes it difficult to detect the most metal-rich upper RGB stars, and so the mean metallicity of the RGB may be underestimated. Line blanketing also causes the AGB to be a near horizontal sequence on CMDs constructed from data with observations shortward of $1\\mu$m (e.g. Figure 5 of Richer et al. 1990). However, the AGB forms a sequence that is closer to vertical on near-infrared CMDs (e.g. Davidge 2000a), thereby making it easier to resolve AGB stars in crowded fields and identify the AGB-tip. There is also a well-calibrated relation between the slope of the RGB and metallicity in the near-infrared (e.g. Ferraro et al. 2000; Kuchinski et al. 1995). Finally, interstellar extinction is less of a concern in the infrared, reducing complications caused by the non-uniform distribution of dust throughout the M31 disk (e.g. Cuillandre et al. 2001; Williams 2002).  The fields that were observed sample a range of stellar densities, from the outer disk, where the incidence of crowding is negligible, to a distance of only two arcmin from the galaxy center, where crowding is a significant concern (e.g. Davidge 2001; Stephens et al. 2001). Three of the fields are dominated by bulge stars, while the fourth is dominated by disk stars. With an angular resolution of 0.08 arcsec FWHM in $K'$, which corresponds roughly to 0.3 parsecs at the distance of M31, these are the highest angular resolution near-infrared images yet obtained of this galaxy.  The paper is structured as follows. Details of the observations, the data reduction procedures, and the photometric measurements are discussed in \\S 2. In \\S 3 and 4 the photometric properties of stars in the disk and bulge fields are discussed, while in \\S 5 the issue of crowding is addressed. A summary and discussion of the results follows in \\S 6. ",
        "conclusions": "Near-infrared images with angular resolutions approaching the diffraction limit of the 8 meter GN telescope have been used to investigate the photometric characteristics of bright evolved stars in the bulge and inner disk of M31. These are the highest angular resolution images yet obtained of M31 at these wavelengths, and the photometric measurements of the brightest red stars are not affected by blending, even in the field that is closest to the center of the galaxy (\\S 5). Thus, these data provide unique insights into the stellar content in the central few arcmin of M31.  \\subsection{Stellar Content Trends in the Disk and Bulge of M31}  Bellazzini et al. (2003) found a remarkable similarity in the metallicity properties of RGB stars in a number of M31 fields that (1) sample a range of galactocentric radii, and (2) are well outside of the bulge. The data discussed in the present paper extend field-to-field comparisons of RGB properties to smaller, bulge-dominated, galactocentric radii. Indeed, the galactocentric radius of Bulge 1 is an order of magnitude smaller than the field closest to the center of M31 studied by Bellazzini et al. (2003).  In \\S 3 and 4 it was demonstrated that the RGBs of Disk 1 and Bulge 1 have $J-K$ colors that are consistent with the RGB of the metal-rich globular cluster NGC 6528, while the RGBs of Disk 2 and Bulge 2 have $J-K$ colors that are redder than the NGC 6528 sequence. Before discussing this result further, it should be noted that other colors give similar results. In particular, we have compared the $(H, J-H)$ CMDs of the four fields observed with ALTAIR with fiducial globular cluster sequences from Table 6 of Valenti, Ferraro, \\& Origlia (2004). The RGBs of Disk 1 and Bulge 1 on the $(H, J-H)$ CMDs agree well with the RGB of the metal-rich globular cluster NGC 6440, which has a metallicity that is only slightly lower than that of NGC 6528, while the RGBs of Disk 2 and Bulge 2 are significantly redder.  The mean metallicities of the four fields observed with ALTAIR, as inferred from the color of the RGB, are likely reliable to no more than a few tenths of a dex.  The uncertainty in the photometric calibration is $\\pm 0.1$ magnitude, which introduces uncertainties of $\\pm 0.3$ dex in [Fe/H]. An even greater source of potential error is the patchy dust absorption that occurs throughout the central regions of M31. In \\S 4.2 it was argued that this introduces uncertainties of $0.1 - 0.2$ magnitudes in the intrinsic RGB colors, with the result that the mean metallicity in each field can not be estimated to better than a few tenths of a dex.  While the color of the RGB can be affected by uncertainties in the photometric calibration and the mean line-of-sight extinction, this is not the case for the scatter at a given brightness within a particular field. After accounting for the dispersion caused by photometric errors, as predicted by artificial star experiments, the widths of the upper RGBs of all four fields on the $(K, J-K)$ CMDs are consistent with a metallicity dispersion with a standard deviation $\\pm 0.5$ dex in [Fe/H]. This is similar to the dispersion detected by Bellazzini et al. (2003) in a large number of fields throughout the disk and halo of M31.  The similarity in stellar content between the four fields observed with ALTAIR is not restricted to metallicity and metallicity dispersion. In \\S 5.2 it was demonstrated that the $K-$band LFs of the 4 ALTAIR fields scale with $r-$band surface brightness. In old solar metallicity populations the integrated $r-$band light is dominated by main sequence stars (e.g. Buzzoni 1989). For comparison, the $K-$band LFs constructed from the ALTAIR data are dominated by stars on the upper RGB and AGB. The comparisons in Figure 15 thus show that the ratio of AGB $+$ RGB stars to main sequence stars, which is a quantity that is sensitive to age, appears not to change over a range of disk-to-bulge ratios.  In the following sub-sections the results obtained from the ALTAIR data are discussed in the context of other observations. Given that disks and bulges are structurally distinct entities, the stellar contents of the disk and bulge of M31 are discussed separately. However, the reader is reminded that there are mechanisms that may couple the stellar contents of disks and bulges (e.g. Kim \\& Morris 2001).  \\subsubsection{The Disk}  That the metallicities of RGB stars in Disk 1 and 2 differ by no more than a few tenths of a dex suggests that the metallicity gradient among RGB stars in the disk of M31 must be modest in size, and this conclusion is consistent with some, but not all, studies of the chemical content in M31. For example, Trundle et al. (2002) find that the chemical contents of bright supergiants do not show radial variations in M31, while Bellazzini et al. (2003) find that the abundance distributions of RGB stars in a number of fields are similar. In contrast, Blair, Kirshner, \\& Chevalier (1982) find radial gradients in the abundances of nitrogen and oxygen in M31 HII regions, which are consistent with radial variations in the C to M giant ratio when R$_{GC} < 20$ kpc if these are interpreted as a metallicity effect (Brewer, Richer, \\& Crabtree 1995).  Radial variations in mean age and metallicity are common in spiral galaxies (e.g. Bell \\& de Jong 2000), and gradients of this nature were likely imprinted during disk formation. However, as noted in the previous paragraph, there are some indications that mixing may have been highly effective throughout a large portion of the M31 disk, at least among stars with ages comparable to objects on the RGB. What processes could have smoothed population gradients in the disk of M31?  Radial flows of gas and stars in disks can be induced by a number of processes, including spiral density waves (e.g. Casuso \\& Berkman 2001; Sellwood \\& Binney 2002), viscosity (Silk 2001; Ferguson \\& Clarke 2001), galaxy-galaxy interactions (Barnes \\& Hernquist 1992; Mihos \\& Hernquist 1996), dynamical friction (Noguchi 1999; 2000) and the infall of angular momentum-rich gas from the halo (Ferguson \\& Clarke 2001). The outward movement of gas clouds may drive the age distribution near the disk edge to younger values (Ferguson \\& Clarke 2001). In fact, intermediate age populations at large galactocentric distances, possibly related to the disk, are seen in the Sc galaxies NGC 2403 and M33 (Davidge 2003a), as well as in the outer regions of the M31 disk (Ferguson \\& Johnson 2001).  Of the various mechanisms that might cause mixing, tidal interactions between M31 and its companions are potentially attractive given the evidence that M31 and its satellites have interacted. For example, there is a warp in the HI disk (e.g. Sawa \\& Sofue 1982), while the star forming history of NGC 205 appears to be coupled with its orbit about M31 (Davidge 2003b). Perhaps most significantly, there are well-defined tidal tails extending from some M31 companions (e.g. Ferguson et al. 2002). However, tidal interactions should also cause mixing of gas, and so this process does not explain the abundance gradients found by Blair et al. (1982).  Clearly, it will be of interest to determine if RGB stars in other disk systems, such as M81, show evidence for an abundance gradient, or are as well mixed as in M31.  \\subsubsection{The Bulge}  The bulge of M31 likely contains a radial metallicity gradient. Davidge (1997) estimated that $\\Delta$[Fe/H]/$\\Delta$log(r) $= -0.5$ from a grid of long slit spectra that mapped the central 1 arcmin of M31. This is shallower than what is seen in the outer regions of the bulge of the Milky-Way (Minniti et al. 1995; Tiede, Frogel, \\& Terndrup 1995), although Ramirez et al. (2000) find that the metallicity gradient in the Galactic bulge may flatten in the central few hundred parsecs.  If the gradient measured by Davidge (1997) is assumed to extend to larger radii then the mean metallicities of Bulge 1 and 2 will differ by 0.15 dex. This is likely a lower limit as the abundance gradient outside of the inner regions of the M31 bulge may steepen, as is the case in the Milky-Way (Ramirez et al. 2000). The Bulge 1 and 2 data may contain signs of a metallicity gradient in the M31 bulge. In \\S 4.2 it was demonstrated that the RGB-tip in Bulge 1 appears to the brighter than that in Bulge 2. While the artificial star experiments indicate that some of the stars in Bulge 1 near the RGB-tip are likely blends, the number of such objects will account for only $\\sim 7\\%$ of the total at this brightness (\\S 2.4), and this is clearly not enough to explain the difference in the RGB-tips in Figure 13. We further note that the difference in RGB-tip brightness is also not a consequence of small number statistics, as numerical simulations indicate that only 45 stars are needed on the upper 2.5 mag of the RGB to measure the RGB-tip brightness to within $\\pm 0.1$ mag (Crocker \\& Rood 1984), and the upper RGBs of Bulge 1 and 2 contain far more stars than this.  A metal-rich population in Bulge 1 that is not seen in Bulge 2 could explain the difference in RGB-tip brightnesses, as the brightness of the RGB-tip in $K$ depends on metallicity, in the sense of becoming brighter and redder with increasing metallicity. The calibration between the $K-$band brightness and metallicity among globular clusters derived by Ferraro et al. (2000) indicates that a 0.15 dex difference in metallicity would change the RGB-tip brightness by roughly 0.1 magnitude, and a shift of this nature would bring the Bulge 1 and 2 LFs in Figure 13 into better agreement. It should be emphasized that the RGB colors of Bulge 1 and 2 are not consistent with this interpretation, although this could be a consequence of reddening internal to M31 (e.g. \\S 4.2). Clearly, spectroscopic observations of individual stars will be the most secure means of measuring metallicities, and integral field spectrographs coupled to adaptive optics systems, such as ALTAIR $+$ NIFS on Gemini North, will be required to overcome crowding in the high stellar density fields near the center of M31.  \\subsection{The Bright AGB Content Near the Center of M31 and the Presence of C Stars}  One of the goals of photometric surveys of nearby spiral galaxies is to probe the star-forming histories of their pressure and rotationally supported components, as this will provide insight into when their global morphological characterics were imprinted. In the case of the Milky-Way, the relative age of the disk and bulge is still a matter of debate. Studies of the main sequence turn-off in the Galactic bulge indicate that it has an age that is similar to that of the majority of globular clusters (Feltzing \\& Gilmore 2000; Kuijken \\& Rich 2002; Zoccali et al. 2003). However, the age of the Galactic disk is less certain. Binney, Dehnen, \\& Bertelli (2000) conclude that the solar neighborhood contains stars that are at least as old as the oldest globular clusters, although studies of white dwarfs suggest an age of only 8 Gyr for the solar disk (Leggett, Ruiz, \\& Bergeron 1998, but see also Hansen 1999).  Peak AGB brightness offers one means of investigating the age of stars near the center of M31, although there are significant problems with this statistic as an age indicator. First, the three fields observed with ALTAIR near the center of M31 contain a mixture of bulge and disk stars. Lacking kinematic information for individual objects, the task of isolating pure samples of disk and bulge stars is problematic. Second, the brightness of the AGB peak depends on metallicity as well as age, so that there is an age-metallicity degeneracy when interpreting the peak brightness. A young or intermediate-age component that is moderately metal-poor may also be missed if there is an older but more metal-rich population with a brighter AGB-tip. Third, efforts to measure the peak brightness of the AGB are complicated by stellar variability. If the brightest AGB stars in the bulge of M31 are LPVs like those in the Galactic bulge then they will have amplitudes of up to $\\pm 1.0$ mag in $K$ (Glass et al. 1995). In fact, it is likely that the brightest stars imaged in the bulge of M31 at any given time are LPVs near the peak of their light curves, and this can be confirmed with multi-epoch observations, as has been done for M32 by Davidge \\& Rigaut (2004).  The brightest AGB stars in Bulge 1 and 2 have M$_K < -8.5$, and M$_{bol} = -5$, and hence are not different from what is seen in the Galactic bulge (Frogel \\& Whitford 1987). If these stars are LPVs near the peak of their light curves then they will have a mean brightness M$_K \\sim -7.5$. This is consistent with the expected AGB peak brightness in an old solar-metallicity system (Girardi et al. 2002). The peak AGB brightness in the M31 bulge is also comparable to that of the brightest AGB star in the metal-rich globular cluster NGC 6553 (Guarnieri, Renzini, \\& Ortolani 1997), which has an old age (e.g. Ortolani et al. 1995). These data thus suggest that an old population dominates in the inner regions of M31.  As noted above, peak stellar brightness does not offer an ironclad means of probing age. In fact, there are a number of sources with M$_K < -8$ in Bulge 1 and 2 and the Disk 2 field that have very red colors, and these provide additional insights into the stellar content near the center of M31 that contradicts the notion that these fields are dominated by an old stellar population. In particular, Hughes \\& Wood (1991) investigated the near-infrared photometric properties of LPVs in the LMC, and found that the vast majority of objects with $J-K > 1.6$ are C stars; of the three objects in their sample with $J-K > 1.6$ that are not C stars, two are very luminous (M$_{bol} < -5.5$) and have periods in excess of 600 days. Davidge (2003b) used $H-K$ and $J-K$ colors to identify C stars in NGC 205, and recovered C stars identified previously by Richer, Crabtree, \\& Pritchet (1984) from narrow-band photometry.  Based on the photometric study by Hughes \\& Wood (1991), we identify objects with $J-K > 1.6$ in the ALTAIR CMDs as C stars. These stars tend to have M$_K$ between --7 and --8.5. Observations indicate that C star production does not occur in systems with ages in excess of 6 Gyr, and that the maximum age for C star production drops as one moves to higher metallicities (e.g. Cole \\& Weinberg 2002). Hence, the presence of C stars indicate that there is likely an intermediate age population in Disk 2 and Bulge 1 and 2.  The identification of the very red objects in M31 as C stars has potential implications for the bulge and inner disk of the Milky-Way, as a population of very red objects is also seen in the $(K, J-K)$ CMD of Baade's Window shown in Figure 17 of Frogel \\& Whitford (1987). The main stellar sample considered by Frogel \\& Whitford (1987) consists of M giants that were selected from grism surveys, and only two of the stars identified as M giants have $J-K > 1.6$. Both of these objects have spectral types from Blanco, McCarthy, \\& Blanco (1984), are photometrically variable, and have two $J-K$ colors listed in Table 1 of Frogel \\& Whitford (1987). Star \\# 87 (spectral type M8) has $J-K = 1.59$ and 1.76, and so may or may not meet the near-infrared color criterion for C star status depending on the epoch of observation. Star \\# 250 (spectral type M7) is clearly more interesting, as it has $J-K > 2$ in both measurements; however, it is also one of the faintest stars classified by Blanco et al.  (1984), and so presumably falls near the faint limit of their data.  The remainder of the stars with $J-K > 1.6$ in the Frogel \\& Whitford sample are LPVs from the sample that was observed by Glass \\& Feast (1982). To the best of our knowledge these objects do not have spectra, presumably because they are relatively faint at visible wavelengths. In fact, the TCD shown in Figure 2 of Frogel \\& Whitford (1987) indicates that a star with $J-K = 1.6$ will have $V-K \\sim 10$. The LPVs in Figure 17 of Frogel \\& Whitford (1987) tend to have $K_0$ between 6 and 7.5 (i.e. M$_K$ between --7 and --8.5), and so will have $V_0$ between 16 and 17.5. After dimming to account for extinction, these stars will have $V$ between 17 and 18.5.  Azzopardi, Lequeux, \\& Rebeirot (1985) conducted a grens survey of Baade's Window, and found 15 C stars with $V$ between 15.5 and 17.5. Azzopardi et al. (1985) argued that these stars are too faint to be C stars associated with the Galactic bulge; however, the arguements given in the previous paragraph indicate that at least some of the objects found by Azzopardi et al. (1985) have $V$ brightnesses that are consistent with them being very red C stars at the distance of the Galactic bulge. To be sure, some of these objects may be dwarf C stars, which are difficult to identify using spectrophotometric information alone (e.g. Downes et al. 2004). The only ironclad means to confirm the nature of the C stars indentified by Azzopardi et al. (1985) is to obtain proper motions for these objects, which will allow their distances to be measured.  \\subsection{The Nature of Intermediate Age Stars Near the Center of M31}  The detection of intermediate age stars in the ALTAIR data is perhaps not a surprise. Spectroscopic studies suggest that there is a centrally-concentrated young or intermediate-age component in M31 (Bica et al. 1990; Davidge 1997; Sil'chenko et al. 1998). The geometric center of M31 has a very red $V-K$ color (Davidge et al. 1997), which may be due to a population of AGB stars that formed during the past Gyr.  Studies of the innermost regions of the Galactic bulge may provide clues about the nature of the central regions of M31. Contrary to expectations (e.g. Figer et al. 2000), star formation occurs near the center of the Galaxy. There is a cluster of young stars around SgrA* (Lebofsky, Rieke, \\& Tokunaga 1982; Krabbe et al. 1991; Davidge et al. 1997; Blum et al. 2003), as well as other young clusters within the central 100 pc (e.g. Cotera et al. 1996; Figer et al. 1999), and at even larger distances (Launhardt, Zylka, \\& Mezger 2002). The structural properties of the region near the Galactic Center (Serabyn \\& Morris 1996), the presence of shells at large R$_{GC}$ that apparently originate from the inner regions of the Galaxy (Bland-Hawthorn \\& Cohen 2003), and the analysis of the stellar content near the Galactic Center (Blum et al. 2003) together suggest that the current star forming activity near the Galactic Center is not a unique event, but is likely an on-going or episodic phenomenon with a time scale shorter than a Gyr. Nuclei similar to SgrA are also common in nearby spiral galaxies, further supporting the arguement that nuclear star formation is episodic (e.g. Davidge \\& Courteau 2002).  The timescale for the disruption of star clusters near the Galactic Center is short (e.g. Figer at al. 1999), and individual stars may be scattered out of the disk plane by interactions with giant molecular clouds (Kim \\& Morris 2001). Consequently, the stars that form along the plane of the disk in the inner regions of a spiral galaxy can diffuse away from their natal locations, and be injected into the surrounding bulge. Processes of this nature could explain the presence of a diffuse intermediate age population in the inner Galactic bulge (Wood \\& Bessell 1983; Harmon \\& Gilmore 1988; Lopez-Corredoira, Garzon, \\& Hammersley 2001; but see also ven der Veen \\& Habing 1990).  The spatial distribution of the C stars seen in Bulge 1 and 2 provide clues about their origins. If they formed in a nuclear region then it can be anticipated that they would be concentrated towards smaller radii. Such a trend is not seen in the ALTAIR data, although the uncertainties are high. The very red objects in the Stephens et al. (2003) sample, which were identified as LPVs in that study, are also uniformly distributed. Perhaps the greatest challenge to any model that tries to associate the C stars with nuclear star formation comes from the $(K, H-K)$ CMDs of the innermost regions of M31 constructed by Davidge (2001). These CMDs have relatively tight, well-defined sequences, with only a modest number of red stars near the bright end, indicating that the number density of C stars may actually drop towards very small radii in M31. Given these arguements, it appears that the C stars did not form in the nuclear regions of M31, but more likely originated over a much larger area.  Studies of late-type spiral galaxies indicate that the density of molecular star-forming material is coupled to disk properties, such as the central disk density (Boker, Lisenfeld, \\& Schinnerer 2003), rather than properties of the inner spheroid, and so it is likely that the C stars in Bulge 1 and 2 formed as part of the inner disk of M31, rather than the inner spheroid. Nevertheless, the presence of C stars in the inner regions of M31 suggests that the bulge of this galaxy may be subject to secular processes. Kim \\& Morris (2001) argue that $1 - 2$ Gyr is required for stars in the inner Galaxy to scatter and form an aspect ratio similar to that of the most luminous AGB stars in the Galactic Bulge. This is shorter than the upper age limit of C stars based on statistical studies (Cole \\& Weinberg 2002), so it is possible that C star progenitors may survive long enough to be redistributed throughout the inner bulge. A radial velocity study of the bright, red stars that occur throughout the central few arcmin of M31 will provide insight into the nature of these objects.  We close by noting that with R$_{GC} = 2$ arcmin, Bulge 1 does not overlap with the central region that has been the target of population studies based on integrated spectroscopic information. Nevertheless, the spatial distribution of the C stars detected here suggests that they are well-mixed throughout the sight line within a few arcmin of the center of M31, and so it seems reasonable to expect that some C stars will almost certainly occur at smaller R$_{GC}$ than sampled by Bulge 1. If this is the case then the C stars seen here may be the brightest members of the intermediate-age population that has been detected spectroscopically.  \\parindent=0.0cm \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccccc} \\tableline\\tableline Field Name & RA & Dec & GSC \\# & $R_{GC}$ & $\\mu_{disk}$ & $\\mu_{bulge}$ & Exposures \\\\ & E2000 & E2000 & & (arcmin) & & & (seconds) \\\\ \\tableline Disk 1 & 00:39:13.3 & +40:29:13.3 & 02788-02207 & 61.74 & 22.2 & -- & $20 \\times 40$ ($J+H$) \\\\ & & & & & & & $24 \\times 40$ ($K$) \\\\ & & & & & & & \\\\ Disk 2 & 00:43:21.7 & +41:21:55.2 & 02805-02131 & 9.09 & 20.8 & 20.7 & $24 \\times 40$ ($J+H$) \\\\ & & & & & & & $27 \\times 40$ ($K$) \\\\ & & & & & & & \\\\ Bulge 2 & 00:42:50.9 & +41:12:31.4 & 02801-02008 & 3.83 & 20.5 & 19.2 & $8 \\times 40$ ($J+H$) \\\\ & & & & & & & $10 \\times 40$ ($K$) \\\\ & & & & & & & \\\\ Bulge 1 & 00:42:33.5 & +41:15:50.1 & -- & 2.05 & 20.2 & 18.1 & $8 \\times 40$ ($J+H$) \\\\ & & & & & & & $10 \\times 40$ ($K$) \\\\ & & & & & & & \\\\ \\tableline \\end{tabular} \\end{center} \\caption{Field Information and Exposure Times} \\end{table*}  \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lccccc} \\tableline\\tableline Field name & $\\sigma_{HK}^{Measured}$ & $\\sigma_{HK}^{Artificial}$ & $\\sigma_{JK}^{Measured}$ & $\\sigma_{JK}^{Artificial}$ & $\\sigma_{JK}^{Residual}$ \\\\ \\tableline Disk 1 & $\\pm 0.14$ & $\\pm 0.01$ & $\\pm 0.10$ & $\\pm 0.04$ & $\\pm 0.09$ \\\\ Disk 2 & $\\pm 0.09$ & $\\pm 0.02$ & $\\pm 0.16$ & $\\pm 0.04$ & $\\pm 0.15$ \\\\ Bulge 2 & $\\pm 0.11$ & $\\pm 0.03$ & $\\pm 0.13$ & $\\pm 0.07$ & $\\pm 0.11$ \\\\ Bulge 1 & $\\pm 0.11$ & $\\pm 0.06$ & $\\pm 0.22$ & $\\pm 0.21$ & $\\pm 0.08$ \\\\ \\tableline \\end{tabular} \\end{center} \\caption{Measured and Predicted Dispersions for Stars with $K$ Between 17.5 and 18.5} \\end{table*}  \\clearpage"
    },
    "0410/hep-ph0410190_arXiv.txt": {
        "abstract": "We present a study of the decay of an intermediate mini black hole at the first impact of a cosmic ray particle with the atmosphere, in the context of D-brane world scenarios with TeV scale gravity and large extra dimensions. We model the decay of the black hole using the semiclassical approximation and include the corrections coming from energy loss into the bulk. Extensive simulations show that mini black hole events are characterized by essentially different multiplicities and wider lateral distributions of the air showers as a function of the energy of the incoming primary, as compared to standard events. Implications for their detection and some open issues on their possible discovery are also briefly addressed. \\vspace{1pc} ",
        "introduction": "At a cursory look, the inclusive cosmic ray spectrum, which has been measured over a wide range of energy, seems to be pretty simple in its functional shape, characterized by a (fast-falling) power-like behaviour over a huge energy range. Several mysteries may be hiding behind this apparent simplicity. Two of these mysteries, which have been puzzling theorists and experimentalists alike for several decades, and have been extensively discussed in this Symposium, are the knee ($10^{15}$ eV) and the ankle ($10^{19}$ eV). A long debate which has spurred from these observations has not reached any final conclusions as to the origin of the peculiar spectral dependence measured in these regions. Physics at the knee may involve, according to some interesting proposals, a critical behaviour of the QCD pomeron \\cite{White}, or, alternatively, disoriented chiral condensates and production of strangelets or other exotics, while the existence or the absence of the GZK cutoff \\cite{gzk} still needs to be completely assessed. Evading the cutoff using string relics \\cite{ccf},\\cite{cfp} has also been proposed as an interesting possibility.  On the other hand, recent formal suggestions from string theory and gauge theory \\cite{ED} lead to the exciting possibility of a brane picture of our world, with a gravity scale lowered from $10^{19}$ GeV down maybe to a fraction of a TeV and with the possibility of envisioning exotic new scenarios where ``black hole resonances'' might form even at colliders such as the LHC \\cite{DL}. In cosmic ray physics, in particular, the energy available at the impact of the primaries with the atmosphere can be very large and up to several hundreds of TeVs, and this points favourably towards a possible test of these new scenarios \\cite{Theodore}.  We have tested the implications of this hypothesis for cosmic ray air showers, by studying, in particular,  the lateral distributions and the multiplicities of events mediated by an intermediate black hole. Here we provide a brief summary of our results, while a detailed analysis can be found in \\cite{CCT}. ",
        "conclusions": ""
    },
    "0410/astro-ph0410550_arXiv.txt": {
        "abstract": "We calculate rates of Type I migration of protoplanets in a non-isothermal three-dimensional protoplanetary disk, building upon planet-disk models developed in previous work.  We find that including the vertical thickness of the disk results in a decrease in the Type I migration rate by a factor of $\\sim2$ from a two-dimensional disk. The vertical temperature variation has only a modest effect on migration rates since the torques at the midplane are weighted heavily both because the density and the perturbing potential are maximized at the midplane.  However, temperature perturbations resulting from shadowing and illumination at the disk's surface can decrease the migration rate by up to another factor of 2 for planets at the gap-opening threshold at distances where viscous heating is minimal. This would help to resolve the timescale mismatch between the standard core-accretion scenario for planet formation and the survival of planets, and could help explain some of the rich diversity of planetary systems already observed. ",
        "introduction": "The issue of planetary migration due to tidal interactions between the planet and a circumstellar disk was first addressed by \\citet{GT80}, who found that the potential of an orbiting planet excites Lindblad resonances in the disk, resulting in the transfer of angular momentum.  They found that the resulting torques on the planet are similar in magnitude, so that whether the planet moves inward or outward from angular momentum exchange depends on how the torques from the inner and outer disk balance out. \\citet{ward86} found that given a decreasing radial temperature profile, the torques from the outer disk are greater than those from the inner disk, resulting in a net torque that forces the planet inward. Planetary migration via this mechanism is referred to as Type I migration to distinguish it from Type II migration, which is the radial drift of planets massive enough to clear all material out of the Lindblad resonances, thereby forming a gap in the disk and moving together with the disk as it is accreted onto the star \\citep{wardA}.  Subsequent work on two-dimensional disks confirms the prediction that planets in protoplanetary disks migrate inward on short timescales on the order of $\\lesssim10^5$ yr for Earth mass planets \\citep[e.g.,][]{artymowicz93b,korycanskypollack,wardA}. \\citet[hereafter \\TTW{}]{TTW} find that including the three-dimensional density structure of an isothermal disk increases the migration timescale by a factor of 2 or 3 over a two-dimensional calculation. \\citet{menougoodman} calculate migration rates in a disk with a radial temperature profile expected in T Tauri disks and also find that the migration timescales increase when the thickness of the disk is accounted for. In a disk model that includes MHD turbulence, the density fluctuations near the planet can be sufficient to change the magnitude and even the sign of the total net torque on a small planet, so that the planet migrates as a ``random walk'' rather than a steady inward drift \\citep{nelsonpapaloizou4}.  While the effect of the density structure in the vicinity of a protoplanet on its migration rate has been studied in detail both analytically and in numerical simulations, the effect of local temperature variations has not \\citep[e.g.,][]{kley99,kdh01,bate}. Most simulations of planets embedded in disks employ a vertically isothermal equation of state, so that the disk has a simple radially dependent temperature structure. Similarly, the analytic calculations discussed above assume vertical isothermality for simplicity. In this paper, we analyze Type I migration torques from Lindblad resonances in a protoplanetary disk where viscous heating and stellar illumination are the primary heating sources.  Further, we examine how temperature variations induced by the presence of the planet itself affect Type I migration.  Disk temperature variations can affect the net Lindblad torques by changing the local pressure gradient in the disk. A pressure gradient shifts the angular frequency of the gas orbiting the star with respect to Keplerian velocities, which in turn shifts the positions of the Lindblad resonances \\citep{wardA}. The forcing function depends on the distance to the planet, so the torques on the planet vary with the locations of the Lindblad resonances.  To date, no one has addressed how a non-isothermal temperature structure will affect the magnitude of the angular momentum transferred at the Lindblad resonances and how that might affect migration rates.  A variation in temperature structure, such as those described in \\citet[hereafter \\paperone{}]{paper1} and \\citet[hereafter \\papertwo{}]{paper2}, due to radiative transfer effects, may affect the balance of migration torques.  In this paper, we address how the temperature structure of the disk affects the locations of Lindblad resonances and the consequences for Type I migration. In \\S\\ref{tidal_torques} we derive the equations used for calculating the torques on the planet from Lindblad resonances in a non-isothermal disk. In \\S\\ref{disk_model} we describe our adopted model disk and briefly discuss how the presence of a planet changes the temperature structure of the disk. We present our results in \\S\\ref{results} and discuss their implications in \\S\\ref{discussion}. Finally, in \\S\\ref{summary} we present our conclusions. ",
        "conclusions": "\\label{summary}  We have analyzed the effect of local disk temperature variations on Type I migration rates of planets embedded in protoplanetary disks.  We include both the vertical structure of the disk as well as temperature perturbations caused by the presence of the planet itself.  In comparison to the results of \\citet{wardA} for a two-dimensional disk model, the primary difference in migration rates is caused by including the vertical thickness of the disk. This is because the perturbing potential is diluted by the vertical distance from the midplane.  When we integrate the total torque at Lindblad resonances over the vertical extent of the disk, we find that this reduces the migration rate by about a factor of two. This is consistent with calculations by \\TTW{}, who find that the vertical dependence of the perturbing potential reduces migration rates by a factor of $\\sim2-3$ in their model disk.  We estimate the effect of including the vertical temperature and density profile of the model disk by comparing it to an equivalent vertically isothermal disk and find that the effect on the migration rates is modest. This is in part because the torque is weighted by disk density, which is highest at the midplane, and also because of the rapid decrease in torque density with disk height. Therefore disks with the same midplane temperature and density profiles will yield similar migration rates.  When we include the effects of radiative transfer on the perturbation caused by the planet on the disk, we find that the resulting temperature perturbations can further decrease the migration rate by up to about a factor of two, depending on the planet's mass and position. As the planet increases in mass, so does the temperature variation, which results in a general decrease in the migration rate. The effect is more prominent in planets further from the star, because viscous heating decreases faster than stellar illumination with increasing distance.  Viscous heating evens out temperature variations, reducing their impact on Lindblad torques.  Thus, the planet at the gap-opening threshold at the farthest distance in our sample is the most effective at reducing its own migration rate.  Our calculation of migration torques given the temperature structure in the photosphere of the disk shows a limiting case, since we assume that only the local temperature of the disk is relevant, and that the azimuthal dependence on temperature is unimportant.  This is justified by the fact that the perturbing potential is strongest closest to the planet, but in principle the azimuthal variation of the temperature could dilute the effect of local temperature variations on migration rates.  Another limitation of our planet-disk model is that we have not considered the non-linear perturbation to the density structure within a few Hill radii of the protoplanet. Since the tidal torques from Lindblad resonances scale with disk density, this effect may be very important. Numerical simulations universally show that spiral structure develops quickly around planets embedded in disks. Radiative transfer from stellar illumination upon these structures can affect the planet and disk both directly through heating, and indirectly by changing the positions and strengths of the Lindblad resonant torques.  Still, our model puts interesting constraints on the effect of the temperature and density structure on Type I migration. We have shown that shadowing and illumination at the surface of a protoplanetary disk can slow migration.  If the perturbations are large enough, they might be able to halt or even reverse Type I migration, and save planet embryos from the ``Shiva catastrophe'' \\citep{wardB}.  Perhaps this can explain why the frequency of giant planets around solar-type stars is relatively high, despite the short timescales predicted for Type I migration."
    },
    "0410/astro-ph0410285_arXiv.txt": {
        "abstract": "We present results from two-dimensional numerical simulations of the interactions between magnetized shocks and radiative clouds. Our primary goal is to characterize the dynamical evolution of the shocked clouds. We perform runs in both the strong and weak magnetic field limits and consider three different field orientations. For the geometries considered, we generally find that magnetic fields external to, but concentrated near, the surface of the cloud suppress the growth of destructive hydrodynamic instabilities. External fields also increase the compression of the cloud by effectively acting as a confinement mechanism driven by the interstellar flow and local field stretching. This can have a dramatic effect on both the efficiency of radiative cooling, which tends to increase with increasing magnetic field strength, and on the size and distribution of condensed cooled fragments. In contrast, fields acting predominately internally to the cloud tend to resist compression, thereby inhibiting cooling. We observe that, even at modest strengths ($\\beta_0\\lesssim100$), internal fields can completely suppress low-temperature ($T<100$ K) cooling. ",
        "introduction": "\\label{sec:introduction}  Shock waves are an important and common feature in both interstellar (ISM) and intergalactic (IGM) media.  They are triggered by such energetic phenomena as jets, supernovas, cloud-cloud collisions, and stellar winds and provide a means for transferring energy from such events into the ambient gas. Since the ISM and IGM are generally inhomogeneous, an important problem in astrophysics is understanding the interaction of these shocks with overdense clumps or clouds. A thorough review of this problem in the unmagnetized and non-radiative limits is provided by \\citet{kle94}.  Of special interest to us is large-scale shock-induced star formation, particularly in the neighborhoods of extragalactic radio jets. One of the first objects demonstrated to show a correlation between a radio jet and regions of active star formation was the nearest radio galaxy, Centaurus A \\citep[e.g.][]{bla75}.  Other examples have been found as the sensitivity and spatial resolution of radio and optical telescopes has improved.  For instance, ``Minkowski's Object'' is a peculiar small starburst system at the end of a radio jet emanating from the elliptical galaxy NGC 541, located near the center of the cluster of galaxies Abell~194 \\citep{van85a}.  Correlations between radio and optical emissions have also been observed in the so-called ``alignment effect'' in distant ($z>0.6$) radio galaxies \\citep{cha87,mcc87}.  These observations are most convincingly explained by models in which shocks generated by the radio jet propagate through an inhomogeneous medium and trigger gravitational collapse in relatively overdense regions, leading to a burst of star formation \\citep{beg89,dey89,ree89}. In a recent study \\citep{fra04a}, we investigated a key component of such models - the radiative shock-induced collapse of intergalactic clouds. For moderate cloud densities ($\\gtrsim 1$ cm$^{-3}$) and shock Mach numbers ($\\lesssim 20$), we found that cooling processes can be highly efficient and result in more than 50\\% of the initial cloud mass cooling to below 100 K. The cold, dense fragments that form are presumably the precursors to active star-forming regions.  In the current work we consider the effects of dynamically important magnetic fields in radiative shock-cloud collisions. Magnetic fields are known to be a pervasive element of the ISM and IGM and are often relevant in characterizing the local and global dynamical behaviors of these media. In shock-cloud interactions, magnetic fields can act to suppress destructive hydrodynamic instabilities by providing additional tension at the interface between the cloud and the post-shock background gas \\citep{nit81,mac94}. Magnetic fields can also limit the growth of disruptive vortices that form in the wake of the cloud, again primarily due to tension in the magnetic field lines as they are wound up within the vortices.  Strong external magnetic fields can also increase the compression of the shocked cloud material, due to the increased external magnetic pressure.  This compression enhances the radiative efficiency of the cloud and allows additional cooling beyond that achievable without magnetic fields.  However, strong internal magnetic fields can resist compression, thereby inhibiting cooling of the cloud. In this paper we explore these competing effects and the general role of magnetic fields in radiative shock-cloud collisions. This paper also reports on the addition of magnetic fields to our astrophysical hydrodynamics code, Cosmos.  We proceed in \\S\\ref{sec:models} by describing the models considered in this work.  In \\S\\ref{sec:methods} we describe our implementation of radiation MHD.  Our results are presented in \\S\\ref{sec:results} and discussed further in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have presented results from a series of two-dimensional shock-cloud simulations with the goal of highlighting the importance of different physical processes, including the interplay between hydrodynamic, radiative, and magnetic effects. These simulations represent the first such calculations we know of that simultaneously consider magnetic fields and radiative cooling.  To facilitate easy comparison, we included runs with different combinations of physical processes active. We summarize our main results as follows:  1.  Unmagnetized, non-radiative clouds are destroyed on a few dynamical timescales through hydrodynamic (Kelvin-Helmholtz and Rayleigh-Taylor) instabilities \\citep{kle94}.  2.  In the cooling-dominated regime, radiative clouds are {\\em not} destroyed.  Instead, they form dense, cold filaments, which are presumably the precursors to active star-forming regions \\citep{mel02,fra04a}.  3.  Tension in magnetic field lines along the surface of a cloud can suppress the growth of hydrodynamic instabilities, thus increasing the cloud's survivability even without radiative cooling \\citep{mac94,jon96}.  This is true for our external field cases ($B_x$ and especially $B_y$) whenever the fields achieve sufficient strength ($\\beta \\lesssim 10^4$ to suppress Rayleigh-Taylor and $\\beta \\lesssim 0.02$ to suppress Kelvin-Helmholtz). On the other hand, internal fields that do not thread through the cloud surface (such as our $B_z$ cases) are unable to suppress the growth of hydrodynamic instabilities.  4.  External magnetic field lines that are stretched over the surface of a cloud can greatly enhance its compression.  For radiative clouds, this can dramatically enhance the cooling efficiency.  For instance, the fraction of cloud material cooling from $10^4$ K to below 100 K increases from $\\sim 0.7$ without magnetic fields (run C) to $>0.9$ with a $\\beta_0=4$, $B_y$ field (run BY4C).  This enhancement is negligible, however, for the $B_x$ field orientation or an initially weak field aligned along $B_y$ ($\\beta_0 > 100$).  5.  Internal magnetic field lines resist compression in a cloud. For radiative clouds, this can dramatically reduce the cooling efficiency.  For instance, in simulations with a $B_z$ field of only modest initial strength ($\\beta_0\\lesssim100$), the cloud material is prevented from cooling below 100 K.  A very strong initial field ($\\beta_0\\sim1$) can even prevent cooling below 1000 K.  Obviously, the field configurations studied in this work are highly idealized.  Real astrophysical fields are likely to have much more complex topologies \\citep{val03}.  In particular, magnetic fields in galactic star-forming clouds are probably linked to underlying turbulent flow patterns \\citep{mac04}.  In that context, magnetic fields are important in stabilizing such clouds against collapse, a role similar to that of the {\\em internal} fields in this work. However, those star-forming clouds are often very cold ($T<100$ K) and subject to important physical process not considered in this work, such as neutral-ion drift, so we caution the reader against attempting to extrapolate our results to that regime.  Due to resolution requirements and computational limitations, the simulations in this work were carried out in two-dimensional Cartesian geometry, and so the clouds represent slices through infinite cylinders.  This special geometry likely affects some of our conclusions.  Here we speculate on how these results may change in three-dimensional simulations:  Since the initial compressive shock in the cloud is highly symmetric, the additional convergence expected in three-dimensional models might lead to stronger compressions and enhanced cooling in radiative clouds even without magnetic fields. Three-dimensional simulations would also provide an additional degree of freedom for fragmentation through dynamical instabilities.  The role of magnetic fields in three dimensions may be more complicated. For a spherical cloud, we lose the distinction between our $B_y$ and $B_z$ runs.  Such transverse fields will cause enhanced compression of the cloud along one direction (the initial direction of the field), but will be unable to prevent expansion in the perpendicular direction. This lateral expansion can enhance the growth rate of the Rayleigh-Taylor instability \\citep{gre99,gre00}, an action that was prevented in this work by the assumed symmetry of the two-dimensional runs. Thus, in contrast to two-dimensional results, magnetic fields in three-dimensional simulations can hasten the destruction of non-radiative clouds. It remains to be seen how radiative cooling would modify this conclusion.  Self-gravity has been neglected in this study on account of the constrained cylindrical geometry. Although the cloud parameters are specifically chosen such that self-gravity is negligible initially, the local free-fall timescale becomes significantly shorter during the latter stages of compression. For the runs with radiative cooling active, self-gravity becomes important at approximately the time we stop the simulations.  It would, therefore, be interesting to follow shock-cloud simulations in three-dimensions with radiative cooling, magnetic fields, and self-gravity included. To do this at comparable resolution to that used in this work will require the use of adaptive mesh refinement to concentrate resolution around the cloud fragments. Adaptive gridding is currently being added to the Cosmos code, and we plan to revisit this problem in future work."
    },
    "0410/astro-ph0410416_arXiv.txt": {
        "abstract": "We report the discovery of SDSSJ1049+5103, an overdensity of resolved blue stars at ($\\alpha_{2000}$, $\\delta_{2000}$) = (162.343, 51.051). This object appears to be an old, metal-poor stellar system at a distance of 45$\\pm 10$ kpc, with a half-light radius of 23$\\pm 10$ pc and an absolute magnitude of M$_V$ = -3.0$^{+2.0}_{-0.7}$.  One star that is likely associated with this companion has an SDSS spectrum confirming it as a blue horizontal branch star at 48 kpc.  The color-magnitude diagram of SDSSJ1049+5103 contains few, if any, horizontal or red giant branch stars, similar to the anomalously faint globular cluster AM 4. The size and luminosity of SDSSJ1049+5103 places it at the intersection of the size-luminosity relationships followed by known globular clusters and by Milky Way dwarf spheroidals.  If SDSSJ1049+5103 is a globular cluster, then its properties are consistent with the established trend that the largest radius Galactic globular clusters are all in the outer halo.  However, the five known globular clusters with similarly faint absolute magnitudes all have half-mass radii that are smaller than SDSSJ1049+5103 by a factor of $\\gtrsim$ 5.  If it is a dwarf spheroidal, then it is the faintest yet known by two orders of magnitude, and is the first example of the ultra-faint dwarfs predicted by some theories. The uncertain nature of this new system underscores the sometimes ambiguous distinction between globular clusters and dwarf spheroidals.  A simple friends-of-friends search for similar blue, small scalesize star clusters detected all known globulars and dwarfs closer than 50 kpc in the SDSS area, but yielded no other candidates as robust as SDSSJ1049+5103. ",
        "introduction": "Milky Way globular clusters are invaluable pieces in the puzzle of galaxy formation. At present, their properties support a general picture of Galactic halo formation as a combination of accretion and dissipative collapse (see review in \\citealt{mackey04}). However, the detailed interpretation of globular cluster (GC) properties in the context of galaxy formation is complex. One outstanding problem is the sometimes ambiguous distinction between GCs and dwarf spheroidal galaxies (dSphs).  For example, a few Milky Way GCs, such as $\\omega$ Cen, have a spread in stellar age and metallicity similar to that seen in many dwarf galaxies \\citep{ashman98}, and have absolute magnitudes that overlap those of known dSph galaxies. A small number of faint GCs have radial profiles that are well fit by an NFW profile (e.g. Palomar 13; \\citealp{cote02}) or have central densities similar to those of dSphs (e.g. Palomar 14; \\citealp{harrisGCcat}) and thus may be the remnants of a stripped dSph.  The relationship between globular clusters and dSphs is particularly interesting in light of recent predictions for low mass substructure around the Milky Way (\\citealp{klypin99,moore99,bullock00,benson02,susa04,kravtsov04}, among others).  It is difficult to determine whether GCs ever contained a substantial amount of non-baryonic dark matter \\citep{ashman98}, which would arguably put them in the category of ultra-faint dwarf galaxies.  If some globular clusters are embedded in extended dark matter halos, the dark matter may not be dynamically important within the extent of the observable stellar distribution.   There are $\\sim 150$ known globular clusters and 9 known dSphs orbiting the Milky Way.  The total number of known clusters has increased by just a few percent over the last twenty-five years \\citep{harris97,harrisGCcat,ortolani00,hurt00,irwin95,ortolani93} and nearly all of the new globular clusters lie at low Galactic latitude. Only one Milky Way dSph has been discovered since 1990.  The lack of new GCs or dSphs at $|b|$ $>$ 30$^{\\circ}$ could lead some to believe that all high latitude systems have been discovered.  However, one anomalously faint GC (AM 4; M$_V$ = +0.2) was discovered serendipitously more than 20 years ago \\citep{madore82}, suggesting that other ultra-faint star clusters may still reside undetected in our halo.  Furthermore, the advent of the Sloan Digital Sky Survey (SDSS; \\citealt{york00}) could lead to the discovery of similar systems, should they exist \\citep{willman02}. In this paper, we report the discovery of SDSSJ1049+5103, a new ultra-faint, stellar system in the outer halo of the Milky Way. We estimate and discuss some properties of SDSSJ1049+5103 in comparison to both globular clusters and Milky Way dwarf spheroidal galaxies. ",
        "conclusions": "In this paper, we report the discovery of SDSSJ1049+5103, a new stellar system that is likely in the outer halo of the Milky Way. Based on comparison with Palomar 5, this new system appears to be $\\sim$ 50 kpc from the Galactic center, have a half light radius of 23 pc, and M$_V = -3.0^{+2.0}_{-0.7}$.  SDSSJ1049+5103 has a size and luminosity that places it at the intersection of the size-luminosity locus followed by Milky Way globular clusters and and that followed by Milky Way dSphs and nearby faint red galaxies.  Both the fundamentally ambiguous distinction between some globular clusters and dSph galaxies and the fact that SDSSJ1049+5103 is unusual relative to the vast majority of GCs leaves open the possibility that it is an extreme dwarf galaxy nearly two orders of magnitude fainter than Ursa Minor, the faintest known MW dSph.  Furthermore, some theories predict the presence of low central surface density, ultra-faint dSphs (e.g. \\citealp{benson02}) such as SDSSJ1049+5103.  If SDSSJ1049+5103 is a globular cluster, then its properties are consistent with undergoing tidal disruption.  Neither the \\citet{willman02} survey nor a friends-of-friends search revealed additional companions similar to, or even a bit fainter than, SDSSJ1049+5103 in the 5000 square degrees analyzed thus far.  This suggests that there is not a substantial unknown population of similar companions closer than $\\sim$ 50 kpc.  We are in the process of obtaining both deep, wide field imaging to accurately measure the spatial distribution of SDSSJ1049+5103 and spectra of individual stars to measure ages, metallicities, and line-of-sight velocities.  Specifically, deeper imaging may distinguish between a King and NFW surface brightness profile and may also reveal tidal features, which would provide strong constraints on its current mass \\citep{moore96}."
    },
    "0410/astro-ph0410620_arXiv.txt": {
        "abstract": "We present a mid-infrared analysis of 35 quasars with spectroscopic redshifts selected from the Spitzer Wide-area InfraRed Extragalactic Survey (SWIRE). We discuss their optical and mid-infrared (MIR) colors, and show that these quasars occupy well defined regions in MIR color-color space. We examine the issue of type-I AGN candidate selection in detail and propose new selection methods based on mid-IR colors. The available multi-band data allows us to construct two new, well-sampled quasar templates, covering wavelengths from the ultraviolet to the MIR. ",
        "introduction": "Efforts to understand the physics that drive the activity in active galactic nuclei (AGN) have been underway for decades, and have recently seen rapid progress thanks to improvements in both instrument design and theoretical modeling. Observationally, new facilities have provided important pieces of information, which when combined with advances in modeling have allowed a reasonably clear picture of the emission mechanisms in AGN to emerge. The optical continuum (up to 1 \\mum) of type-I AGN, is seen to be a power law, which can be explained by the existence of an accretion disk around a central supermassive black hole. At longer wavelengths, one starts observing emission thought to arise from a dusty torus, which, depending on the geometry, peaks at wavelengths between 15 and 100 \\mum. This emission is thought to be the reprocessed emission of the UV/optical radiation from the accretion disk by the particles composing the torus, namely silicate and graphite grains (e.g. \\citealt{granato94}; \\citealt{efstathiou95}; \\citealt{nenkova02}).  Studies with the Infrared Space Observatory (ISO) showed that the mid-infrared (MIR) window can be used to distinguish AGN from starbursts \\citep{rigo99,laurent00,farrah02,farrah03}. Furthermore, recent work conducted on a sample consisting of 25 quasars detected by the European Large Area ISO Survey (ELAIS) at 15 \\mums \\citep{afonso04} showed the potential of combined optical and MIR observations in the construction of quasar samples, and in achieving a better understanding of the physical properties of AGN. These studies also showed, however, the need for deeper multi-band observations in the MIR in order to refine physical constraints on the torus models and unified schemes.  With the advent of Spitzer it is now possible to obtain MIR photometry for very large samples of galaxies. In particular, the Spitzer Wide-area InfraRed Extragalactic Survey (SWIRE; \\citealt{lonsdale03}; \\citealt{lonsdale04}) offers an unprecedented opportunity to study the MIR properties of AGN thanks to both the large areal coverage ($\\sim 50$ deg$^2$) and to the number of bands in which the fields are being observed (IRAC bands 3.6, 4.5, 5.8 and 8 \\mums and MIPS bands 24, 70 and 160 \\mum). In this paper we investigate properties of a complete optically selected spectroscopically confirmed type-I quasar sample from the Sloan Digital Sky Survey (SDSS), consisting of 35 objects. The paper is structured as follows. Section \\ref{sample} describes the MIR, optical and near-IR data available for this sample; these properties are then used in Section \\ref{midselection} in order to determine type-I AGN candidate selection criteria from multi-color data. Section \\ref{seds} discusses the SEDs of the objects. Section \\ref{bhmass} deals with the Black Hole (BH) mass and IR luminosity estimates. Notes on individual objects are given in Section \\ref{notes}. Finally, Section \\ref{discuss} presents a discussion of the results. ",
        "conclusions": "\\label{discuss} This paper presents an analysis of a sample of 35 spectroscopically confirmed SDSS quasars within the SWIRE EN1 field, observed with Spitzer as part of the SWIRE Legacy Survey. Even though SWIRE EN1 is a small part of the whole 50 deg$^2$ SWIRE survey ($\\sim$18\\%) and SDSS DR2 only covers a third of it, this study gives a good initial indication of the AGN science that can be carried out with SWIRE data.  All SDSS quasars within the SWIRE EN1 field have Spitzer counterparts. The distinct location of the spectroscopically confirmed quasars on MIR color-color plots greatly contributes to the selection of type-I AGN candidates, independent of their optical properties. They are redder than the majority of the population at least up to 8 \\mums (restframe from 1.7 to 6.1 \\mums for the redshift range covered) and continue forming a rather compact clump in the redder colors. Appropriate color-color combinations can be found in order to separate quasars of specific (but large) redshift ranges from the various types of contaminant galaxies. A comparison between model predictions and selected quasar candidates indicates the existance of a large number of obscured AGN, some of which probably reside among the proposed candidates.  BH masses, estimated from emission lines, show a tendency with redshift up to $z \\sim 1$, however the small number statistics do not allow any conclusion on the dependence or not on bolometric luminosity.  Even though large multi-wavelength spectroscopic efforts are still required in order to understand the nature and propeties of AGN, we demonstrate that Spitzer MIR data contribute significantly towards this goal."
    },
    "0410/hep-ph0410065_arXiv.txt": {
        "abstract": "Assuming that cosmic rays entering the Earth's atmosphere contain a small admixture of nuggets of strange quark matter in form of {\\it strangelets} one can explain a number of apparently \"strange\" effects observed in different cosmic rays experiments. We shall demonstrate here that the mass spectrum of such strangelets filles the \"nuclear desert\" gap existing between the heaviest elements observed in Universe and the next \"nuclear-like objects\" represented by neutron and strange stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/hep-ph0410247_arXiv.txt": {
        "abstract": "We provide a simple but robust argument that primordial black hole (PBH) production generically does {\\em not} exceed astrophysical bounds during the resonant preheating phase after inflation. This conclusion is supported by fully nonlinear lattice simulations of various models in two and three dimensions which include rescattering but neglect metric perturbations.  We examine the degree to which preheating amplifies density perturbations at the Hubble scale and show that at the end of the parametric resonance, power spectra are universal, with no memory of the power spectrum at the end of inflation. In addition we show how the probability distribution of density perturbations changes from exponential on very small scales to Gaussian when smoothed over the Hubble scale --  the crucial length for studies of primordial black hole formation -- hence justifying the standard assumption of Gaussianity. ",
        "introduction": "Primordial black holes (PBHs) span a wide range of mass scales and are typically much smaller than the solar mass ($ \\sim 10^{33} {\\rm g} $) and may be formed in the early universe~\\cite{carr}. PBHs may form from  the gravitational collapse of large density fluctuations at horizon (i.e. Hubble scale $k = aH$) crossing  in the radiation dominated universe.  A PBH formed at the Planck time $ \\sim 10^{-43} {\\rm sec} $ will have a mass $ \\sim 10^{19} {\\rm GeV} $, while masses around $ \\sim 10^{15} {\\rm g} $ are formed at $ \\sim 10^{-23} {\\rm sec} $ ( see for example~\\cite{carr2} ).  The evaporation time for a PBH mass $ \\sim 10^{15} {\\rm g} $ is nearly the present age of the universe, so PBH with smaller masses than this would have evaporated in the past, unloading a potentially vast amount of entropy. The success of the standard cosmology and the observation of the cosmic rays severely constrains the abundance of PBHs for various masses and provides useful constraints on inflationary and early universe physics.  For example, Big Bang Nucleosynthesis limits the PBH abundance in the mass range $ 10^{6} \\sim 10^{13} {\\rm g} $, by limiting the entropy from  Hawking radiation or requiring that it does not modify the cosmological composition of the light elements~\\cite{starobinskii}. PBH of mass $ 10^{15} {\\rm g} $ evaporating now will emit particles such as $ \\gamma$-rays which are constrained by observations of the extragalactic $\\gamma$-ray background which imply $ {\\Omega}_{PBH}  < 10^{-8} $~\\cite{page}. For $ > 10^{15} {\\rm g} $, the limit on the PBH abundance is obtained from requiring that $ {\\Omega}_{PBH} $ does not exceed unity.  These observational constraints on PBHs provide a powerful probe of the primordial fluctuations. The upper bound on the abundance of PBHs directly leads to that on the density fluctuation at horizon crossing when PBH are formed. Therefore the scales of fluctuations relevant to PBH formation are much smaller than those associated with the Cosmic Microwave Background and the large-scale structure. This makes studying PBHs important. In the past, for example, constraints on the density perturbation spectrum were obtained by studying PBH formation~\\cite{bringmann}. While the requirement that PBHs are not over-produced yields useful information about the early universe, PBH can be an interesting dark matter candidate in smaller abundances~\\cite{ivanov}.  In this paper we will show that typically PBH are not over-produced during the violent non-equilibrium phase of preheating that follows the end of many inflationary models. This follows from three key observations: (1) the peak of the density perturbation spectrum typically lies at scales smaller than the Hubble scale. (2) The peak corresponds to density contrasts of order unity. (3) The slope of the spectrum around the horizon size is three (in $3d$). Putting these together we typically find that at the horizon scale relevant for PBH formation, the density contrast is around an order of magnitude too small to over-produce PBH.  Nevertheless, the density perturbation on the horizon scale is significantly enhanced by preheating (by several orders of magnitude) compared with the no-resonance case and hence preheating is important in understanding the potential astrophysical and cosmological implications of PBH.  In studying the production of PBHs, one usually assumes that the probability distribution of density fluctuations at horizon crossing is Gaussian. This assumption is critical because the density perturbations which collapse to black holes are very rare: several $ \\sigma $ fluctuations (otherwise PBH will be over-produced in all cases) and therefore production of PBHs is sensitive to the tail of the distribution. Indeed Bullock and Primack~\\cite{bullock} found that in some inflation models large perturbations are suppressed relative to a Gaussian distribution, resulting in a significant change in a number of PBH. We study the validity of the Gaussian assumption in a later section.  Among various possible scenarios that might over-produce PBHs, we focus on preheating after inflation. Preheating is a process in which energy transfer occurs rapidly from inflaton field to another field due to the non-perturbative effects during the oscillating phase of the inflaton~\\cite{traschen}. This process significantly differs from the usual reheating scenario, where inflaton decays perturbatively to another particles, in a sense that in preheating most energies of inflaton field converts to created particles only during the several oscillations of inflaton and even the massive particle which is much heavier than the inflaton can be created. It has been understood that parametric resonance occurs generically at the first stage of reheating~\\cite{kofman1}. The parametric resonance does not last long because the rapid increase of created particles eventually affect the motion of background field and the created particles scatter off each other, removing the particles from the resonance band. By these effects, the resonance becomes inefficient and the decaying process of inflaton is described by usual single-body decay theory which finally leads to the thermal equilibrium state~\\cite{felder00}.  There are many works on preheating. The first stage of preheating, where the backreaction is negligible and the linear approximation is valid, was studied in detail in~\\cite{kofman2, pgreene, shtanov}. Parametric resonance including  metric perturbations in order to see the behavior on super horizon scales were studied in~\\cite{taruya, kodama, bassett1, bassett2, finelli1, jedamzik, liddle}. As in the case without metric perturbations, there is a crucial difference between the single and multiple field cases. The analysis of fully non-linear preheating including gravity is very difficult. Before now, different approximations such as mean-field approximation~\\cite{bassett2, jedamzik, zibin, bassett3, bassett4,tsujikawa, boyanovsky}, lattice simulation~\\cite{lattice} without metric perturbation and one dimensional fully non-linear calculations~\\cite{finelli2, parry, easther} have been done for studying the various effects caused by preheating on the present universe.  Green and Malik~\\cite{agreen} argued that PBHs {\\em will} be overproduced due to the amplification of fluctuations during preheating for many parameter regions in a two-field massive inflation model based on the results of~\\cite{liddle2}, which takes into account of the second order fluctuation of $ \\chi $ field. Put simply, their results suggested that the backreaction timescale was smaller than the timescale for the over-production of PBH.  Bassett and Tsujikawa~\\cite{bassett3} studied PBH production in the two-field massless inflation model including the effect of backreaction via the Hartree-Fock approximation  and found that PBH overproduction might occur, if the  probability distribution of fluctuation at horizon crossing was assumed to be chi-squared (which lowers the threshold mass variance, $\\sigma$). However, they found that PBH were not over-produced if the distribution was assumed to be Gaussian {\\em and} the density field was smoothed on the scale of the horizon. Nevertheless, that analysis was limited since it neglected the mode-mode coupling  effects of rescattering and hence there was an open question both as to the underlying probability distribution of the density fluctuations and the contribution of rescattering to horizon-scale fluctuations.  We address both of these issues in this paper. We study PBH production due to preheating via two and three-dimensional lattice simulations which automatically include the effects of backreaction and rescatterings. We modified the C++ code LATTICEEASY written by Felder and Trachev~\\cite{latticeeasy}. In lattice simulations, the evolution equations for the scalar fields (and also the scale factor) are solved in real (as opposed to Fourier) space ($ N=2048^2 $ for 2 dimensional simulation and $ N=128^3 $ for 3 dimensional simulation). Metric perturbations were not included, so we cannot apply this method to the dynamics on super horizon scales. We followed the evolution of both the scalar fields and the total density perturbation. We found that PBH are not overproduced and, interestingly, that the power spectrum at the end of preheating has a universal feature, that is, it is determined by the preheating dynamics and does not depend on the initial conditions. We also studied the probability distribution of the density perturbation at horizon crossing and found that it remained Gaussian which seems to be valid even in the tail of distribution.  A brief summary of our paper is as follows. Section II gives a brief review of calculating the PBH abundance formed from the large density perturbation. Section III describes the model of preheating in this simulation. In section IV, we consider the initial conditions of scalar fields. Section V shows the numerical results of lattice simulation. We will interpret the non-linear behavior of fluctuations and study the production of PBHs. Section VI discusses the probability distribution of fluctuations at horizon crossing we assume Gaussian in usual and Section VII is a conclusion. ",
        "conclusions": "We have studied the formation of black holes during preheating after inflation. Preheating provides a challenging framework to illucidate various complex physical processes such as non-equilibrium, non-perturbative field theory in expanding backgrounds. In addition, since preheating is generic in some regions of parameter space for many inflationary models, the issue of whether primordial black holes (PBH) are overproduced is an important one.  To address this we have performed two and three dimensional lattice simulations of several different inflaton potentials (conformal and massive) and parameter regions. These simulations automatically incorporate all non-linear effects such as backreaction and rescattering of fields. We found no evidence for over-production of PBH, in contrast to earlier expectations. In addition we found that although highly non-Gaussian on very small scale, the spectrum of density perturbations is effectively Gaussain on the horizon (i.e. Hubble) scale.  Our results can be understood simply. For all cases, we found that when the amplitude of density perturbations at about the shortest wavelength in the resonant band becomes of order unity,  the growth due to parametric resonance terminates due to backreaction.  At the end of preheating the final density spectrum is universal, with a blue power spectrum, $ \\propto k^D $ (D = 2,3) on horizon scales, depending on whether the simulation was two ($D=2$) or three ($D=3$) dimensional.  Since the peak of the spectrum is scales significantly shorter than the horizon scale, the amplitude of the density perturbation at the horizon scale extrapolated from the peak is typically about an order of magnitude  below the threshold for PBH overproduction.  We gave an explanation of this universality in the slope of the final power spectrum on large scales as a result of loss of coherence due to parametric resonance.  These results argue for the view that generically PBHs will not be produced so much as to violate astrophysical constraints, even in the case of strong preheating. This result also applies to the cases with other models and  parameters regions if our interpretation of the power spectrum on large scales is universally correct.  The result obtained in this paper is different from the claim by Green and Malik~\\cite{agreen}, where they found that PBHs are likely to be overproduced. They estimated the time when backreaction becomes significant as well as the time when the amplitude of density perturbations exceeds the threshold separately based on linear approximation. Comparison of these two times was used to give a criteria for PBH formation. However, for example, in Ref.~\\cite{kofman2} the time when the backreaction becomes significant in the case with $m=10^{-6} m_{pl}$ and $ g=10^{-3} $ is estimated to be $ \\sim 90 $, which is slightly later than our numerical result (See Fig.~\\ref{backreaction}). Since the growth of perturbation amplitude is exponential, a small error in the back reaction time can lead to wrong conclusions. In calculating the abundance of produced PBHs, only $ 10 \\sim 20 $ percent error of backreaction time can lead to the opposite conclusion. By contrast, our conclusion is based on self-consistent simulations and rather robust qualitative observations. There is no delicate comparison of different time scales.  In this paper, we considered standard slow roll inflation. In such cases, preheating occurs at rather high energies. Therefore the mass of PBHs formed in the present context is too small to avoid evaporation before the Big Bang Nucleosynthesis. Hence, those PBHs are not subject to any observational constraint even if PBHs are produced abundantly, unless PBHs leave Planck mass relics~\\cite{barrow}. (Hence we have assumed that PBH will leave Planck mass relic throughout this paper. ) Even if relics are formed, the constraint on the mass fraction of produced black holes is $ \\beta < 10^{-20} $~\\cite{carr5}, we can therefore conclude that production of PBHs by preheating does not give a serious constraint on such simple models of preheating as discussed in this paper.  However, our qualitative results will also apply for preheating at lower energies. Let us consider the possibility of making more massive black holes by preheating, say, in the case of massive inflaton models. Now we consider to vary the inflaton mass $m$ and the value of $\\phi$-filed at the beginning of preheating $\\phi_0$ within $\\phi_0 \\alt m_{pl}$ (As an example of realizing small value of $\\phi_0$, we can consider the hybrid inflation model~\\cite{linde}.). The key quantity is the ratio of $ k_{max} $ given in (\\ref{con4}) to the Hubble parameter, which is estimated as \\begin{equation} \\frac{k_{max}}{H} \\sim \\frac{m_{pl}}{\\phi} q^{1/4}. \\label{con5} \\end{equation} This ratio is independent of $m$ and is the smallest for $\\phi\\approx m_{pl}$. Then the same estimate given in (\\ref{remass7}) applies as an upper bound for $\\delta_H$. On the other hand, lowering $\\beta$ down to $10^{-30}$ only change the upper bound on $\\sigma$ from 0.031 to 0.026. Thus we can say that overproduction of more massive PBHs due to parametric resonance is also unlikely.  We have also confirmed Gaussianity of the probability distribution of density perturbations at the horizon scale, which is assumed in the estimate of the production rate of PBHs. The appearance of Gaussianity is in accordance with the interpretation of the spectrum on large scales. If perturbations become uncorrelated beyond the shortest resonant scale, perturbations at the horizon scale are given by average of many statistically independent random variables. Thus Gaussianity  naturally follows from the central limit theorem."
    },
    "0410/astro-ph0410402_arXiv.txt": {
        "abstract": "The NeXT  mission has been proposed to study high-energy non-thermal phenomena in the universe. The high-energy response of the super mirror  will enable us to perform the first sensitive  imaging observations  up to 80 keV. The focal plane detector, which combines a fully depleted X-ray CCD and a pixellated CdTe  detector, will provide spectra and images in the wide energy range from 0.5 keV to 80 keV. In the soft gamma-ray band upto $\\sim$ 1 MeV, a   narrow field-of-view Compton gamma-ray telescope utilizing several tens of layers of thin Si or CdTe detector will provide  precise spectra with much higher sensitivity than present instruments. The continuum sensitivity will reach several $\\times$ 10$^{-8}$ photons/s/keV/cm$^{2}$ in the   hard X-ray region and a few $\\times$ 10$^{-7}$ photons/s/keV/cm$^{2}$ in the soft $\\gamma$-ray region. ",
        "introduction": "The hard X-ray  and gamma-ray bands  have long been recognized as  important windows for exploring the energetic universe. It is in these energy bands that non-thermal emission, primalily due to accelerated high energy particles, becomes dominant. However, by comparison with the soft X-ray band, where the spectacular data from the XMM-Newton and Chandra satellites are revolutionizing our understanding of the high-energy universe, the sensitivities of hard X-ray missions flown so far, or currently under construction, have not dramatically improved over the last decade. In order to study the energy content of non-thermal emission and to draw a more complete picture of the non-thermal universe, observations by highly sensitive missions are important.  The NeXT (Non-thermal Energy eXploration Telescope) mission proposed in Japan is a successor to the Astro-E2 mission, and is optimised to study the high-energy non-thermal universe (Fig.\\ref{Fig:NeXT}) \\cite{Ref:Proposal,Ref:Kunieda}. NeXT will carry three hard X-ray telescopes (HXTs) for the wide band X-ray imager (WXI), one soft X-ray telescope (SXT)  for the soft X-ray spectrometer (SXS), and a soft $\\gamma$-ray detector (SGD). With the first imaging observations up to 80\\,keV, NeXT is expected to achieve two orders of magnitude improvement in the sensitivity in the hard X-ray region. The SGD utilizes a new concept of a narrow-FOV Compton telescope\\cite{Ref:Takahashi_Yokohama,Ref:Takahashi_SPIE} to improve the sensitivity in the soft $\\gamma$-ray region for matching with the sensitivity of the HXT/WXI combination. The extremely low background brought by the SGD will allow us to measure the precise $\\gamma$-ray spectrum upto 1 MeV. Furthermore, as described in the separate paper\\cite{Ref:Tajima-pol}, the SGD will be  sensitive to the polarization of the  incident $\\gamma$-rays, and will potentially be the first hard X-ray polarimeter in orbit. In addition to these high energy instruments, the soft X-ray spectrometer (SXS) will map out the velocity field in objects  by means of X-ray line profiles. The SXS is based on a transition edge sensor (TES) and has both spectral and spatial resolution much better than the X-ray calorimeter (XRS) onboard the Astro-E2\\cite{Ref:Mitsuda}. In this paper, we describe  the current ideas about the WXI and the SGD.    \\begin{figure} \\centerline{\\includegraphics[height=7cm]{NeXT.eps}} \\caption{Artist's drawing of the NeXT satellite. The focal  length of the HXT is 12 m. } \\label{Fig:NeXT} \\end{figure} ",
        "conclusions": "The line and continuum sensitivities of the NeXT mission for 100 ks observation are  shown in Fig.\\ref{Fig:Continuum} and Fig.\\ref{Fig:Line}, together with those of  other missions. The continuum sensitivity could reach several $\\times$ 10$^{-8}$ photons/s/keV/cm$^{2}$ in the   hard X-ray region and a few $\\times$ 10$^{-7}$ photons/s/keV/cm$^{2}$ in the soft $\\gamma$-ray region. The high-energy response of the super mirror  of NeXT will enable us to perform the first sensitive  imaging observations  up to 80 keV. By combining  an X-ray CCD and a  CdTe pixel  detector, the WXI is an ideal solution for the focal plane detector of the super mirror, providing fine imaging capability and high spectral resolution with almost full detection efficiency. The narrow field-of-view Compton $\\gamma$-ray telescope realized by the SGD extends the bandpass to well above the cutoff for the hard X-ray telescope and allows us to study the high energy end of the particle spectrum through the sensitive observation of the  $\\gamma$-ray spectrum up to 1 MeV. By combining these detectors, we expect to achieve an unprecedented level of sensitivity in the hard X-ray and sub-MeV $\\gamma$-ray region for both line and continuum emission.  \\begin{figure} \\centerline{\\includegraphics*[width=8cm,angle=90]{sensitivity_cont_crab2.ps}} \\caption{The expected WXI and SGD senstivities for continuum emissions from a point source, assuming an  observation time of 100 ks.} \\label{Fig:Continuum} \\end{figure}  \\begin{figure}  \\centerline{\\includegraphics*[width=8cm,angle=90]{sensitivity_line_gamma.ps}} \\caption{The expected WXI and SGD senstivities for line emissions from a point source, assuming an  observation time of 100 ks.} \\label{Fig:Line} \\end{figure}"
    },
    "0410/astro-ph0410128_arXiv.txt": {
        "abstract": "{ The emission from galactic black holes and Seyfert galaxies is generally understood in term of two-phase models (Haardt \\& Maraschi 1991, 1993). Such models postulate that a hot plasma ($\\sim$10$^{9}$\\,K) coexists with relatively colder material ($\\sim$ 10$^{6}$\\,K) in the inner part of the accretion flow. We present the first simulated broad-band spectra produced by  such a system and accounting simultaneously for energy balance and Comptonisation in the hot phase, together with reflection, reprocessing, ionization and thermal balance in the cold phase. This was made possible by coupling three radiative transfer codes: a non-linear Monte-Carlo code (NLMC), a photo-ionization code TITAN and a linear Monte-Carlo code NOAR. The equilibrium comptonisation spectrum appears to be sensitive to the shape of the reprocessed spectrum that, in turn, depends on the ionization parameter, but also on the structure of the irradiated cold material. This is illustrated by a comparison of simulations assuming constant density or a constant pressure in the cold phase. We also compare our results with simplified models where reprocessing is approximated by a blackbody spectrum. Our detailed treatment leads to noticeably different spectral energy distributions (SEDs) characterised by harder X-ray spectra. { Even at low ionization parameters ($\\xi \\sim 300$ erg~s$^{-1}$~cm) the commonly used blackbody approximation is poor, leading to X-ray spectra that are too soft}. The effect, however, seems not to be strong enough to reconcile the slab corona model with the hardest observed spectra, { unless the reflector has a constant density and the ionization parameter is large, of the order of 10$^{4}$ erg~s$^{-1}$~cm.} ",
        "introduction": "At least two distinct mediums are required to produce the main features observed in the broad band spectra of radio quiet active galactic nuclei (AGN) and  black hole binaries (BHBs) in the low/hard state. Their hard X-ray  emission exhibits a power-law spectrum cutting-off at a few hundred keV, which is generally interpreted as thermal Comptonisation spectra in a very hot ($\\sim$10$^{9}$ K) optically thin plasma with Thomson depth $\\tau_{\\rm T}\\sim$1 (Sunyaev \\& Titarchuk 1980; {  Haardt \\& Maraschi 1991; Stern et al. 1995b; Poutanen \\& Svensson 1996; Zdziarski et al. 1998)}. The big blue bump observed in the UV spectra of Seyfert galaxies (Walter \\& Fink 1993), as well as the soft X-ray excess in BHBs (e.g. Balucinska-Church et al. 1995) are indicative of thermal emission from optically thick and much colder material with a temperature in the range  10$^{5}$--10$^{7}$ K. The presence of relatively cold material in the innermost part of the accretion flow is corroborated by the existence of reflection features in the X-rays such as the Fe fluorescence line around 6.4 keV and a reflection bump peaking at about 30 keV (George \\& Fabian 1991, Nandra \\& Pounds 1994).  The nature and geometry of these two phases are uncertain. The cold medium is generally believed to form an accretion disc (Shakura \\& Sunyaev 1973), alternatively it could consist of cold dense clouds located inside or around the hot plasma (Rees 1987, Collin-Souffrin et al. 1996, Malzac 2001, Malzac \\& Celotti 2002). The hot phase could constitute the accretion disc corona ( Svensson \\& Zdziarski  1994, \\.Zycki et al. 1995, Sincell \\& Krolik 1997, R{\\'o}{\\.z}a{\\'n}ska \\& Czerny 2000, and ref. therein) or the hot inner part of the accretion disc itself (Shapiro \\& Lightman 1976).  It was soon realized that, if close to each other, the hot and the cold phases should be strongly radiatively coupled. This led to the development of two-phase models (Haardt \\& Maraschi 1991, 1993), where the cold phase constitutes the main source of seed photons for the Comptonisation process. The temperature in the hot phase, and thus the detailed shape of the high energy spectrum, is controlled mainly by the flux of soft cooling photons. On the other hand, the thermal reprocessing of the high energy radiation impinging on the cold phase provides a dominant contribution to soft emission from the cold phase.  Due to the complexity of the Comptonisation process, the importance of geometric effects, and the requirement of taking accurately into account the coupling between the two phases, it was necessary to develop sophisticated numerical codes in order to compute the detailed spectrum produced by the two-phase system in different configurations (Stern et al. 1995a, Poutanen \\& Svensson 1996).  The methods developed in this context, such as the non-linear Monte-Carlo method (Stern et al. 1995a), provide an accurate treatment of the hot phase emission in energy balance. On the other hand, the emission from the cold phase is not detailed: a pure blackbody spectrum with a fixed temperature and reflection on neutral material are generally assumed.  Actually the emission due to X-ray reprocessing in the cold phase differs widely from a simple blackbody, in particular there is a complex line emission. Moreover, due to the hard X-ray irradiation, a ionised skin is likely to form on the surface layers of the cold material (Ross \\& Fabian 1993, Collin-Souffrin et al. 1996, Nayakshin et al. 2000, Ballantyne et al. 2001). This ionised material affects the line emission as well as the shape and amplitude of the reflection component (\\.Zycki et al. 1994, Dumont et al. 2000, Nayakshin et al. 2000, Nayakshin \\& Kallman 2001). Ionization is also likely to affect the equilibrium in the hot phase. Indeed the hard X-ray albedo for ionised material is larger and, as a consequence, the flux of thermal reprocessed soft photons is lower, affecting the temperature in the hot plasma.  The first attempts to include these effects in the two-phase model calculations used a very simplified approach to the computation of the reprocessed spectrum  (e.g. Nayakshin \\& Dove 2001). Indeed, a detailed self-consistent computation of the emission from irradiated material is an heavy task. It requires one to solve the radiative transfer equations in the cold medium taking into account the energy and ionization balance including all the atomic processes. Another difficulty comes from the fact that the cold material is optically thick and its  temperature and ionization structure at equilibrium is strongly inhomogeneous.  The detailed computation of the structure and spectrum of irradiated optically thick material was performed by coupling  the photo-ionization code TITAN  and the Monte-Carlo code NOAR  (accounting for Compton scattering) (Dumont et al. 2000). It enabled one to show the importance of an accurate radiative transfer treatment for the irradiated cold material emission.  However, in these calculations the feedback from the cold radiation on the hot phase equilibrium was not considered. As irradiating spectrum, a simple power-law with a fixed slope was assumed.  Here we add a self-consistent treatment of the hot phase emission by coupling the TITAN/NOAR codes with a third code based on the Non-Linear Monte-Carlo method (Malzac \\& Jourdain 2000).  This enables us to propose the first full treatment of the radiative coupling, accounting accurately and self--consistently for both the hot and cold phase emission.  \\begin{figure}  \\includegraphics[width=\\columnwidth]{snotigeos+s.eps}  \\caption{The infinite slab accretion disc corona and spherical cloud geometry. The hot slab is divided into 10 horizontal layers with homogeneous temperatures. The sphere is divided into 5 concentric shells.} \\label{jmalzac1-C2_fig:fig1} \\end{figure}    \\begin{figure*}[!h]  \\includegraphics[width=\\textwidth]{spxi.eps}  \\caption{Angle averaged escaping spectra for  covering fractions $C=0.3$ (left-hand side) and $C=0.5$ (right-hand side). Thin curves: reprocessed/reflected spectrum. Dashes: comptonised spectrum. Thick curves: total observed spectrum (including the transmitted component). The upper panels show the results from NLMC calculations with the cold phase emission approximated by a blackbody + neutral reflection. The blackbody temperature is $kT_{r}$=10 eV. The other panels are the results from the TITAN/NOAR/NLMC calculations for the following parameters of the cold medium: $n\\,=\\,3\\,10^{14}\\,$cm$^{-3}$, $N_{H}\\,=3\\,10^{25}$\\,cm$^{-2}$ ($\\tau_c=24$). In all of the simulations the Thomson optical depth of the hot plasma was fixed to $\\tau_{\\rm T}=1$. The assumed ionization parameter $\\xi$ is indicated in each panel together with the resulting 2-10 keV photon index of the comptonised emission $\\Gamma$ and the volume averaged temperature of the hot phase $kT_{\\rm e}$ (in keV).} \\label{jmalzac1-C2_fig:fig2} \\end{figure*}  \\begin{figure*}[!h]  \\includegraphics[width=\\textwidth]{albedo.eps}  \\caption{X-ray albedo (i.e. fraction of the reflected/reprocessed luminosity emitted above 1 keV) as a function of the ionization parameter $\\xi$ (left) and spectral index,  $\\Gamma_p$ (right), of the simulated spectra obtained from fits with the PEXRAV model (see text and Table~\\ref{tab:fits}). The results are plotted using a different symbol for each model, as shown on the figure. Filled symbols correspond to the blackbody and neutral reflection models.} \\label{fig:albedo} \\end{figure*}  \\begin{figure*}[!h]  \\includegraphics[width=\\textwidth]{para.eps}  \\caption{Reflection amplitude $R$ and spectral index $\\Gamma_p$ of the simulated spectra obtained from fits with the PEXRAV model (see text and Table~\\ref{tab:fits}). The results are plotted using a different symbol for each model, as shown in the figure. Filled symbols correspond to the blackbody and neutral reflection models. Left panels: $R$ and $\\Gamma$ as a function of $\\xi$. Upper right panel: in the constant pressure case, $R$ and $\\Gamma_p$ as a function of the hot skin Thomson optical depth $\\tau_{\\rm skin}$. Bottom right panel: $R$ versus $\\Gamma_p$, the dashed curve shows the best fit power law approximation to the observed $R$-$\\Gamma$ correlation in AGN as given by Zdziarski et al. 1999. } \\label{fig:fits} \\end{figure*} ",
        "conclusions": "We calculated the equilibrium spectra in the context of the two-phase models for the emission from Seyfert galaxies and black hole binaries. For the first time, we included a detailed treatment of reprocessing and ionization and thermal balance in the cold phase.  The resulting broad band spectrum differs significantly from what obtained using the usual ``blackbody + neutral reflection'' approximation. In particular, our realistic treatment of reprocessing leads to higher hot-phase temperatures associated with  harder X-ray spectra. {  Surprisingly the effect is strong even at low ionization parameters. This is because even if the X-ray albedo is not significantly affected by the weak ionizing illumination, the shape of the reprocessed spectrum below 1 keV is much broader than a simple blackbody.}   For the slab model, we performed simulations with two different prescriptions for the cold phase, namely a constant density and a constant pressure. The comparisons between the two cases indicate that the overall spectral shape is also quite sensitive to the physical structure of the cold phase. A slab model at constant density can account for the hardest observed spectra for large values of the  ionization parameters. In the case of constant pressure disc  however, the hardening is not strong enough to reconcile the model with the even harder spectra observed in Seyfert galaxies and Galactic black holes.  The equilibrium spectra of a two-phase system thus appear to be significantly affected by the shape of the reprocessed spectrum. even at low ionization parameters. A detailed treatment of reprocessing appears to be required for most observationally relevant cases."
    },
    "0410/astro-ph0410291_arXiv.txt": {
        "abstract": "We present here the combined radio spectrum from the Giant Metrewave Radio Telescope (GMRT) and Very Large Array (VLA) of type-IIb SN 1993J (age 11 years) and a type-Ic  SN 2003bg (age 1 year). In SN 1993J, we find a break in the spectrum at 4 GHz and associate it with Synchrotron cooling break.  Hence, we determine the magnetic field independent of equipartition assumption between relativistic electrons and magnetic energy. We also see a hint of break in the spectrum of SN 2003bg between 22-40 GHz. The spectrum of SN 2003bg is well described by the synchrotron self absorption model. ",
        "introduction": "Synchrotron aging can provide wealth of information related to plasma conditions in young supernovae. We discuss two supernovae in this context - an eleven years old type-IIb supernova SN 1993J and a one year old type-Ic  SN 2003bg. ",
        "conclusions": "The relativistic particle energy density is far below than the magnetic energy density (by a factor of 1/10000). This is quite unlike the expectation of efficient shock acceleration of relativistic particles in SNe. The expected fraction of post-shock pressure in relativistic particles is $W \\ge 0.1$ as 10\\% of the energy ejected by the SNe is needed to power the cosmic rays in the galaxy \\citep{mck87}. It indicates that perhaps the efficiency of acceleration is either highly variable in SNe, or it will evolve with time in SN 1993J."
    },
    "0410/astro-ph0410544_arXiv.txt": {
        "abstract": "{Measurement of cosmic microwave background polarization is today a major goal of observational cosmology. The level of the signal to measure, however, makes it very sensitive to various systematic effects. In the case of Planck, which measures polarization by combining data from various detectors, the beam asymmetry can induce a conversion of temperature signals to polarization signals or a polarization mode mixing. In this paper, we investigate this effect using realistic simulated beams and propose a first-order method to correct the polarization power spectra for the induced systematic effect. ",
        "introduction": "After the success of \\emph{COBE} (Smoot et al, \\cite{smoot}) and \\emph{WMAP} (Bennett {\\em et al.} \\cite{Bennett} , the Planck mission, to be launched by ESA in early 2007, is the third generation space mission dedicated to the measurement of the properties of the Cosmic Microwave Background (CMB). About 20 times more sensitive than \\emph{WMAP}, Planck will observe the full sky in the millimeter and sub-millimeter domain in nine frequency channels centered around frequencies ranging from 30 to 70 Ghz (for the Low Frequncy Instrument, or LFI) and from 100 to 850 GHz (for the High Frequency Instrument or HFI). Of these channels, the seven at lowest frequencies --- from 30 to 350 GHz, are polarization sensitive.  Temperature anisotropies have been detected by many experiments now, the most recent of which detect a series of acoustic peaks in the CMB spatial power spectrum (de Bernardis et al, \\cite{boom1}, Hanany et al, \\cite{maxima1}, Beno\\^{i}t et al, \\cite{archeops1}, Hinshaw et al, \\cite{wmapcl}), confirm the Gaussianity of observable CMB fluctuations (Komatsu et al, \\cite{wmapgauss}, though wavelet methods have detected presence of non-Gaussianity in \\emph{WMAP} data, Vielva et al, \\cite{vielva}), and demonstrate the spatial flatness of the Universe (Netterfield et al, \\cite{Netterfield}, Lee et al, \\cite{maxima2}, Beno\\^{i}t et al, \\cite{archeops2}, Spergel et al, \\cite{wmapcp}, Spergel et al, \\cite{wmapcp3}). This provides compelling evidence that the primordial perturbations indeed have been generated during an inflationary period in the very early Universe. The next challenge now is the precise measurement of polarization anisotropies and, in particular, the detection of the pseudo-scalar part of the polarization field (the $B$ modes of CMB polarization) which are expected to carry the unambiguous signature of the energy scale of inflation and of the potential of the inflationary field. The Planck mission will be the first experiment able to constrain significantly these $B$ modes over the full sky and hence to measure them on very large scales.  The first detection of CMB polarization at one degree angular scale of resolution, at a level compatible with predictions of the standard cosmological scenario, has been announced by Kovac {\\em et al} (\\cite{Kovac}) (see also Leitch {\\em et al} \\cite{Leitch}). Since then, CBI and CAPMAP have also obtained significant detection of CMB E--mode polarization (Readhead {\\em et al.}  \\cite{Readhead} and Barkats {\\em et al.} \\cite{Barkats}). More recently, the {\\emph Boomerang} (Piacentini {\\it et al} \\cite{Piacentini}, Montroy {\\it et al} \\cite{Montroy}) and \\emph{WMAP} (Kogut {\\em et al.} \\cite{Kogut}, Page {\\it et al.}  \\cite{Page}) teams have obtained a measurement of the temperature-polarization correlation and E--mode spectrum compatible with cosmological model. No significant constraint on B--mode at degree angular scales exist today.  While the measurement of the temperature and polarization auto and cross power spectra of the CMB carries a wealth of information about cosmological parameters and about scenarios for the generation of the seeds for structure formation, some near--degeneracies exist which require extremely precise measurements. In particular, a very precise control of systematic errors is required to constrain parameters which impact these anisotropies and polarization fields at a very low level.  Many sources of systematic errors are potentially a problem for polarization measurements. In particular, the shape of the beams of the instrument need to be known with extreme precision. In addition, when the measurements of several detectors are combined to obtain polarization signals, it is required that the responses of these detectors be matched precisely in terms of cross--calibration, beam shape, spectral response, etc.  Measurements of Planck telescope beams in the actual operation conditions are not to be made on ground. Also, there are no polarized astrophysical sources for the in-flight beam calibration. Therefore, one should rely on numerical simulations of the beams and self-correcting algorithms of data processing that should allow efficient elimination of systematic errors. In polarization measurements, a significant systematic error would arise due to elliptical shapes of telescope beams which appear, mainly, due to ellipsoidal shape of telescope mirrors introducing astigmatic aberrations and other beam imperfections even with the otherwise ideal mirror surfaces.  In this paper, we investigate the impact of beam imperfections on the measurement of polarization power spectra. We then discuss a method for first-order correction of the effect of these imperfections. To illustrate this method, we apply it to the case of Planck HFI polarization measurement.  The remainder of this paper is organized as follows. In section \\ref{sec:problem}, we discuss the issue of beam shape mismatch for the detection of CMB polarization. Section \\ref{sec:readouts} is dedicated to the computation of simulated readouts using the realistic beams described in Section~\\ref{sec:problem} and the reconstruction of polarized power spectra. In Section \\ref{sec:correction}, we present a method to correct for the systematic bias in the $B$ mode power spectrum induced by the asymmetry of the beams. Finally, Section \\ref{sec:conclusions} draws the conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we have shown the effect of asymmetric telescope beams on the bolometric measurements of polarization of incoming radiation by considering the case of the Planck satellite mission. We have used electromagnetic simulation of the optical system (including telescope and horns) to compute the main beam shapes of the different detectors of Planck. These beams are roughly Gaussian elliptical, with a major axis 10\\% larger than the minor axis and with essentially different orientations of the beam ellipses for the two horns to be combined to measure the full set of Stokes parameters, $I$, $Q$ and $U$.  By simulating the scan of a patch of the sky by Planck with these realistic, simulated beams, we have estimated the bias induced on the $E$ and $B$ mode polarization spectra due to their asymmetric shapes. We first remark that the $E$ mode power spectrum is very well recovered (once corrected for an effective symmetric beam), the bias being around 0.1\\% of the signal in the multipole range $300<l<2000$, where lies the most interesting part of the signal. On the other hand, the $B$ mode is affected by a bias around 10\\% at the peak of the lensing signal ($l\\sim 1000$) and increasing for higher $l$, up to 100\\% of the signal at $l\\sim 2500$. This bias has two origins. First, it is produced by the difference of beam patterns of two different horns combined to measure $Q$ and $U$. This difference induces mainly an error on the polarization angle, which turns to a mixing of $E$ and $B$ modes. Since, in general, $E \\gg B$, we observe finally a leakage from $E$ to $B$. The second origin of the bias is the minor difference of beam patterns of two PSB channels with orthogonal polarizations within the same horn, which induces a temperature to polarization leakage.  Finally, we have proposed a way to correct the $B$ mode power spectrum from the above bias in a one-step correction which uses the measured $T$ and $E$ maps to compute the expected leakage into $B$ when they are observed with a model of the instrument's beams. The efficiency of this correction depends on the precision of the beam knowledge: for example, using elliptical Gaussian fits of the actual beams allows us to reduce the bias from 10\\% to 3\\% at the lensing signal peak, $l\\sim 1000$. In all cases, this first order correction has been shown to reduce significantly $B$ mode contamination. More refined treatments, currently being investigated, are expected to be yet more efficient if needed."
    },
    "0410/astro-ph0410708_arXiv.txt": {
        "abstract": "We present observations of the giant HII region complex N159 in the LMC using IRAC on the {\\it Spitzer Space Telescope}. One of the two objects previously identified as protostars in N159 has an SED consistent with classification as a Class I young stellar object (YSO) and the other is probably a Class I YSO as well, making these two stars the youngest stars known outside the Milky Way. We identify two other sources that may also be Class I YSOs. One component, N159AN, is completely hidden at optical wavelengths, but is very prominent in the infrared. The integrated luminosity of the entire complex is L $\\approx 9\\times10^6$~L$_{\\odot}$, consistent with the observed radio emission assuming a normal Galactic initial mass function (IMF). There is no evidence for a red supergiant population indicative of an older burst of star formation. The N159 complex is 50~pc in diameter, larger in physical size than typical HII regions in the Milky Way with comparable luminosity. We argue that all of the individual components are related in their star formation history. The morphology of the region is consistent with a wind blown bubble $\\approx 1-2$~Myr-old that has initiated star formation now taking place at the rim. Other than its large physical size, star formation in N159 appears to be indistinguishable from star formation in the Milky Way. ",
        "introduction": "The Magellanic Clouds, the nearest moderate sized galaxies to our own Milky Way, fortuitously lie in a direction well out of the plane and are relatively free of foreground extinction. As a result, they present a unique opportunity to study the process of star formation in a galaxy other than our own. N159 is a group of HII regions cataloged by \\citet{hen56} who gave separate letters to distinguish between the clumps of nebulosity found within the complex. The first extragalactic protostar was discovered in the N159 complex \\citep{gat81}, as well as the the first extragalactic Type~I~OH maser \\citep{cas81}, and the region contains an H$_{2}$O maser \\citep{sca81}.  Radio observations of the LMC with the Molonglo Synthesis Telescope \\citep[MOST,][]{mil84} revealed high surface brightness, optically thin, thermal radio emission characteristic of HII regions powered by ionizing photons from OB stars. A higher resolution radio map by \\citet{hun94} confirm the morphology seen in the MOST maps. CO observations \\citep{coh88, bol00, joh98} reveal a very high concentration of molecular gas at the location of N159.  \\citet{gat81} searched N159 at broadband {\\it K} (2.2~\\micron ) looking for very red objects that might be the analog of sources such as BN in the Milky Way. They found such an object, with $K-L^\\prime = +2.7$~mag color and no detectable Br$\\gamma$ emission at 2.16~$\\micron$, which they suggested was a protostar. Subsequently, \\citet{gat82} found a candidate protostellar object in the Small Magellanic Cloud (SMC) while \\citet{jon86} found a second candidate in the N159 complex. Our understanding of YSOs has advanced considerably since these discoveries (for a discussion of low mass YSOs, see \\citet{ada87, shu87} and for high mass YSOs, see \\citet{chu02, gar99}). Currently YSOs are classified in an evolutionary sequence from Class 0 to III. Class 0 stars, the youngest, show no near infrared (near-IR) counterpart and are detected primarily at sub-MM wavelengths. Such objects will be impossible to find in N159 at present because of beam confusion with current far-IR and sub-mm facilities. Class I YSOs, however, typically peak in the mid-IR and have detectable near-IR flux \\citep{shu87, whi03}. A Class I YSO may, or may not, depending on mass and evolutionary stage, be producing ionizing radiation \\citep{chu02}.  The first far-IR detection of N159 was made with the Air Force Geophysics Laboratory (AFGL) survey \\citep{pri76} and a single beam observation from the {\\it Kuiper Airborne Observatory} (KAO) was made by \\citet{wer78}. \\citet{jon86} mapped several HII regions, including N159, using the KAO at 40\\arcsec \\ resolution and found very close correspondence between the far-IR and MOST maps, as expected for dusty HII regions. H$\\alpha$ observations of N159 have revealed locations with high surface brightness that were studied by \\citet{hey82} in N159A and by \\citet{hey99} in N159-5 (The Papillon). The stellar content associated with N159A was studied at optical wavelengths by \\citet{deh92} while the stellar content around the Papillon was studied at near-IR ({\\it JHK}) wavelengths by \\citet{mey04}.  There are three questions regarding star formation in N159 we wish to explore using Infrared Array Camera \\citep[IRAC,][]{fazio04} images obtained with the {\\it Spitzer Space Telescope} \\citep[{\\it Spitzer},][]{wer04}. First is the nature of the two protostars discovered by \\citet{gat81} and \\citet{jon86}. In light of our new measurements and a more contemporary understanding of very young stars, how should the two protostars in N159 be classified? Is there any evidence for additional protostars missed in the earlier near-IR surveys?  A second question we wish to address is the nature of the upper end of the Initial Mass Function (IMF) in the LMC. Based on far-infrared (far-IR) observations of giant LMC HII regions, \\citet{jon86} suggested N159 was producing more ionizing photons for its luminosity than it should, given their estimate of the total luminosity. Perhaps the IMF was weighted to producing more very massive stars than is the case in the Milky Way. Subsequent studies in the optical, in particular \\citet{mas98}, argue that the LMC has a normal IMF at the extreme high mass end. Our {\\it Spitzer} observations fill in a portion of the spectral energy distribution (SED) missing from the earlier study by \\citet{jon86} enabling a more accurate assessment of N159's luminosity.  Thirdly, we seek to understand the morphology of N159, which is larger in physical size ($\\approx 50$~pc) than typical for Galactic HII regions. Since there is a significant reservoir of molecular gas in the vicinity of N159, the individual components of the complex may be unrelated and simply the product of randomly initiated star formation at locations of dense molecular gas. Alternatively, the morphology and stellar content may suggest a scenario where the star formation within N159 was triggered by a single, relatively recent event less than 2~Myr ago.  In this paper we present {\\it Spitzer} images of N159 using the four channels of IRAC with band centers at $3.5, \\ 4.5, \\ 5.7,$ and $7.9~\\micron$. In \\S2 we describe the observations and data reduction. In \\S3 we examine the general morphology of the nebula and make comparisons to observations at other wavelengths. In \\S4 we study the point sources seen in the IRAC maps, including the two protostars. Finally, we discuss the global properties and evolutionary state of the N159 complex and present a scenario for wind driven star formation on 50~pc scales. ",
        "conclusions": "}  Based on our new {\\it Spitzer} observations of the protostars in N159, we classify P2 as a Class~I YSO with an age of $\\approx 10^{5}$~yrs and a luminosity of $3 - 5~\\times~10^{3}$~L$_{\\odot}$. The original protostar, P1 is also likely a Class~I YSO (or perhaps a Class~II YSO), but its [5.7]-[7.9] color is not as red as typical Milky Way Class~I YSOs. Two new sources have the colors of Class~I YSOs and deserve attention in future observations.  The total luminosity of the entire N159 complex is $\\approx 9~\\times~10^{6}$~L$_{\\odot}$. The ratio of total radio flux density to total luminosity is consistent with typical dusty HII regions observed in the Milky Way. There is no evidence for an excess of ionizing photons indicative of a top heavy IMF. There is no evidence of a previous episode of significant star formation that would have produced core burning supergiants that would still be visible today.  Star formation in N159 appears identical to star formation in the Milky Way, in particular an association with molecular gas and the presence of relatively high mass protostars. N159 is somewhat different, however, in its large physical size. The entire complex is over 50pc in diameter with at least four separate concentrations of ionizing flux and mid-IR luminosity. We argue that all of these separate HII regions are related to one another, the result of triggered star formation on the scale of tens of parsecs."
    },
    "0410/astro-ph0410222_arXiv.txt": {
        "abstract": "Outflows  up to 1500 \\kms\\ of atomic neutral hydrogen are detected in a growing number of radio galaxies. Outflows with similar velocities are also detected in ionized gas, suggesting a common origin for the extreme kinematics of these two phases of the gas.  The high detection rate of such outflows  in young (or restarted) radio sources appears to be related to the existence of a dense ISM around these objects.  Such a dense ISM can have important consequences for the evolution of the radio source and the galaxy as a whole.  Here we summarize the recent results obtained and the characteristics derived so far for these outflows.  We also discuss possible mechanisms (e.g. interaction between the radio plasma and the ISM and adiabatically expanding broad emission lines clouds) that can be at the origin of these phenomena. ",
        "introduction": "The presence of neutral hydrogen in the central regions of active galaxies is often explained as being due to gas in circumnuclear tori/disks, i.e. with gas relatively settled and regularly rotating (see e.g. Pihlstr\\\"om, Conway, Vermeulen 2003, Mundell et al. 2003 and Morganti et al. 2004a for a review).  However, the conditions of the \\HI\\ are often more complex and disturbed  kinematics of the atomic neutral hydrogen is observed in a growing number of radio sources.  The presence of this ``unsettled'' gas can now be detected thanks to the possibility of performing, with sensitive systems, \\HI-21 cm observations using broad spectral bands (e.g. bands that can cover up to $\\sim 4000$ \\kms\\ in the case of the recently upgraded WSRT).  The presence of gas with extreme kinematics in the central regions of active galactic nuclei (AGN) has recently attracted particular attention, in particular on the presence and occurrence of fast gas outflows.  Fast nuclear outflows of {\\sl ionized} gas appear to be a relatively common phenomena in active galactic nuclei (see e.g. Crenshaw et al.\\ 2000, Kriss et al. 2004, Capetti et al. 1999, Krongold et al.\\ 2003, Veilleux et al.  2002, Tadhunter et al. 2001 Elvis 2000 and ref. therein).  They are mainly observed in optical, UV and X-ray observations.  Gas outflows associated with AGN provide energy feedback into the interstellar medium (ISM) that can profoundly affect the evolution of the central engine as well as that of the host galaxy. The mass-loss rate from these outflows can be a substantial fraction of the accretion rate needed to power the AGN.  Thus, the physics of these phenomena needs to be understood in order to understand AGN. It is not too surprising to find such outflows also in radio galaxies (see Tadhunter et al. 2001, Holt et al. 2003 and Morganti et al. 2003a for a summary of recent results).  However, it is extremely intriguing that is several radio sources fast outflows of {\\sl neutral} hydrogen ( up to 2000 \\kms) have been discovered.  The physical conditions of this gas provide new and important clues on the mechanism responsible for AGN-related outflows. Here, we summarize the cases found so far and the characteristics of the neutral hydrogen associated with such outflows. We also discuss some of the possible mechanisms that can be at the origin of this phenomenon.  \\begin{figure*} \\centerline{\\psfig{figure=morganti_fig1.ps,angle=0,width=16.0cm} } \\caption{Examples of broad \\HI\\ absorption has been detected in radio galaxies.  The systemic velocity as derived from optical emission lines is also indicated. In two cases (3C~459 and PKS~1814-637) we have marked both the systemic velocity (corresponding to the velocity of the {\\sl narrow} lines) as well as the velocity of the {\\sl broad} component also detected in the optical emission lines (Holt et al. in prep.). See text for details.} \\end{figure*} ",
        "conclusions": ""
    },
    "0410/gr-qc0410053_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cosmic magnetic fields are ubiquitous at all large intragalactic scales. It is a well-known observational fact that our galaxy and many other spiral galaxies are endowed with coherent magnetic fields of $\\mu $G (microgauss) strength \\cite{Sofu, Asseo, Kronberg, Beck, Grasso}, having approximately the same energy density as the cosmic microwave background radiation (CMBR). There is also evidence for larger magnetic fields of similar strength within clusters \\cite{Perely, Kim}. The presence of magnetic fields at larger scales has also been confirmed \\cite{Giovannini,Vallee}. These magnetic fields play an important role in various astrophysical processes, such as the confinement of cosmic rays and the transfer of angular momentum away from protostellar clouds so that they can collapse and become stars. Magnetic fields are also present in the intracluster gas of rich clusters of galaxies, in quasistellar objects (QSO's) and in active galactic nuclei. They may influence the formation process of large-scale structure \\cite% {Subramanian,Tsagas}.  It is widely believed that galactic magnetic fields are amplified and sustained by a dynamo mechanism \\cite% {Kronberg,Beck,Grasso,Parker1,Parker2,Zel'dovich}, in which the cyclonic turbulent motion of ionized gas combined with the differential rotation of the galaxy exponentially amplifies a ``seed'' magnetic field. This continues until the backreaction of the motion of the plasma offsets the growth of the field, stabilizing it to dynamical equipartition strength. However, while the dynamo mechanism provides an amplification mechanism, it does not explain the origin of galactic magnetic fields, and requires a ``coherent'' seed magnetic field for it to be effective. Indeed, it has been shown that seed magnetic fields that are too incoherent may undermine the action of the dynamo \\cite{Kulsrud}. Most dynamo scenarios require a minimum coherence length equal to the dimension of the largest turbulent eddy, usually around $% \\sim 100$ pc. If the mechanism has functioned over the whole age of the galaxy($\\sim $ 10G yr) a seed field of $10^{-19}$G is required. If recent observations are correct and the universe is dominated by a dark-energy density component \\cite{Supernova, Jaffe}, then galaxies are older than previously thought and the seed magnetic field may be as low as $% B_{seed}\\sim {10}^{-30}$G \\cite{Davis}.  A contrasting view is that the primeval magnetic flux trapped in the gas that collapsed to form the galaxy is responsible for the existence of galactic magnetic fields. This hypothesis also requires the existence of a seed magnetic field, one that is as great as the field observed today \\cite% {Piddington, Ohki}. Several scenarios have been suggested for creation of the required seed magnetic field, the most important of which involve battery \\cite{Kulsrud} or vorticity \\cite{Harrison} effects. The battery mechanism requires a large-scale misalignment of density and pressure gradients usually related to active galactic nuclei (AGN) or starburst activity. Therefore, it is difficult to realize in the majority of galaxies. The vorticity mechanism is based on the relative motions of photons and electrons induced by vorticity that was present before decoupling. Of course this mechanism assumes the existence of primeval vorticity. In addition, large-scale vortical motions can be effective only if ionization of the plasma is considerable, which does not occur at the galaxy-formation epoch.  Throughout most of the history of the universe the average time $\\tau $% \\thinspace between particle interactions has been much smaller than the expansion time scale, $\\tau \\ll t_{Hubble}$ . Consequently the universe has been a good conductor \\cite{Turner1}, and any primeval cosmic magnetic field would have evolved in a manner that preserved magnetic flux: $Ba^{2}\\sim $ constant, where $a$ is the scale factor. Hence the dimensionless ratio $% r=B^{2}/\\left( 8\\pi \\rho _{\\gamma }\\right) $ is almost constant and provides a convenient measure of magnetic field strength. If there had been a pregalactic cosmic magnetic field that collapsed with the gas that formed the galaxy, its strength must have increased as $\\mathbf{[\\rho }_{gal}% \\mathbf{/\\rho }_{tot}(t)]^{2/3}$, where $\\rho _{tot}(t)$ is the average cosmic mass density at time $t$. As $\\rho _{tot}\\propto a^{-3}$ and $\\rho _{gal}/\\rho _{tot}=10^{6}$ today ($t_{0}=0$), it follows that the strength of the magnetic field at the time of formation $t_{form.}$ must have been $% 10^{4}[a(t_{form.})/a(t_{0})]^{2}B_{cosmic}$ or $B_{gal.}\\simeq 3r^{1/2}10^{-2}$G. This yields $r\\simeq 10^{-34}$ for initiating the galactic dynamo, or alternatively $r\\simeq 10^{-8}$ for seeding the galactic magnetic field itself while avoiding the necessity of a galactic dynamo. If the existence of dark energy in the universe is confirmed, the minimum $r$ required to seed the dynamo mechanism reduces to $10^{-56}$.  Inflation offers the hope of furnishing a mechanism for kinematically and dynamically producing the seed for cosmic magnetic fields. It provides the kinematic means for producing long-wave-length effects at very early times through microphysical processes operating on scales less than the Hubble radius. Since an electromagnetic wave with $\\lambda _{phys}\\geq H^{-1}$ has the appearance of static $\\mathbf{E}$ and $\\mathbf{B}$ fields, very long wavelength photons ($\\lambda _{phys}\\gg H^{-1}$) can lead to large-scale magnetic fields (which then become supported by currents). Of course the electric field generated during the inflationary stage not only is not amplified, but is actually damped down due to the large conductivity of the primeval plasma. Another reason inflation is considered to be a prime candidate for field amplification is the fact that during inflation the universe is devoid of charged particles. Hence, the magnetic flux is not necessarily conserved and $r$ can increase. Furthermore, inflation can superadiabatically amplify the energy density($\\simeq kd\\rho /dk$) of the minimally coupled field \\cite{Turner2}. Then the energy density decays as $% a^{-2}$, rather than the usual result $a^{-4}$(``adiabatic result'').  However the conformal flatness of the Robertston-Walker metric prevents the background gravitational field from producing particles, provided the underlying matter theory is conformally invariant. A pure U(1) gauge theory with the standard lagrangian $\\mathcal{L}=-\\frac{1}{4}F_{\\mu \\nu }F^{\\mu \\nu }$ is conformally invariant, from which it follows that $\\rho _{B}\\propto B^{2}$ always decreases as $a^{-4}$. During the inflationary epoch, the total energy density in the universe is dominated by vacuum energy and therefore the energy density in any magnetic field during inflation is significantly suppressed. In fact it can be shown that $r=10^{-104}\\lambda _{Mpc}$ \\cite{Turner1}, where $\\lambda_{Mpc}\\equiv \\lambda/1 Mpc$.  Several proposals have been given to break the conformal invariance of the theory: (i) coupling the electromagnetic field to a non-conformally-covariant charged field \\cite{Ratra}, (ii) coupling the electromagnetic field to gravity via either gauge non-invariant terms such as $RA_{\\mu }A^{\\mu }$, $R_{\\mu \\nu }A^{\\mu }A^{\\nu }$ or gauge invariant ones like $R_{\\mu \\nu \\lambda \\kappa }F^{\\mu \\nu }F^{\\lambda \\kappa }/m^{2}$% , $R_{\\mu \\nu }F^{\\mu \\kappa }F_{\\kappa }{}^{\\nu }/m^{2}$ or $RF^{\\mu \\nu }F_{\\mu \\nu }/m^{2}$ \\cite{Turner1}, (iii) invoking effects due to the quantum conformal anomaly \\cite{Dolgov,Chanowitz}, (iv) creating primordial magnetic fields at either the QCD transition epoch \\cite{Quashnock} or the electroweak transition \\cite{Vachaspati}, (v) breaking conformal invariance via nonzero vacuum expectation values of flat directions in minimally supersymmetric standard models (MSSM) \\cite{Enqvist}.  Attempts to realize the first possibility were carried out by coupling the electromagnetic field to the scalar field $\\Phi $ responsible for inflation via a term $\\propto e^{\\Phi }F^{\\mu \\nu }F_{\\mu \\nu }$ \\cite{Ratra}. This investigation showed that in the exponential potential for the inflaton% \\begin{equation} V(\\Phi )=(\\frac{6-q}{3})\\frac{16\\pi }{{m_{pl}}^{2}}\\rho _{b\\Phi }^{0}e^{(-% \\sqrt{\\frac{q}{2}}(\\Phi -\\Phi _{0}))},  \\label{1} \\end{equation}% it is possible to generate an intergalactic magnetic field whose present strength (depending on values of parameters of the model) lies between $% 10^{-65}$ to $10^{-10}$ on a scale of $1/1000$ that of the Hubble scale. In (% \\ref{1}) $m_{Pl}$ is the planck mass, $\\Phi ^{0}$ is the value of the scalar field and $\\rho _{b\\Phi }^{0}$ is the homogeneous scalar field energy density when the scale factor is $a_{0}$. Thus in this scenario one can have the desired galactic magnetic field by resorting to the dynamo mechanism.  Considering next possibility (ii), if we add gauge non-invariant terms to the action, the $U(1)$ gauge invariance will be broken. To avoid the phenomenological disasters this can cause one can endow the photon with a mass squared of the order of $H^{2}$ (well below present limits of detectability). It has been shown \\cite{Turner1} that such a term can create primeval fields with strength as large as $r\\sim 10^{-8}$ .  Gauge invariant modifications to the action have much better theoretical motivation. For example, all $RF^{2}$ terms can be obtained by calculating the effective Lagrangian for QED in curved space-time to one loop order \\cite% {Drummond}. At early times, when $R^{1/2}\\sim H\\sim \\rho _{tot}^{1/2}/m_{pl}\\gg 10^{-11}m_{pl}$, these terms govern the behavior of the electromagnetic field. However at later times, when $R^{1/2}\\ll 10^{-11}m_{pl}$, they are negligible compared to the standard $-\\frac{1}{4}% F^{2}$ term. In a power-law inflationary background the amplitude of large-scale fields is not large enough to be astrophysically interesting % \\cite{Turner1}.  The third possibility has proven to be promising for gauge theories with large groups and a greater number of bosons than fermions. In such theories, it has been shown that this mechanism for breaking conformal invariance in quantum electrodynamics can create a sufficient amount of primordial magnetic field \\cite{Dolgov}. However, in the simplest version of the grand unified $SU(5)$ model with three generations of fermions the magnetic field produced is below the requirement of the dynamo mechanism.  Phase transitions at different cosmological epochs (grand unification, the electroweak transition \\cite{Vachaspati} or the quark confinement epoch \\cite% {Quashnock}) have also been considered. However, since the generating mechanisms are causal, the coherence of the created magnetic field cannot be larger than the particle horizon at the time of the phase transition. Because all the above transitions occurred very early in the universe's history, the comoving size of the horizon is rather small. The best case is the QCD transition, for which the horizon corresponds to $\\sim $ 1 a.u. Consequently the real magnetic fields generated lack sufficient coherence.  The MSSM flat directions, made up of gauge invariant combinations of squarks and sleptons, acquire non-vanishing vacuum expectation values (vev) during inflation. These flat directions endow the standard model gauge fields with mass and break the conformal invariance. The quantum fluctuations of these flat directions, in contrast to the their classical vevs, induce fluctuations in the gauge degrees of the freedom that cannot be gauged away. The gauge field fluctuations that are stretched outside the horizon during inflation, provide us with a seed (hyper)magnetic field after they re-enter the horizon. They give rise to $U(1)_{em}$ magnetic field with strength of $% 10^{-30}$ G, as required by the dynamo mechanism \\cite{Enqvist}.  Here we consider an alternative mechanism that is based on a hypothesis of a minimal fundamental length scale. Minimal length breaks conformal invariance and so it might be expected that primordial magnetic fields can be produced during inflation. One suggestion \\cite{Kempf} for implementing minimal length into the inflationary scenario in the context of trans-planckian physics \\cite{Martin} is based on the hypothesis of a generalized uncertainty principle: \\begin{equation} \\Delta x\\Delta p\\geq \\frac{1}{2}\\left( 1+\\beta ~(\\Delta p)^{2}\\right) , \\label{mur} \\end{equation}% where $\\sqrt{\\beta }$ is the ultraviolet cutoff on the order of the Planck or string length. In this paper we employ this formalism to implement minimal length into the action of electrodynamics. This translates into a UV cutoff which, once implemented, has the sole effect of modifying the evolution of the electromagnetic field. As we will demonstrate, the formalism is not able to create a squeezing effect for the electromagnetic field. Therefore the energy density of the electromagnetic field attenuates adiabatically, $\\rho _{B}\\propto a^{-4}$.  However, it has recently been shown \\cite{Ashoorioon} that terms in the action that are total time derivatives are not invariant under the influence of the minimal length hypothesis \\cite{Kempf}. Consequently such terms contribute to the equations of motion of the matter fields. We consider in this paper an example of a total time derivative that, under the influence of the UV cutoff, causes the photon to gain a large negative mass during inflation. This effect goes to zero at the end of inflation and so such a mass is undetectable today. We shall show that this approach is successful in providing the dynamo mechanism with sufficient primordial seed magnetic field. Even in absence of the dynamo mechanism, one can adjust a free parameter in the action to account for the observed magnetic field of galaxies today. ",
        "conclusions": "The origin of magnetic fields with $\\mu $G strength that are observed on intragalactic scales remains an intriguing mystery. As the observed magnetic field is coherent on such cosmological scales, the first cosmological process one might think of as being able to produce such prevalent fields is inflation. However the conformal invariance of the electromagnetic field prohibits the quantum fluctuations of the electromagnetic field from squeezing and amplifying during inflation.  The existence of a minimal length breaks this conformal invariance. We have proposed a scenario based on this observation that can provide the requisite initial magnetic seed for the astrophysical dynamo mechanism. With a proper choice of the free parameter within the theory one can avoid the need for the dynamo mechanism.  The scenario is based on the observation that incorporating minimal length at the level of first quantization, as was done for the first time in \\cite% {Kempf}, does not render total time derivatives invariant under the influence of minimal length. Therefore one can have actions that are equivalent at the continuous space-time level, but are distinct from one another once the presence of minimal length is introduced. We added a prototype for such a total time derivative term to the action of electromagnetism that respects the behavior of the photon throughout the history of the universe except for the inflationary era. During inflation this term induces a huge mass for the photon. We found that to match this model with observation we must tune the free parameter of the model, $M_{1}$% , to be extremely small. Since $M_{1}^{-2}$ is proportional to the coupling of electromagnetism to the background geometry during inflationary epoch, the small size of $M_{1}$ is indicative of gravity and electromagnetism being strongly coupled at that time. We note that the numerical value of $% M_{1}$ is approximately the inverse Hubble length, but we have found no deeper explanation for this coincidence at the level of the model presented here.  Of course it is conceivable that other gauge bosons of the standard model can inherit the same tachyonic instability that we have considered for the photon. \\ However all other gauge bosons are non-Abelian and so will experience screening effects that we expect will tend to dampen out this instability \\cite{Biro}. \\ A detailed calculation of this effect remains an interesting subject for future study.  Of course the main drawback for this model is its arbitrariness in the choice of total time derivative. It would be really interesting if one were able to find candidates from existing models of fundamental physics. Our main goal here was that of demonstrating that specific characteristics of space-time at Planckian epochs can create observable phenomena in the universe at much later cosmological times."
    },
    "0410/astro-ph0410014_arXiv.txt": {
        "abstract": "CG J1720-67.8 is an ultra compact group of several galaxies with a low velocity dispersion, and displaying the hallmarks of mutual interaction and possible tidal dwarf galaxy formation. In hierarchical models, the system is a possible precursor to a massive elliptical galaxy.  In this paper, we use new optical integral field spectroscopic and radio observations to investigate the evolutionary status of the group in more detail: global star-formation rates are estimated using H$\\alpha$ and 1.4 GHz radio continuum measurements; H\\,{\\sc i} observations provide an upper limit to the global neutral gas content; optical broadband colours and spectra provide ages and stellar mass estimates for the tidal dwarf candidates; the bidimensional H$\\alpha$ velocity field is used to trace the kinematics of the group and its members, which are compared with numerical simulations of galaxy encounters. The observations suggest a model in which multiple interactions have occurred, with the latest strong encounter involving at least two major components within the last 200 Myr. Debris from the encounter fuels ongoing star formation at the global level of $\\sim20$ M$_{\\odot}$ yr$^{-1}$, with self-gravity within the tidal tail possibly providing a mechanism to enhance the star formation rate of the tidal dwarf candidates, with bursts of star-formation in clumps of mass $\\sim2\\times 10^7$ M$_{\\odot}$ appearing within the last 10 Myr. The amount of time required for final merging of all group components remains uncertain. ",
        "introduction": "One of the most important problems in cosmology concerns the formation and evolution of galaxies. Sixty per cent of galaxies are found in groups \\citep{tu87}. This fact alone indicates that the group environment is of utmost importance for galaxy evolution. The most recent views of hierarchical structure formation in the Universe involve galaxy groups as building blocks of clusters and as places where interactions pre-process galaxies, which eventually become part of the cluster population \\citep[e.g. ][]{mi04}. Due to the high densities and low velocity dispersions (with respect to clusters) that characterize them, compact groups (CGs) of galaxies offer the most favorable environmental conditions for interaction and merging processes to take place and lead to profound galaxy transformations. The evolution of CGs is believed to end in bright elliptical galaxy formation through subsequent mergers \\citep[e.g., ][]{kc98,ba98}. Therefore, groups can be regarded as putative factories of field elliptical galaxies. This evolutionary scenario is supported by the recent identification of ``fossil groups'' or ``over-luminous elliptical galaxies'' in the X-ray domain. More than 10 fossil groups are known to date, although only a few of them have been studied in detail \\citep[see e.g., ][and references therein]{s04}. The most likely interpretation of these fossil systems is that they are the end-results of galaxy merging within normal groups \\citep{j03}.  However, despite considerable efforts, details of the processes leading from CGs to elliptical galaxies are still uncertain and, particularly, the latest stages of evolution of CGs are profoundly unknown. In fact only very few high density (i.e. evolved) CGs have been found to date. The most well studied sample of CGs is that compiled by \\citet{hi82} which encompasses groups in differing evolutionary stages. However, none of these systems can be said to be close to final merging, with the possible exception of HCG 31 \\citep{ler04} and HCG 79 which is estimated to achieve final coalescence in less than 1 Gyr \\citep{ni00}. The rarity of highly evolved CGs suggests that, once triggered, the coalescence phase must be relatively short. On the other hand, the fraction of fossil groups with respect to normal groups would point to a relatively slow rate of formation of fossil groups, consistent with the long dynamical friction time-scale of L$^{\\ast}$ galaxies \\citep{j03}. Hence, apart from the importance of understanding the physical processes prevailing in transition objects, it is of great significance to find and study galaxy groups that show characteristics of a very advanced phase of evolution such as extremely high density and low velocity dispersion.  In recent work \\citep{t03a,t03b}, we have presented a photometric and spectroscopic analysis of the ultra-compact galaxy group CG J1720-67.8 \\citep{wtk99}. The group, characterized by a very high density and the presence of an outstanding tidal tail hosting a number of tidal dwarf galaxy (TDG) candidates, reveals properties indicative of an advanced evolutionary stage. Interpreting its properties in the framework of hierarchical clustering theories would suggest that the group is close to evolving into to a single merged galaxy, although the evolutionary history of this interesting galaxy group is not yet established.  In this paper we present new observations of CG J1720-67.4 both in the optical and radio domain and use them in combination with previously obtained data to advance our understanding of the evolutionary state of the group. In particular, a mosaic of integral field spectra covering virtually the whole extent of the group is used to investigate the group kinematics. New long-slit spectra of the central/southern part of the tidal tail, obtained at the 6.5-m Magellan telescope, serves to further investigate the kinematics of some tidal dwarf galaxy (TDG) candidates. The star formation activity of the group is studied using new H$\\alpha$ and radio-continuum data. H\\,{\\sc i} observations are used to obtain information on the mass of neutral gas in the group, an additional clue to the evolutionary state of the system. Finally, we investigate the properties of the tidal features by means of $BVR$ photometry.  For a description of data obtained at the ESO 3.6 m telescope with EFOSC2 in the $BVR$ bands and their reduction and photometric calibration, we refer the reader to \\citet{t03a}. Our observations and analysis of integral field data, long-slit spectra and radio data are described in Sections 2, 3, and 4, respectively, where results from the individual observations are also reported. Our observational results are used to discuss the star formation activity (Section 5), the structure of the diffuse light and tidal features (Section 6) and the group's kinematics and dynamics in the framework of its interaction history (Section 6). Our main results are summarised in Section 7. ",
        "conclusions": "In this paper and our previous works \\citep{wtk99,t03a,t03b} we have collected a series of multiwavelength observations with the main aim of establishing the evolutionary state of the ultra-compact galaxy group CG J1720-67.8. Our studies have revealed this system as an interesting candidate to represent the rarely observed phase in the evolution of compact groups that precedes the final merging. In particular in this work we have considered the kinematics of the group components -- based on both integral field and long slit optical spectroscopy -- and the intensity and distribution of star formation activity -- based on H$\\alpha$ and radio continuum observations -- together with the photometric properties of the whole group and its tidal debris to put some constraints on the interaction history of CG J1720-67.8. We use encounter simulations in the literature as well as examples of merging galaxy pairs to interpret our observations and outline some possible evolutionary scenarios. However, the presence of tails and bridges involving all three galaxies suggests a more complex interaction than those usually simulated considering only galaxy pairs. In CG J1720-67.8 we have most probably a three-body encounter or an even more complicated case with one of the galaxies being already the outcome of a relatively recent merger.  Although other interpretations are possible, as discussed extensively in Section 7.1, we favour a scenario in which Galaxy 4 and Galaxy 1 experienced a prograde-retrograde close encounter within the last 200 Myr. During this encounter, the strong tidal tail and the bridge between the two galaxies were formed, and possibly also the cone-like plume with the apex pointing approximately to the centre of Galaxy 1. The interaction process triggered strong star formation episodes across the bodies and in the centres of both galaxies. More recent star formation events ($<$ 10 Myr) have taken place in the condensations of gas and stars that formed under the action of self-gravity within the tidal tail. These condensations have properties consistent with tidal dwarf galaxies in the process of formation. Although these objects cannot be confirmed as real TDGs, they do show possible signs of self-gravitation.  The role of Galaxy 2 in this recent encounter is not clear. The small bridge departing from Galaxy 1 in its direction suggests that this galaxy was involved in the interaction. The presence of a further bridge between Galaxies 2 and 4 is possible but cannot be established with certainty because of the significant overlap between the two galaxies in projection on the sky. Our new data have confirmed the presence of low level star formation concentrated in the central part of Galaxy 2. This fact, together with the surface brightness distribution indicating a dominant de Vaucouleurs bulge, suggested the interesting hypothesis that Galaxy 2 is a merger remnant still showing the remnant of a central burst of star formation triggered by the merger event. If this was the case, the faint secondary tail seen on the western side of the group could have emanated from Galaxy 2's progenitors. This is also an alternative origin for the northern cone-like plume. Since our data do not have sufficient spatial resolution to disentangle any possible double nucleus or other substructure within Galaxy 2, the merger hypothesis remains speculative. The proximity of the three galaxies in space and their strong interactions suggests that they will merge in a relatively short time. Furthermore, several indications of significant  H\\,{\\sc i} deficiency in the group (Section 7.3), according to the evolutionary scenario proposed by \\citet{vm}, support the idea that the galaxy system is already evolved. Dedicated dynamical modelling is necessary to evaluate the validity of any of the proposed scenarios and predict the future evolution of the system.  \\paragraph*"
    },
    "0410/astro-ph0410364_arXiv.txt": {
        "abstract": "Coalescence of binary supermassive black holes (SBHs) would constitute the strongest sources of gravitational waves to be observed by LISA. While the formation of binary SBHs during galaxy mergers is almost inevitable, coalescence requires that the separation between binary components first drop by a few orders of magnitude, due presumably to interaction of the binary with stars and gas in a galactic nucleus. This article reviews the observational evidence for binary SBHs and discusses how they would evolve. No completely convincing case of a bound, binary SBH has yet been found, although a handful of systems (e.g. interacting galaxies; remnants of galaxy mergers) are now believed to contain two SBHs at projected separations of $\\lap 1$ kpc. $N$-body studies of binary evolution in gas-free galaxies have reached large enough particle numbers to reproduce the slow, ``diffusive'' refilling of the binary's loss cone that is believed to characterize binary evolution in real galactic nuclei. While some of the results of these simulations -- e.g. the binary hardening rate and eccentricity evolution -- are strongly $N$-dependent, others -- e.g. the ``damage'' inflicted by the binary on the nucleus -- are not. Luminous early-type galaxies often exhibit depleted cores with masses of $\\sim 1-2$ times the mass of their nuclear SBHs, consistent with the  predictions of the binary model. Studies of the interaction of massive binaries with gas are still in their infancy, although much progress is expected in the near future. Binary coalescence has a large influence on the spins of SBHs, even for mass ratios as extreme as $10:1$, and evidence of spin-flips may have been observed. ",
        "introduction": "\\label{sec:intro}  With an ever-increasing number of secure detections, supermassive black holes (SBHs) have evolved, in the span of a few years, from exotic possibilities to well-established components of galaxies. While it was understood since the 1960's that the energy sources of quasars must be gravitational \\cite{robinson-65}, it was thirty years before the existence of SBHs was firmly established, through measurements of the Keplerian rise in the rotation velocity of stars or gas at the very centers of galactic nuclei \\cite{kr-95}. It is now generally accepted that the formation and evolution of galaxies and SBHs are tightly intertwined, from the early phases of proto-galactic formation \\cite{silk-98}, through hierarchical build-up in CDM-like cosmogonies \\cite{haehnelt-02}, to recent galaxy mergers \\cite{mm-01}.  SBHs appear to be linked in fundamental ways to the dynamics of the stellar component in galaxies, both on large and small scales. An astonishingly tight correlation exists between SBH mass and the central velocity dispersion of the stellar component, $M_{\\bullet}\\sim\\sigma^{\\alpha}$, $\\alpha\\approx 4.5$ \\cite{fm-00}; the correlation with the velocity dispersion averaged over kiloparsec scales is weaker but still impressive \\cite{gebhardt-00,ferrarese-02}. Similar correlations exist between SBH mass and bulge luminosity \\cite{mclure-02,marconi-03} and central concentration of the light \\cite{graham-01,erwin-04}, indicating that SBHs ``know'' about the depth of the gravitational potential well in which they live. These tight correlations probably reflect a degree of feedback in the growth of SBHs \\cite{silk-98}.  On small scales, SBHs are embedded in stellar cusps, parsec-scale regions where the stellar density increases approximately as a power law with distance from the SBH into the smallest resolvable radii \\cite{crane-93,ferrarese-94,mf-95,gebhardt-96}. Faint galaxies have steep nuclear density profiles, $\\rho\\sim r^{-\\gamma}$, $1.5\\lap\\gamma\\lap 2.5$, while bright galaxies typically have weaker cusps, $\\gamma\\lap 1$. Steep cusps form naturally as the growth of the SBH pulls in stars \\cite{peebles-72}. In small dense galaxies where the star-star relaxation time is shorter than $10^{10}$ yr, steep cusps may also form via collisional relaxation \\cite{bw-76,preto-04}. Weak cusps may be remnants of strong cusps that were destroyed by binary SBHs during galaxy mergers; in fact the structure and kinematics of galactic nuclei are now believed to be fossil relics of the merger process \\cite{merritt-04}.  Larger galaxies grow through the agglomeration of smaller galaxies and protogalactic fragments. If more than one of the fragments contained a SBH, the SBHs will form a bound system in the merger product \\cite{bbr-80,roos-81}. This scenario has received considerable attention because the ultimate coalescence of such a pair would generate an observable outburst of gravitational waves \\cite{thorne-76}. The evolution of a binary SBH can be divided into three phases \\cite{bbr-80}: 1. As the galaxies merge, the SBHs sink toward the center of the new galaxy via dynamical friction where they form a binary. 2. The binary continues to decay via gravitational slingshot interactions \\cite{saslaw-74} in which stars on orbits intersecting the binary are ejected at velocities comparable to the binary's orbital velocity, while the binary's binding energy increases. 3. If the binary's separation decreases to the point where the emission of gravitational waves becomes efficient at carrying away the last remaining angular momentum, the SBHs coalesce rapidly.  The transition from (2) to (3) is understood to be the bottleneck of a SBH binary's path to coalescence, since the binary will quickly eject all stars on intersecting orbits, thus cutting off the supply of stars. This is called the ``final parsec problem'' \\cite{mm-03a}. But there are other possible ways of continuing to extract energy and angular momentum from a binary SBH, including accretion of gas onto the binary system \\cite{armitage-02} or refilling of the loss cone via star-star encounters \\cite{yu-02,mm-03b} or triaxial distortions \\cite{poon-04b}. Furthermore there is circumstantial evidence that efficient coalescence is the norm. The $X$-shaped radio sources \\cite{dennett-02} are probably galaxies in which SBHs have recently coalesced, causing jet directions to flip. The inferred production rate of the $X$-sources is comparable to the expected merger rate of bright ellipticals, suggesting that coalescence occurs relatively quickly following mergers \\cite{ekers-02}. If binary SBHs failed to merge efficiently, uncoalesced binaries would be present in many bright ellipticals, resulting in 3- or 4-body slingshot ejections when subsequent mergers brought in additional SBHs. This would produce off-center SBHs, which seem to be rare or non-existent, as well as (perhaps) too much scatter in the $\\mh-\\sigma$ and $\\mh-L_{bulge}$ relations \\cite{haehnelt-02}.  While the final approach to coalescence of binary SBHs is not well understood, much of their dynamical effect on the surrounding nucleus takes place very soon after the binary forms. The binary quickly (in less than a galactic crossing time) ejects from the nucleus a mass in stars of order its own mass \\cite{quinlan-96,mm-01} significantly lowering the central density on parsec scales. There is reasonably quantitative agreement between this model and the observed structure of nuclei: the ``mass deficit'' -- the stellar mass that is ``missing'' from the centers of galaxies, assuming that they once had steep cusps like those observed at the centers of faint ellipticals -- is of order the black hole mass \\cite{milos-02,ravin-02,graham-04}.  While the binary SBH model is compelling, there is still not much hard evidence in its support. Observationally, no bona fide binary SBH (i.e. gravitationally bound pair of SBHs) has definitely been detected, although there is circumstantial evidence (precessing radio jets; periodic outburst activity) for SBH binaries in a number of active galaxies, as reviewed briefly below (see \\cite{komossa-03b} for a more complete review of this topic). But the binary SBH model has one great advantage: the postulated effects are accessible to observation, since they extend to scales of $1-100$ pc, the distance out to which a binary SBH can significantly influence stellar motions. Much of the recent theoretical work in this field has been directed toward understanding the influence of a binary SBH on its stellar surroundings and looking for evidence of that influence in the distribution of light at the centers of galaxies.  Following the definition of terms and time scales in \\S\\ref{sec:prelim}, we present a brief overview of the observational evidence for binary SBHs in \\S\\ref{sec:obs}. Interaction of a binary SBH with stars is discussed in \\S\\ref{sec:stars}. The possibility of multiple SBHs in galactic nuclei, and the implications for coalescence, are discussed in \\S\\ref{sec:mult}. \\S\\ref{sec:nbody} summarizes $N$-body work on the evolution of binary SBHs, with an emphasis on the question of binary wandering. Observational evidence for the destruction of nuclear density cusps is reviewed in \\S\\ref{sec:cusp}. In some galaxies, the predominant source of torques leading to decay of the binary may be gas; this topic is reviewd in \\S\\ref{sec:gas}. Finally, the influence of binary coalescence on SBH spins is summarized in \\S\\ref{sec:spin}. ",
        "conclusions": "\\begin{enumerate} \\item No completely convincing case of a bound, binary SBH has yet been disccovered. ``Dual'' SBHs, i.e. two, widely-separated SBHs in a single system, have been seen in some interacting galaxies and binary quasars. However none has a projected separation less than $\\sim 1$ kpc. Strong limits can be placed on the binarity of the Milky Way SBH. \\item The evolution of binary SBHs in gas-poor galaxies is dominated by the gravitational slingshot ejection of stars that pass near the binary, carrying away energy and angular momentum. Once the binary has ejected all stars on intersecting orbits, continued hardening depends on a refilling of the binary's ``loss cone.'' Possible mechanisms for loss-cone refilling include star-star gravitational scattering, chaotic orbits in non-axisymmetric potentials, and perturbations from additional massive objects. \\item Most $N$-body simulations of binary evolution have been based on such small particle numbers that the binary's loss cone was refilled at a spuriously high rate by gravitational encounters, Brownian motion of the binary and other finite-$N$ effects. As a result, many results from these studies can not usefully be extrapolated to the large-$N$ regime; in particular, they do not make useful predictions about the binary hardening rates expected in real galaxies. \\item $N$-body studies do appear to make robust predictions about the ``mass deficit,''  the mass in stars ejected from a galactic nucleus by an evolving binary. Observed mass deficits are of order $1-2$ times the SBH mass, consistent with $N$-body predictions. \\item Studies of the interaction of binary SBHs with gas are in their infancy. Major uncertainties are the amount, distribution and thermodynamic state of gas very near the centers of galaxies containing massive binaries. \\item Binary coalescence can have a large influence on SBH spins; even mass ratios as extreme as $10:1$ can substantially spin up or re-orient a SBH. Evidence for spin-flips may have been observed in the so-called $X$-shaped radio sources.  \\end{enumerate}"
    },
    "0410/astro-ph0410687_arXiv.txt": {
        "abstract": "\\noindent We derive the {\\it minimal} seed magnetic field which unavoidably arises in the radiation and matter eras, prior to recombination, by the rotational velocity of ions and electrons, gravitationally induced by the non-linear evolution of primordial density perturbations. The resulting magnetic field power-spectrum is fully determined by the amplitude and spectral index of density perturbations. The {\\it rms} amplitude of the seed-field at recombination is $B \\approx 10^{-23} (\\lambda/{\\rm Mpc})^{-2} $ G, on comoving scales $\\lambda \\gsim 1$ Mpc. ",
        "introduction": "Magnetic fields are present in most astrophysical systems \\cite{obsreview,Torn,Grasso,Widrow,Carilli}, but their origin is still unknown. Spiral galaxies typically contain magnetic fields of about $10^{-6}$ G that are aligned with the spiral density waves \\cite{obsreview}. A plausible explanation is that these fields have been produced from the exponential amplification of an initially weak field by a mean-field dynamo \\cite{dynamo,moffatt}, in which a seed-field was amplified by the differential rotation of the galaxy in conjunction with magnetohydrodynamic turbulence. Magnetic fields of comparable amplitude are also found on cluster scales \\cite{bolo}.  The seed-field strength required at the time of completed galaxy formation for a galactic dynamo to produce the present magnetic field strength of a few $\\mu$G is usually quoted in the range $\\sim (10^{-23}$ -- \\mbox{$10^{-19}$) G} \\cite{obsreview,Torn,Grasso,Widrow,Carilli}. These estimates, though, are obtained by considering the dynamo amplification in a flat Universe with zero cosmological constant for ``typical'' values of the parameters of the $\\alpha\\omega$--dynamo. According to Ref.~\\cite{dlt} these lower bounds can be relaxed to about $10^{-30}$ G, for a Universe with a dark energy component (e.g. a cosmological constant or quintessence), which appears to be favored by recent results from high-redshift Supernovae observations, cosmic microwave background experiments and large--scale structure  data.   Most proposed models of primordial seed-field generation either fail to meet these requirements or invoke non-standard coupling between the electromagnetic field and the inflaton scalar sector responsible for a primordial period of accelerated Universe expansion (see, e.g. Ref.~\\cite{Widrow} and references therein).  The most economic and conservative physical mechanism for producing such a seed-field was proposed by Harrison \\cite{Harrison}. The mechanism relies on the fact that weak magnetic fields are generated during the radiation era in regions that have non-vanishing vorticity. The main problem with Harrison's mechanism is that the required vorticity of the plasma does not have any dynamical origin, but has to be put in by hand, as an initial condition. In this paper we show that rotational velocity of the plasma is unavoidably induced gravitationally by the nonlinear mode-mode coupling of primordial density perturbations generated during inflation, arising at second order in perturbation theory. Therefore, we predict that a minimal seed magnetic field is generated  by the differential rotational velocity of ions and electrons. The resulting magnetic field power-spectrum is fully determined by the amplitude and the spectral index $n_S$ of primordial scalar perturbations. The {\\it rms} amplitude of the seed-field at recombination is $B \\approx 10^{-23} (\\lambda/{\\rm Mpc})^{-2}$ G, on comoving scales $\\lambda \\gsim$ Mpc.  The plan of the paper is as follows. In Section II we briefly recall Harrison's mechanism and we show how the seed magnetic field is produced. In Section III we compute the power-spectrum of such a seed-field in terms of the power-spectrum of primordial scalar perturbations generated during inflation. Finally, in Section IV we draw our conclusions. ",
        "conclusions": "In this paper we have discussed a new mechanism for the generation of the cosmic seed magnetic field. It acts during the evolution of the Universe up to the epoch of hydrogen recombination. The underlying physical phenomenon is the vorticity of charged particles driven by gravitational vector modes. The latter, in turn, arise from the non-linear evolution of purely scalar (density) primordial perturbations. Hence, our mechanism is a generic prediction of the standard hierarchical structure formation scenario and does not require any ad-hoc assumption.  Maybe the most important feature of the created magnetic field discussed in this paper is that its power-spectrum is fully determined by the power-spectrum of the primordial density perturbations. As a consequence, a significant signal is expected over cosmological scales much greater than those encountered in other mechanisms acting during the early evolution of the Universe."
    },
    "0410/astro-ph0410478_arXiv.txt": {
        "abstract": "We present and discuss the properties and the main results of a cosmological model with a spontaneously-broken chiral symmetry. The model contains and relates dynamically, two spin-zero fields. The scalar field can provide the dynamical basis for inflation in the early universe. The pseudoscalar, Goldstone field can provide an early, small residual vacuum energy density, the absolute value of which we estimate to be similar to the present, empirically small vacuum energy density. The small energy scale for this effective cosmological constant is estimated separately, by relating it dynamically to the empirical, small scale of neutrino mass. CP invariance is broken spontaneously. This provides a natural basis for the early generation of an antineutrino-neutrino asymmetry, whose magnitude we estimate, and find to be significant. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410152_arXiv.txt": {
        "abstract": "Using available multiwavelength data for the pulsars of different ages detected in the optical range,  we analyze the efficiencies of the conversion of the pulsar spindown power ${\\rm \\dot{E}}$ into the observed non-thermal luminosity ${\\rm L}$ in different spectral domains. We find that these pulsars show significantly non-monotonic dependencies of the optical and X-ray  efficiencies (${\\rm\\eta = L/\\dot{E} }$) versus pulsar age with a pronounced  minimum at the beginning of the middle-age epoch (${\\rm \\sim 10^4}$~yr) and  comparably higher  efficiencies of  younger  and older pulsars. In addition, we find a strong  correlation between the optical and 2--10 keV X-ray luminosities of these pulsars that implies the same origin of their nonthermal  emission in both spectral domains. ",
        "introduction": "\\label{} The properties of the observed pulsar emission show that radio and gamma-ray  photons   have  nonthermal  origin  and   are  produced  by relativistic particles in different parts of the pulsar magnetosphere. The soft  X-ray-optical radiation can be a combination of several  nonthermal  spectral components  from  the magnetosphere  and thermal components from the whole surface of a cooling neutron star (NS) and from its hot polar caps heated by the relativistic particles from the magnetosphere. The nonthermal component has been found to dominate the optical emission of several young, middle-aged, and old pulsars.  It is believed that nonthermal components are powered by NS rotational energy loss $\\dot{E}=4\\pi^2 I \\dot{P}/{P}^3$, where $I\\approx 1.1\\times 10^{45}\\ gm\\  cm^{2}$ is the inertia momentum of a NS of 1.4 $M_\\odot$ and 10~km radius, $P$ and $\\dot{P}$ are the pulsar period and its derivative, respectively. $\\dot{E}$ is frequently called also the spindown luminosity and denoted as $L_{sd}$. The ratio $\\eta =L/\\dot{E}$, where $L$ is the photon luminosity, is used to characterize the efficiency of the conversion of $\\dot{E}\\equiv L_{sd}$ into the pulsar emission at a given spectral domain.  In this paper, we analyze the evolution of the nonthermal luminosities and efficiencies of pulsars detected in the optical range. Such analysis has been performed about 10~yr ago in pioneering works  \\cite{Mereg,Goldoni}. However, considerable progress in the pulsar distance measurements, and much higher quality of recent optical and X-ray observations allowing to resolve better the nonthermal and thermal emission components from the pulsars and to identify in these ranges several new  objects enable us to update significantly previous analysis and to get qualitatively new results on the evolution of old pulsars. \\begin{table*}[t] \\caption{The dynamical ages ${\\rm \\tau =P/2\\dot{P}}$, distances ${\\rm d}$, spindown luminosities ${\\rm \\dot{E}~(or~L_{sd})}$, and observed nonthermal luminosities in the radio, ${\\rm L_{R}}$, optical, ${\\rm L_{Opt}}$, X-rays, ${\\rm L_{X}}$, and gamma-rays, ${\\rm L_{\\gamma}}$, of seven radio pulsars detected in the optical range \\cite{Zhar1,ZharS}. Figures in brackets are ${\\rm \\pm1\\sigma}$  uncertainties of the values. } \\begin{tabular}{llllllll} \\hline\\hline Source     &log ${\\tau}$ & d                &  log ${\\rm \\dot{E}}$&log ${\\rm L_{R}}^a$\t& log ${\\rm L_{Opt}}$ &log ${\\rm L_{X}}$ & log ${\\rm L_{\\gamma}}$\\\\ \\vspace{-0.3cm} &                  &                  &                    &                  &            &           &         \\\\ &[yr]              &  [pc]            & [{\\rm erg~s$^{-1}$}]& $\\rm [ mJy~kpc^2]$      & [{\\rm erg~s$^{-1}$}]  &  [{\\rm erg~s$^{-1}$}] &  [{\\rm erg~s$^{-1}$}] \\\\ \\vspace{-0.3cm} &                  &                  &                     &                  &            &           &         \\\\ &                  &                  &         &  408 MHz  &    B band  &    2-10 keV        &  $ \\ge$400MeV           \\\\ \\hline          \\vspace{-0.4cm} &                                    &         &                           &                   &                  &               &  \\\\% \\vspace{-0.5cm} Crab       &  3.1\t      &2.0$\\times10^3$&  38.65  & 3.41(4)\t  &33.23(5)\t\t  &$36.67(^{+20}_{-26})$  & $35.7(^{+1}_{-3})$   \\\\  \\vspace{-0.3cm} &\t              &\t\t         &\t   &\t\t&\t  \\\\ B0540-69   & 3.2     &5.0$\\times10^4$& 38.17   & 3.30(9)&33.47(15)\t\t   &$36.99(^{+19}_{-23})$  &$ \\le 35.97$ \\\\ \\vspace{-0.3cm} &\t\t      &\t      \t         &\t   &\t\t&\t  \\\\ Vela       & 4.1    & ${293(^{+19}_{-17})} $   & 36.84   & 2.64(20)   &28.3(3)\t   & $31.2(^{+36}_{-38})^i$ & $33.9(^{+1}_{-3})$  \\\\ \\vspace{-0.3cm} &   \t\t      &\t\t         &\t   &\t\t&\t  \\\\ B0656+14   &  5.0   & ${288(^{+33}_{-27})}$    & 34.58   & -0.27(9)\t&27.53(8)\t\t    & $30.30(^{+36}_{-28})$ & $32.37(^{+10}_{-30})$ \\\\ \\vspace{-0.3cm} &\t\t      &\t\t         & \t   &\t        &\t \\\\ Geminga    & 5.5      &${153(^{+59}_{-34})}$  & 34.51   &${0.375(^{+27}_{-23})}$&  ${26.95(^{+16}_{-10})}$& $29.35(^{+38}_{-36})$ &$32.95(^{+10}_{-30})$ \\\\ \\vspace{-0.3cm} &\t\t      &\t\t         & \t   &\t        &\t \\\\ B1929+10   & 6.5\t  & 361$^{10}_{-8}$  & 33.59   & 1.52(5)    &${27.26(^{+20}_{-33})}$&$29.86(^{+13}_{-15})$ & $\\le 32.57$\t\\\\  \\vspace{-0.3cm} &\t\t      &\t\t         & \t   &\t        &\t \\\\ B0950+08   &  7.2\t      &262(5)       & 32.75   & 1.44(16)\t&26.88(8)\t&$29.28(^{+13}_{-18})$ & $ \\le 32.51$ \\\\  \\hline  \\end{tabular} \\begin{tabular}{l} {\\footnotesize$^a$ The radio  luminosity is defined traditionally as  ${\\rm L_{_{R}} =S_{408}d^2_1~[mJy~kpc^2]}$, where ${\\rm S_{408}}$ is the observed flux } \\\\ {\\footnotesize density from  a pulsar   at 408 MHz in ${\\rm mJy}$, and ${\\rm d_1}$ is its distance in kpc. At a typical radio band  FWHM } \\\\ {\\footnotesize of  about 100 kHz  the conversion factor to  standard luminosity units is ${\\rm \\approx9.51\\times10^{21}~[erg~s^{-1}]/[mJy~kpc^2]}$.} \\end{tabular} \\label{Pulsars} \\end{table*} \\begin{figure}[t] \\setlength{\\unitlength}{1mm} \\resizebox{8.7cm}{!}{\\begin{picture}(80,115)(0,0) \\put ( 3,18) {\\includegraphics[width=56mm, clip]{zharikovfig1.eps}} \\put (-3, 27){ \\parbox[t]{75mm} { \\caption{{\\it From top to bottom:} evolution of the  pulsar spindown, radio, optical, X-ray, and gamma-ray  luminosities and respective efficiencies, with  dynamical age. Filled circles  are the data  taken  from \\cite{Zav}. Hereafter errorbars and arrows show ${\\rm\\pm1\\sigma}$  value uncertainties and ${\\rm 1\\sigma}$ upper limits, respectively. }}} \\end{picture} } \\label{fig1} \\end{figure} \\begin{figure}[t] \\setlength{\\unitlength}{1mm} \\resizebox{8.7cm}{!}{ \\begin{picture}(80,76)(0,0) \\put ( -1.5,38) {\\includegraphics[width=69.5mm, clip]{zharikovfig2.eps}} \\put ( 0,-7) {\\includegraphics[width=68mm, clip]{zharikovfig3.eps}} \\end{picture} } \\caption{{\\sl Top:} radio efficiencies of 500 pulsars from ATNF \\cite{http} {\\it vs} dynamical age; triangles mark the pulsars with more accurate parallax-based distances and  circles select optical pulsars from Table 1. {\\sl Bottom:} X-ray efficiencies {\\it vs} dynamical age; lines are the linear regression fits of evolution tracks of younger and older pulsars and numbers show the line  slopes.} \\label{fig2} \\end{figure} \\begin{figure}[tb] \\includegraphics[width=73mm,clip]{zharikovfig4.eps} \\vspace{-1.1cm} \\caption{Relationship between the optical and X-ray efficiencies in the B band and 2-10~keV ranges.} \\label{fig4} \\vspace{-0.7cm} \\end{figure} \\begin{figure}[tb] \\includegraphics[width=73mm,clip]{zharikovfig5.eps} \\vspace{-1.1cm} \\caption{The evolution of the optical and X-ray luminosities ratio with dynamical  age.} \\label{fig5} \\vspace{-.7cm} \\end{figure} ",
        "conclusions": "Our multiwavelength analysis of the nonthermal luminosity and efficiency distributions {\\it vs} dynamical age and spindown power of the pulsars detected in the optical range suggests significantly nonmonotonic evolution of their emission in the optical and X-ray domains and/or a bimodal distribution of the respective efficiencies {\\it vs} spindown power. Less energetic old pulsars are much more efficient photon emitters than the middle-aged ones and their efficiency is comparable with that for the much younger and energetic Crab-like objects. The bimodal structure is  not significant but still resolved in a larger sample of pulsars identified  in X-rays, where it may be smoothed by large distance and flux uncertainties and needs further study.  % An apparent evolution track of old pulsars in the optical and X-rays is roughly compatible with the track seen in the radio range where the evolution is accompanied by a significant increase of the radio-efficiency with NS  age. If it is true, this implies % an universal engine that drives the non-thermal evolution of old pulsars by a unified way increasing their efficiency with age in the all spectral domains.\\\\  This work has been partially supported by CONACYT project 36585-E, and RFBR grants 02-02-17668, 03-02-17423, 03-07-90200 and RLSS 1115.2003.2"
    },
    "0410/hep-th0410222_arXiv.txt": {
        "abstract": "The gravitational wave (GW) signals emitted by a network of cosmic strings are reexamined in view of the possible formation of a network of cosmic superstrings at the end of brane inflation. The reconnection probability $p$ of intersecting fundamental or Dirichlet strings might be much smaller than 1, and the properties of the resulting string network may differ significantly from those of ordinary strings (which have $p=1$). In addition, it has been recently suggested that the typical length of newly formed loops  may differ by a factor $\\epsilon \\ll 1$ from its standard estimate. Here, we analyze the effects of the two parameters $p$ and $\\epsilon$ on the GW signatures of strings. We consider both the GW bursts emitted from cusps of oscillating string loops, which have been suggested as candidate sources for the LIGO/VIRGO and LISA interferometers, and the stochastic GW background, which may be detectable by pulsar timing observations. In both cases we find that previously obtained results are \\textit{quite robust}, at least when the loop sizes are not suppressed by many orders of magnitude relative to the standard scenario. We urge pulsar observers to reanalyze a recently obtained 17-year combined data set to see whether the large scatter exhibited by a fraction of the data might be due to a transient GW burst activity of some sort, e.g.  to a  near cusp event. ",
        "introduction": "Cosmic strings can be formed as linear defects at a symmetry breaking phase transition in the early universe and can give rise to a variety of observable phenomena at the present cosmic age. String formation, evolution, and observational effects have been extensively studied in the 1980's and 90's (for a review see \\cite{book,kibblehindmarsh}). This exploration was to a large degree motivated by the string scenario of structure formation \\cite{Zeldovich,AV81a}, which requires strings of the grand unification energy scale, $\\eta \\sim 10^{16}$ GeV, or $G\\mu \\sim (\\eta/M_p)^2 \\sim 10^{-6}$. ( Here, $G$ is Newton's constant, $\\mu$ is the string tension and $M_p\\sim 10^{19}$ GeV is the Planck mass. $G\\mu$ is a dimensionless parameter characterizing the gravitational interactions of strings.) This scenario is now disfavored by the CMB observations, but strings of a somewhat lower energy scale would be consistent with the data, and their detection would of course be of great interest. The current CMB bound on strings is \\cite{Pogosian,Smoot} $G\\mu < 6.1 \\times 10^{-7}$. There is also a case of potential string detection with a similar value of $G\\mu$. Two nearly identical galaxies are observed at angular separation of 1.9 arc sec, suggesting gravitational lensing by a cosmic string with $G\\mu\\sim 4\\times 10^{-7}$ \\cite{Sazhin}.  Cosmic strings can also be detected through the gravitational wave (GW) background produced by oscillating string loops \\cite{AV81b}. This background, which ranges over many decades in frequency, has been extensively discussed in the literature \\cite{AV81b,Hogan,VV85,BB90,CA,Caldwell96}. The analysis of eight years of millisecond pulsar timing observations has led to setting rather stringent (95\\% confidence level) limits on the GW contribution to the cosmological closure density $\\Omega_{g} \\equiv \\rho_{g} / \\rho_c \\equiv 8 \\pi G \\rho_{g}/ (3 H_0^2)$: $\\Omega_{g} h^2 < 6 \\times 10^{-8}$ according to the original analysis \\cite{KTR94}, or $\\Omega_{g} h^2 < 9.3 \\times 10^{-8}$ according to the Bayesian approach of \\cite{MZVL96}. (Here, $h \\equiv H_0/(100 {\\rm km/s/Mpc})$. Note that $h^2 \\simeq (65/100)^2 \\simeq 0.42$.)  As we shall review below, the corresponding bound on $G\\mu$ is $ c^{3/2}G\\mu < 10^{-7} $, where $c$ denotes the (mean) number of cusp events per oscillation period of a string loop.  Until recently, it appeared that the gravitational effects of strings with $G\\mu\\ll 10^{-7}$ are too weak to be observable. However, it has been shown in \\cite{DV1,DV2} (following a suggestion in \\cite{BHV}) that GW bursts emitted from cusps of oscillating loops should be detectable by LIGO and LISA interferometers for values of $G\\mu$ as low as $10^{-13}$.  During the last few years, there have been several important developments that motivate us to reexamine the GW signatures of strings. First, there has been a renewed interest in the possibility \\cite{Witten85} that fundamental strings of superstring theory may have astronomical sizes and play the role of cosmic strings. In particular, it has been argued \\cite{Tye1,Tye2,Dvali} that fundamental (F) and D-string networks can naturally be formed at the end of brane inflation. In this scenario \\cite{DvaliTye}, inflation is driven by the attractive potential between a D-brane and an anti-D-brane, and strings are produced when the branes eventually collide and annihilate. The rather stringent requirements allowing for the stability of such cosmic superstrings can be met in some scenarios \\cite{Dvali,polchinski,polchinski2}. The predicted string tensions are \\cite{Tye1,Tye2,KKLMMT,polchinski} $10^{-11}\\lesssim G\\mu\\lesssim 10^{-6}$ and appear to be comfortably within the range of detectability by LIGO and LISA. However, the analysis of GW bursts in \\cite{DV1,DV2} may not be directly applicable in this case, since the properties and evolution of F and D-string networks may differ in significant ways from those of ``ordinary'' cosmic strings.  If both F and D-strings are produced, they can form an interconnected network, in which F and D-strings join and separate at 3-way junctions \\cite{polchinski,Dvali}. Each junction joins an F-string, a D-string, and a (1,1)-string, which is a bound state of F and D. $(p,q)$-strings, which are bound states of $p$ F-strings and $q$ D-strings with $p,q>1$ can also be formed. The evolution of FD-networks is similar to that of monopole-string $Z_3$ networks, in which each monopole is attached to three strings. Simulations of $Z_3$ network evolution suggest \\cite{VV87,Spergel} that the typical distance between the monopoles scales with the cosmic time, $L\\sim \\gamma t$, where the coefficient $\\gamma$ depends on the rate of energy loss by the network. But if the main energy loss mechanism is gravitational radiation, as the case may be for an FD network, then it has been argued in \\cite{VV87} that the energy dissipation by the network is rather inefficient, so it quickly comes to dominate the universe.  If this picture is correct, then models predicting FD networks are ruled out. It is conceivable, however, that the main energy loss mechanism of FD networks is not GW emission, but chopping off of small nets, similar to closed loop production by ordinary strings. The negative verdict on this type of models can then be avoided. This issue can only be resolved with the aid of new, high-resolution numerical simulations of FD networks.  In this paper we shall focus on models where only one type of string is formed. Still, the string evolution may differ from that of ordinary strings, because the reconnection probability $p$ for intersecting strings may be significantly smaller than 1. When ordinary strings intersect, they always reconnect \\cite{Shellard,Matzner}. For intersecting F-strings, the reconnection probability is suppressed by the string coupling, $g_s^2 <1$. Moreover, strings moving in a higher-dimensional bulk can avoid intersection much more easily than strings in 3 dimensions \\cite{Dvali,Tye2}. The string propagation in the bulk is expected to be restricted by bulk potentials, but the effective reconnection probability can still be reduced by an order of magnitude or so. Analysis in \\cite{polchinski2} suggests reconnection probabilities in the range \\beq 10^{-3}\\lesssim p\\lesssim 1 \\label{p} \\eeq for F-strings and \\beq 0.1\\lesssim p\\lesssim 1 \\eeq for D-strings. We thus need to analyze the effect of a reduced reconnection probability on the GW burst statistics and on the stochastic GW background.  Another interesting recent development has been the analysis by Siemens and Olum \\cite{Olum1} of the gravitational radiation from  counter-streaming wiggles on long strings. They showed that this radiation is much less efficient in damping the small-scale wiggles than originally thought. This may result in much smaller sizes of closed loops produced by the network \\cite{Olum2} than previously assumed. This effect can be quantified by the dimensionless  parameter $\\epsilon$ defined in Eq.~(\\ref{epsilon}) below. Refs.~\\cite{DV1,DV2} had assumed (besides $p=1$), the ``standard'' value $\\epsilon =1$, and we need to see how a different value of $\\epsilon$ might affect the GW burst statistics.  Last, but certainly not least on our list of recent developments is the potential improvement in the sensitivity to a GW background of pulsar timing observations. Indeed, these have been recently extended to a 17-year data set\\cite{Lommen}. We shall discuss below what kind of limits on $G\\mu$ follow from this extended data set.  The main goal of the present paper is to analyze how the amplitude and frequency of GW bursts from strings and the intensity of the stochastic GW background depend on the parameter $\\epsilon$ measuring the characteristic size of closed loops, and on the string reconnection probability $p$. The paper is organized as follows. In the next Section we outline the relevant features of string evolution. GW bursts from cusps are discussed in Section III. The stochastic background is discussed in Section IV, where we also discuss bounds on $G\\mu$ based on millisecond pulsar observations. Our conclusions are summarized in Section V. ",
        "conclusions": "It has been argued \\cite{Tye1,Tye2,Dvali,KKLMMT,polchinski} that F- and D-string networks can naturally be formed at the end of brane inflation \\cite{DvaliTye}, with string tensions in the range $10^{-11} \\lesssim G \\mu \\lesssim 10^{-6}$. We focussed on models where only one type of string (F or D) is formed, and considered the gravitational wave (GW) signatures of cosmological networks of such strings. We studied how the finding \\cite{DV1,DV2} that GW bursts emitted from cusps of oscillating loops should be detectable by LIGO and LISA interferometers for values of $G \\mu$ as small as $10^{-13}$ might be modified by two separate effects. First, the reconnection probability $p$ for intersecting F or D strings might be, contrary to the case of ordinary field-theory strings, significantly smaller than 1 \\cite{Dvali,Tye2,polchinski2}. Second, Refs.~\\cite{DV1,DV2} had assumed that the characteristic size of newly formed loops was $ l(t) \\sim \\alpha t$, with $\\alpha \\sim 50 G \\mu$, as expected from from standard GW radiation-reaction arguments \\cite{BB90}. However, recent analyses \\cite{Olum1,Olum2} have suggested that gravitational radiation is less efficient than originally thought, and might result in a much smaller typical size for newly formed loops: $ l(t) \\sim \\alpha t$, where $\\alpha \\sim \\epsilon 50 G \\mu$, with $\\epsilon \\ll 1$.  A detailed analysis of the effects of the two parameters $p$ and $\\epsilon$ on on the detectability of GW bursts by LIGO or LISA has shown that the results of \\cite{DV1,DV2} are \\textit{quite robust}, at least when $\\epsilon \\gtrsim 10^{-11}$ in the case of LIGO, or $\\epsilon \\gtrsim 10^{-7}$ in the case of LISA. See Figs. 1-4.  In particular, it is notable that a smaller reconnection probability can \\textit{only increase} the detectability of cosmic superstrings. However, very small values of $\\epsilon$ might lead to cutting off the burst signals in GW detectors when the string tension is $G \\mu \\sim 10^{-10}$ or less.  We have also considered the detectability, via pulsar timing observations, of the stochastic GW background produced by oscillating string loops. When the loop-length parameter $\\alpha$ is large enough, namely $\\alpha = 50 \\epsilon G \\mu > 10^{-9}$, for the GW frequencies emitted by string loops at redshift $z \\sim 1$ to fall within the pulsar sensitivity band $ f \\sim 1/( 10 {\\rm yr})$, we find that the intensity of the GW background \\textit{increases} $\\propto p^{-1} \\epsilon^{-1/3}$, when either $p$ or $\\epsilon$ gets smaller than 1. In addition, we find that present pulsar timing experiments have the potential of detecting string tensions as small as $ G \\mu \\sim 3.6 \\times 10^{-10} c^{-3/2} p^{3/2} \\epsilon^{1/2}$, where $c$ denotes the number of cusps per string oscillation. We urge pulsar observers to reanalyze a recently obtained 17-year combined data set to see whether the large scatter exhibited by a fraction of the data might be due to a transient GW burst activity of some sort. We note that future versions of the ``pulsar timing array'' \\cite{Lommen,Backer1,Backer2,SKA} might further improve the sensitivity of pulsar observations. The ultimate sensitivity of pulsar timing experiments might then become limited  (by the ability of pulsar timing to probe the frequencies emitted by string loops) to the level $G \\mu \\sim 2 \\times 10^{-11} \\epsilon^{-1}$. In other words, if the suggestion of \\cite{Olum1,Olum2} is confirmed, and the reduced  loop-length parameter $\\epsilon \\equiv \\alpha/(50 G \\mu)$ turns out to take very small values $\\epsilon \\ll 1$ (possibly of the form $\\epsilon \\sim (50 G \\mu)^m$, with $m>0$), and if $ G \\mu$ itself takes small values, $ G \\mu \\sim 10^{-10}$ \\cite{KKLMMT,polchinski}, the GW frequency band probed by pulsar timing experiments might fall in the domain $\\alpha = 50 \\epsilon G \\mu <10^{-9}$ where the GW spectrum from loops is cut off. In such a case, only higher-frequency GW experiments, such as LISA or, even better, LIGO, might be able to detect GW from string loops.  Our conclusions show that it is urgent to develop a new generation of string network simulations able to determine how the crucial loop-length parameter $\\alpha$ depends on $G\\mu$ and $p$. It would also be quite important to determine the average number $c$ of cusp events per loop oscillation. It is only when $\\alpha$ and $c$ are known that it will be possible to discuss with any reliability the detectability of cosmic superstrings by GW experiments."
    },
    "0410/astro-ph0410109_arXiv.txt": {
        "abstract": "Radio waves are imprinted with propagation effects from ionized media through which they pass.  Owing to electron density fluctuations, compact sources (pulsars, masers, and compact extragalactic sources) can display a wide variety of scattering effects.  These scattering effects, particularly interstellar scintillation, can be exploited to provide \\emph{superresolution}, with achievable angular resolutions ($\\lesssim 1\\,\\mu$arc sec) far in excess of what can be obtained by very long baseline interferometry on terrestrial baselines. Scattering effects also provide a powerful \\emph{sub-AU} probe of the microphysics of the interstellar medium, potentially to spatial scales smaller than 100~km, as well as a tracer of the Galactic distribution of energy input into the interstellar medium through a variety of integrated measures.  Coupled with future $\\gamma$-ray observations, SKA observations also may provide a means of detecting fainter compact $\\gamma$-ray sources.  Though it is not yet clear that propagation effects due to the intergalactic medium are significant, the SKA will either detect or place stringent constraints on intergalactic scattering. ",
        "introduction": "\\label{sec:turb.intro}  All radio observations of Galactic and extragalactic objects are conducted while viewing these objects through the Galaxy's interstellar medium (ISM).  Dispersion of pulsar signals and optical observations, particularly of the H$\\alpha$ emission line \\cite{hrtmjp03}, indicate that the interstellar medium contains a diffuse ionized component in addition to classical H\\,\\textsc{ii} regions.  This interstellar plasma occupies a significant fraction ($\\sim 0.2$) of the volume of the \\hbox{ISM}, and the energy required to keep it ionized is considerable, roughly 15--20\\% of the luminosity of all of the O stars in the Galaxy.  The interstellar plasma is also known as the Warm Ionized Medium (WIM) or the Diffuse Ionized Gas (DIG).  Significantly for the \\hbox{SKA}, the interstellar plasma will produce observable effects on every line of sight over at least a fraction of the entire proposed operating frequency range.  The SKA can allow unprecedented use of interstellar scattering and interstellar scintillation (ISS) for study of sources and intervening media:  \\begin{itemize} \\item Resolving source structure on angular scales $\\lesssim 10^{-6}$~arcsec will probe the inner jet regions in active galactic nuclei and individual source regions in pulsar magnetospheres.  Astrophysically, such observations may be crucial to an understanding of acceleration of particles in these and other kinds of sources.  The SKA's sensitivity over a wide range of frequencies is needed to exploit the resolving power of \\hbox{ISS}.  \\item Probing microturbulence in the ionized ISM through sampling of scattering phenomena along large numbers ($\\gtrsim 10^4$) lines of sight.  The microturblence plays important roles in the energetics of the ISM and in the propagation of cosmic rays in the Galaxy's magnetic field.  \\item Probing the intergalactic medium (IGM) through detection of angular broadening in excess of that expected from foreground Galactic gas. The SKA allows angular broadening measurement of large samples of active galactic nuclei (e.g., $10^6$ lines of sight) that will map out contributions from individual intervening galaxies, galaxy clusters, Ly-$\\alpha$ clouds, and from a pervasive but clumpy ionized IGM that comprises most of the baryonic matter in the Universe. \\end{itemize}  Radio-wave scattering from density fluctuations in the interstellar plasma offers both a means of characterizing the ISM on sub-AU scales as well as a powerful probe of source structure.  In the remainder of this section, we describe observable effects and astrophysical measures to quantify the interstellar plasma.  Subsequent sections discuss interstellar scintillation and superresolution of source structure (\\S\\ref{sec:turb.structure}), the microphysics of the interstellar plasma (\\S\\ref{sec:turb.physics}), interstellar scattering as a means to measure distances (\\S\\ref{sec:turb.distribute}), the possibility of intergalactic scattering (\\S\\ref{sec:turb.igm}), the potential impact of scintillations on SKA observations (\\S\\ref{sec:turb.problems}), and key observations to be carried out with the SKA (\\S\\ref{sec:turb.summary}).  \\subsection{Observable Effects}\\label{sec:turb.effects}  The phase of a radio wave of wavelength~$\\lambda$ after having propagated through a plasma is \\begin{equation} \\phi = \\int dz\\,\\lambda\\,r_en_e(z), \\label{eqn:turb.phase} \\end{equation} where $r_e$ is the classical electron radius and~$n_e$ is the plasma density.  Soon after the discovery of pulsars, it became clear that their signals were displaying intensity scintillations (e.g., \\cite{r70}), akin to the scintillation of star light at optical wavelengths viewed through the Earth's atmosphere.  Intensity scintillations result from constructive and destructive interference and occur if the phase of the wave has been perturbed during its propagation with structure on small transverse scales. As eqn.~(\\ref{eqn:turb.phase}) shows, phase perturbations imply that the propagating wave has been scattered by interstellar electron density fluctuations.  Figure~\\ref{fig:turb.geometry} illustrates the relevant geometry.  \\begin{figure} \\includegraphics[width=\\columnwidth]{turb.geometry_clr.eps} \\vspace{-1.25cm} \\caption{Geometry of radio-wave scattering in the ISM that is responsible for a variety of seeing and scintillation effects.} \\label{fig:turb.geometry} \\end{figure}  Since the detection of pulsar intensity scintillations, a rich variety of additional interstellar scattering observables has been recognized. These include \\cite{r90,cl91}: \\begin{description} \\item[Intensity Scintillation:] The original indication of interstellar density fluctuations, intensity scintillations have been detected not only from pulsars but potentially also from active galactic nuclei (AGN) and from masers.  \\item[Angular broadening (``seeing''):] Compact radio sources, both Galactic and extragalactic, show angular broadening by amounts ranging from $\\lesssim 1$~mas to~1~arcsec at~1~GHz, depending on the distance and direction.  The angular smearing scales approximately as $\\nu^{-2}$.  \\item[Pulse broadening:] Multipath propagation causes a multiplicity of arrival times, usually seen as an exponential-like ``tail'' to pulses from pulsars.  \\item[Angular wandering:] Focussing and defocussing of rays by AU-scale structures can cause apparent shifts in source positions.  \\item[Pulse time of arrival (TOA) fluctuations:]  Changes in geometry caused by proper motions of a pulsar and intervening material induce variations in the amount of plasma along the line of sight.  Also, variable scattering causes variable arrival times.  \\item[Spectral Broadening:]  Broadening of narrow spectral lines from a combination of scattering and time-variable geometry.  This effect is very small in the ISM ($\\lesssim 1$ Hz) and has not been found because there are no known sources that produce spectral lines narrow enough for this effect to be significant.  It has been measured from spacecraft viewed through the interplanetary medium \\cite{hc83}, and, if extraterrestrial transmitters are ever found, it may be important for them. \\end{description}  \\subsection{Integrated Measures}\\label{sec:turb.measures}  As we discuss in more detail below, radio measurements provide four integrals involving the electron density and magnetic field.   Differential arrival times from pulsars provide the \\emph{dispersion measure}, \\begin{eqnarray} \\DM &=& \\int_0^D ds\\, \\ne(s) \\quad ({\\rm  pc\\, cm^{-3}}), \\end{eqnarray} (for $D$ in pc and $n_e$ in cm$^{-3}$).  A magnetized plasma produces Faraday rotation, which is described by the \\emph{rotation measure}, \\begin{eqnarray} \\RM &=& \\frac{e^3}{2\\pi m_e^2c^4} \\int_0^D ds\\, \\ne\\bpar,\\nonumber\\\\ &=& 0.81\\,\\mathrm{rad\\,m^{-2}}\\, \\int_0^D ds\\, \\ne\\bpar, \\label{eqn:turb.rm} \\end{eqnarray} the integral of the magnetic field component parallel to the line-of-sight $\\bpar$ (in $\\mu$G) and which is sensitive to the correlation of electron density and magnetic field.  The \\textit{emission measure} can be measured from recombination line and free-free absorption or emission observations in the radio and observations of H$\\alpha$ in the optical \\begin{eqnarray} \\EM = \\int_0^D ds\\, \\ne^2 \\quad ({\\rm  pc\\, cm^{-6}}) \\end{eqnarray} (for $D$ in pc and $n_e$ in cm$^{-3}$).  As we discuss below, the density fluctuations responsible for interstellar scattering can be described in terms of a power spectrum. The \\textit{scattering measure} is the integral of $\\cnsq$, the coefficient of the wavenumber spectrum for electron-density fluctuations, $\\dne$, \\begin{eqnarray} \\SM = \\int_0^D ds\\, \\cnsq(s) \\quad (\\mathrm{kpc\\,m}^{-20/3}). \\label{eqn:turb.sm} \\end{eqnarray} Cordes et al.~\\cite{cwfsr91} give relations between \\hbox{DM}, \\hbox{SM}, and \\hbox{EM}.  Finally, diffuse $\\gamma$-ray emission is given by \\begin{equation} I_\\gamma = \\int ds\\,n_n(s)q_\\gamma(s), \\label{eqn:turb.g-rayI} \\end{equation} where $n_n$ is the nucleon density which is dominated by the hydrogen number density~$n_H$, i.e., \\hii, \\hi, and~H${}_2$; $q_\\gamma$ is the gamma-ray emissivity per nucleon, which is proportional to the cosmic-ray density; and $ds$ is a path-length element along the line of sight.  Although not yet exploited, combining future radio and $\\gamma$-ray observations may allow fluctuations in $q_\\gamma$, resulting from magnetic field fluctuations, to be traced.  From large sets of measurements of these measures, but, in particular \\DM, \\SM, and \\RM, the spatial distribution and volume filling factor of the electron density and its fluctuations, e.g., as parameterized by a wavenumber spectrum for quantities such as $\\dne$ and $\\delta\\bpar$, can be probed as can the large-scale structure and fluctuations in the magnetic field. ",
        "conclusions": "An important aspect of many of kinds of observations relevant to superresolution imaging and turbulent phenomena is that they represent a class of \\emph{non-imaging} analyses that the SKA should be capable of performing.  In general, these observations require that some (large) fraction of the SKA's collecting area be capable of being operated in a phased-array mode, in which the voltages from the individual collectors are delayed appropriately and then summed.  Exploiting the SKA to study the structure of both extragalactic sources and pulsars as well as of the magnetoionic Galactic interstellar medium requires that a variety of observables be measurable.  In many cases these observables are similar to what is needed to exploit observations of transient radio sources (Cordes et al., this volume) and GRB afterglows (Weiler, this volume), and, as pulsars form an important target, similar to the requirements for exploiting observations of pulsars (Cordes et al., this volume; Kramer et al., this volume).  \\begin{description} \\item[Dynamic spectra] Probing source structure with superresolution as well as the scattering properties of the ISM requires multiepoch dynamic spectra (Figure~\\ref{fig:turb.ds1133+16}).  In general for probing source structure, the key promise of the SKA is its sensitivity.  Many of the relevant observations can be conducted at the present time or in the near future, but only on a small number of sources, so it is not clear to what extent these sources represent typical or representative examples.  The utility of dynamic spectra for a source is that, if there is sufficient sensitivity to form dynamic spectra as a function of time, modifications in the DISS characteristics in the differential dynamic spectra indicate the presence of structure on sub-$\\mu$as angular scales.  For some applications, the dynamic spectra should be obtained for all Stokes parameters.  For pulsars, dynamic spectra formed a function of pulse phase can be used to assess the presence of multiple emitting regions within the pulsar magnetosphere (Figure~\\ref{fig:turb.psr.geometry}).  For GRB afterglows, a similar analysis could be used to detect and monitor the formation and evolution of components within the emitting shock regions.  For \\hbox{AGN}, if they do display \\hbox{DISS}, the relevant scales probed are comparable to those on which jet formation is thought to occur near the central engine.  For the case of \\hbox{AGN}, dynamic spectra observations will complement those of future missions such as Constellation-X and \\hbox{GLAST}, which will observe the highly-energetic emission from the bases of AGN jets.  Particularly for the case of pulsars, dynamic spectra also can be utilized to extract parameters such as the scintillation bandwidth, scintillation time scale, and scintillation velocity, and the arc phenomenon (Figure~\\ref{fig:turb.ss1133+16}) seen in pulsar dynamic spectra can be used to constrain the electron density power spectrum. Although these parameters can be extracted from the current pulsar census with existing telescopes, the much larger pulsar census that the SKA will perform opens up the possibility of qualitatively new analyses in determining both the Galactic (\\S\\ref{sec:turb.galactic}) and line-of-sight (\\S\\ref{sec:turb.distribute}) distributions of scattering material.  These observations of dynamic spectra require time resolutions $\\sim 1$~s and frequency resolutions $\\sim 1$--100~kHz over wide bandwidths.  In order to obtain dynamic spectra as a function of pulsar pulse phase, sampling times of~10 to~100~$\\mu$s are needed.  \\item[Pulse broadening] Measurement of pulse broadening of pulsar pulses enables characterization of SM along heavily scattered lines of sight, complementing those for which dynamic spectra can be used, and constrains the wavenumber spectrum for $\\delta n_e$ and the distribution along the line of sight of scattering material.  As with dynamic spectra, pulse broadening can be determined for a small number of pulsars with existing telescopes, but the SKA offers the promise of a greatly increased number of pulsars toward which pulse broadening can be determined.  In particular, measurements of pulse broadening in the the much larger pulsar census provided by the SKA are important for the possibility of qualitatively new analyses in determining the Galactic distribution of the interstellar plasma (\\S\\ref{sec:turb.galactic}).  ``Clumps'' in the interstellar plasma, as introduced in the NE2001 model, produce enhanced pulse broadening and so may enable the Galactic spiral arms to be traced, in part, by pulse broadening.  These observations require time resolutions, channelization, and post processing identical to that needed for pulsar timing (Cordes et al., this volume).  \\item[Angular broadening] Probing the Galactic (and intergalactic?) and line-of-sight distributions of scattering material as well as comparing radio-wave and $\\gamma$-ray observations requires high angular resolution in order to measure angular broadening.  One of the strengths of angular broadening measurements toward AGN is that it provides a direct measure of the scattering strength along the line of sight.  Other measures of scattering, e.g., pulse broadening, contain line-of-sight weightings for the scattering material.  As for pulse broadening, angular broadening is a powerful probe of the Galactic plasma and, in terms of a qualitative improvement for the Galactic interstellar plasma (\\S\\ref{sec:turb.galactic}), a large census of \\SM\\ values for \\hbox{AGN} would complement the large pulsar census that the SKA will perform.  Moreover, pulsars will not be detectable over cosmological distances, while AGN clearly are, so angular broadening offers the promise of a direct probe of the IGM (\\S\\ref{sec:turb.igm}).  For angular broadening measurements, the most important aspect is transcontinental or intercontinental baselines.  As described in \\S\\ref{sec:turb.igm}, intercontinental baselines at~0.33~GHz or baselines of order 1000~km below~0.2~GHz or both are minimal requirements for angular broadening studies.  \\item[Light curves] In many current concepts, the SKA is envisioned as having a ``core'' and ``outlying stations.''  Not all observations will be able to make use of the entire SKA, e.g., use of the outlying stations produces too low of a surface brightness sensitivity for some kinds of H\\,\\textsc{i} observations.  Thus, a small number of outlying stations could be used to form a sub-array and tasked to measure the flux density of a variety of objects, such as to monitor a large number of AGN for IDV or search for ESEs. \\end{description}  \\begin{table*} \\caption{Requirements for the SKA\\label{tab:turb.require}} \\begin{tabular}{lcc} \\noalign{\\hrule\\hrule} Item     & Requirement   & Motivation \\\\ \\noalign{\\hrule} Time Resolution      & 10--100~$\\mu$s & dynamic spectra \\\\ &                & pulse broadening \\\\ Frequency Resolution & 1--100~kHz     & dynamic spectra \\\\ Configuration        & phased-array   & pulsar observations \\\\ & operations possible & \\\\ Configuration        & sub-arrays possible & light curves\\\\ Configuration        & $> 1000$~km baselines & angular broadening \\\\ \\noalign{\\hrule} \\end{tabular} \\end{table*}"
    },
    "0410/astro-ph0410615_arXiv.txt": {
        "abstract": "Motivated by the idea that the recently detected near-infrared (1.2-4 $\\mu$m) excess over the contribution of known galaxies is due to redshifted light from the first cosmic stars (Salvaterra \\& Ferrara 2003), we have used the effect caused by photon-photon absorption on gamma-ray spectra of blazars to put constraints on extragalactic background light (EBL) from the optical to the far-IR bands. Our analysis is mainly based on the blazar \\astrobj{H~1426+428}, for which we assume a power-law unabsorbed spectrum. We find that an EBL model with no excess over known galaxies in the near-infrared background (NIRB) is in agreement with all the considered blazars; however, it implies a very peculiar intrinsic spectrum for \\astrobj{H~1426+428}. Additional data on the blazars \\astrobj{1ES1101-232}, \\astrobj{H~2356-309} and \\astrobj{PKS~2155-304} exclude the existence of a strong NIRB excess consistent with Kelsall's model of zodiacal light subtraction (ZL); the COBE/DIRBE measurements, after Wright's model ZL subtraction, represent a firm NIRB upper limit. The constraints on the optical EBL are weaker, due to the fact that predictions from different optical EBL models are often comparable to the experimental errors. In the mid-infrared the {\\it SPITZER} measurement of $\\nu{}I_{\\nu{}}$=2.7 nW m$^{-2}$ sr$^{-1}$ at 24 $\\mu{}$m, gives a good fit for all the considered blazars. ",
        "introduction": "The study of the  Extragalactic Background Light (hereafter EBL) might provide unique information to understand many crucial astrophysical questions, including, among others, the first cosmic star formation (Salvaterra \\&{} Ferrara 2003), the evolution of  galaxies (Totani \\&{} Takeuchi 2002; Totani et al. 2001), the role of  dust emission (Granato et al. 2001). Yet, the EBL determination in the optical and infrared bands poses challenging observational and theoretical difficulties, often leading to discrepant interpretations. For example, Madau \\& Pozzetti (2000) suggest that the optical EBL can be completely explained by galactic counts, while Bernstein et al. (2002) claim the presence of an optical excess. Moreover, the presence of an excess (Matsumoto et al. 2000; Wright 2001) in the near infrared background (hereafter NIRB) crucially depends on the subtraction of the zodiacal light (Kelsall et al. 1998; Wright 1997, 1998; Wright \\& Reese 2000). For the middle and far infrared background (hereafter MIRB and FIRB) various models have been proposed (Totani \\&{} Takeuchi 2002; Silva et al. 2004; Primack et al. 1999) which cannot be discriminated on the basis of currently available data (Elbaz et al. 2002; Metcalfe et al. 2003, and references here). Finally, very strong constraints to the MIRB can be provided by the {\\it SPITZER}\\footnote{http://www.spitzer.caltech.edu/index.shtml} satellite (Papovich et al. 2004), which indicates a background total flux of $2.7^{+1.1}_{-0.8}$ nW m$^2$ sr$^-1$ at 24 $\\mu$m. In the next years {\\it SPITZER} (and other projects like ASTRO-F and SPICA) will likely be able to reconstruct the complete map of the MIRB and FIRB. \\\\ Waiting for the upcoming infrared facilities, an alternative way to put independent constraints to the EBL is to resort to the so-called photon-photon absorption of the high energy tail of the blazar spectra. It is known that gamma rays in the GeV-TeV energy bands can be absorbed by softer (mainly optical and infrared) photons, via electron-positron pair production (Nikishov 1962). Stecker, de Jager \\& Salamon (1992) pioneered the method on the spectrum of 3C 279 to first study the EBL. Their attempt stimulated a plethora of works on the subject  (Stecker, de Jager \\& Salamon 1993; Madau \\& Phinney 1996; Stecker \\& de Jager 1997; Aharonian et al. 2003; Konopelko et al. 2003; Stecker 2003; Costamante et al. 2004; Dwek \\& Krennrich 2005) that were encouraged by the availability of new measurements of blazar spectra in the TeV regime based on the Cherenkov imaging technique (Aharonian et al. 1999b; Krennrich et al. 1999 and Aharonian et al. 2002b for \\astrobj{Mkn 421};  Aharonian et al. 1999a, 2001a, 2002b for \\astrobj{Mkn 501}; Aharonian et al. 2003 for \\astrobj{H~1426+428}).  In this paper we calculate the optical depth due to photon-photon absorption using the most recent optical and infrared background measurements (Bernstein et al. 2002; Matsumoto et al. 2000; Papovich et al. 2004), combined with state-of-art theoretical models of the NIRB (Salvaterra \\& Ferrara 2003) and of the MIRB-FIRB (Totani \\& Takeuchi 2002). We compare these results with the Cherenkov imaging data available for six blazars (\\astrobj{H~1426+428}, \\astrobj{Mkn 421}, \\astrobj{Mkn 501}, \\astrobj{PKS~2155-304}, \\astrobj{1ES1101-232}, \\astrobj{H~2356-309}), to derive some constraints on the EBL. The final aim of the study is to assess if the most recent EBL data and models exclude or require the presence of an excess with respect to galaxy counts in the optical and in the near-infrared regime. In addition, we try to test the predictions of the proposed EBL models, especially in the optical, near and middle infrared ranges. \\\\ The paper is organized as follows. Sec. 2 presents the method adopted to derive the optical depth for the photon-photon absorption and to compare it with blazar data. In Sec. 3 we briefly review the most important EBL data and the EBL models considered. Sec. 4 is devoted to the description of  the Cherenkov data available for the three blazars that we analyze. In Sec. 5 we discuss the results and Sec. 6 summarizes our conclusions. ",
        "conclusions": "In this paper, we have explored the effects produced by photon-photon absorption on blazar spectra to put constraints on extragalactic background light from the optical to the far infrared bands. We mostly focused on the high redshift blazar \\astrobj{H~1426+428}, for which we adopted a power-law unabsorbed spectrum. This might be a simplifying assumption, as some theoretical models indicate a more complex spectral shape (Inoue and Takahara 1996; Tavecchio, Maraschi \\& Ghisellini 1998; Fossati et al. 2000). On the other hand, the data for \\astrobj{H~1426+428} are not sufficiently accurate to suggest the existence of a cut-off or to exclude a power-law spectrum.  We have considered three different models of the NIRB, based respectively on galaxy counts (C1) and on the presence of an excess (DW and MK). In the case of \\astrobj{H~1426+428} the DW model provides the best fit both using the method presented in Sec. 5.2, and the alternative one of Sec. 5.6. The MK model has a higher $\\chi{}^2$ but cannot be rejected. For the same reasons, C1 model cannot be excluded either; however, in this case a very peculiar blazar intrinsic spectrum must be assumed (Fig. 8). We conclude that the presence of a NIRB excess with respect to galaxy counts, at the level given by the DW model, seems to be required to fit the blazar spectrum.   The very recent HESS data of the blazars \\astrobj{PKS~2155-304} (Aharonian et al. 2005a), \\astrobj{1ES1101-232} and \\astrobj{H~2356-309} (Aharonian et al. 2005b) allow us to put stronger constraints to the NIRB. In particular, from the analysis of all these blazar, we conclude that the MK models, based on the Kelsall method of ZL subtraction, must be rejected, since they imply too hard blazar spectra. The spectra of \\astrobj{PKS~2155-304} and \\astrobj{1ES1101-232} are marginally consistent with DW models, based on Wright ZL subtraction, which can be considered as an upper limit to the NIRB. The model C1, without NIRB excess, is favored by all these new data. These findings are quite different from those we derived for \\astrobj{H~1426+428}. Since the available data for  \\astrobj{H~1426+428} are old, new measurement of this blazar are eagerly required, to shed some light on this puzzling inconsistency.   The derived constraints on the optical EBL are weaker, due to the fact that deviations between different optical EBL models are comparable to the experimental errors. The fit tends to favor models without an optical excess over galactic light, contrary to the result obtained by  Bernstein et al. (2002). A more solid conclusion on the amplitude of the optical EBL has to await for the next generation of Cherenkov Telescopes as {\\it MAGIC, VERITAS, HESS, GLAST}, or infrared satellites ({\\it CIBER}).  Finally, in the mid-infrared the {\\it SPITZER} measurement of $\\nu{}I_{\\nu{}}$=2.7 nW m$^{-2}$ sr$^{-1}$ at 24 $\\mu{}$m, combined with the EBL model by Totani \\& Takeuchi (2002), allows us to obtain a good fit for all the blazars available. Again, a tremendous advance on the determination of the MIRB/FIRB is expected from   {\\it SPITZER} and the next generation of infrared satellites as {\\it ASTRO-F}, and {\\it SPICA}.\\\\   In summary, recent measurements of blazars in the TeV range seem to exclude the existence of a strong NIRB excess consistent with Kelsall's model of ZL subtraction. The COBE/DIRBE measurements, after Wright's model ZL subtraction, represent a firm NIRB upper limit."
    },
    "0410/astro-ph0410565_arXiv.txt": {
        "abstract": "This paper introduces a novel feature extraction technique for the analysis of spectral line Stokes profiles. The procedure is based on the use of an auto-associative artificial neural network containing non-linear hidden layers. The neural network extracts a small subset of parameters from the profiles (features), from which it is then able to reconstruct the original profile. This new approach is compared to two other procedures that have been proposed in previous works, namely principal component analysis and Hermitian function expansions. Depending on the target application, each one of these three techniques has some advantages and disadvantages, which are discussed here. ",
        "introduction": "\\label{sec:intro}  A new breed of techniques for the diagnostics of solar magnetic fields has been emerging during the last few years. The fundamental idea that inspired them is the recognition that the observable polarization (Stokes) line profiles exhibit conspicuous {\\it patterns} of features associated with the underlying physical conditions in the atmosphere where they form. This has recently motivated a considerable interest in the field of pattern recognition. Potential applications to solar magnetic field diagnostics have been explored in a number of recent papers by \\citeN{RLAT+00}; \\citeN{SNLAL01}; \\citeN{CS01}; \\citeN{SNSA02}; \\citeN{SLA02}; \\citeN{SN03}; \\citeN{dTILA03}; \\citeN{CLAT+03}.  Pattern recognition techniques were originally developed in the fields of artificial intelligence and machine vision, but have found a vast number of applications over a broad spectrum of disciplines (see, e.g., the review by \\citeNP{RG03}). Solar physicists have relied on least-squares fitting procedures to invert the observed spectra since the late 70s, although this only became a generalized practice in the early 90s (for more details see the reviews by \\citeNP{SN01a}; \\citeNP{dTI03}). The least-squares approach is still as useful as ever and is not likely to be replaced by the newer pattern recognition in the near future. The reason is simply that one can obtain much more detailed information from a given set of profiles by using a least-squares inversion. It also offers more flexibility in the physical model used for the inversion. Research on new inversion techniques is intended to develop alternative algorithms that will {\\it complement}, and not {\\it replace} previous least-square procedures.  Nevertheless, it is justified to devote significant efforts towards the development of alternative inversion algorithms. Without going into much detail here (see, e.g., \\citeNP{SN02} for a more thorough discussion), it should be noted that the next generation spectro-polarimeters (both ground-based and space-borne) will deliver enormous dataflows that simply cannot be analyzed with traditional inversion codes. The need for larger fields of view, higher spatial resolution and better time cadence is driving new instrumental developments. These, in turn, demand new analysis techniques that are fast, ideally operating in real-time, and robust, in the sense of not requiring any human supervision. Other important considerations that motivate research on pattern recognition are related to the need for establishing long-term databases of observations (properly inverted), or the desire to obtain real-time maps of the solar magnetic field at the observing site.  As an intermediate step in this process, some authors have sought a suitable set of features in the Stokes profiles of spectral lines and an appropriate algorithm to extract such features. Feature extraction is of great potential interest, e.g. for data compression (which would help circumvent limited spacecraft telemetry), dimensionality reduction, profile classification, inversion pre-processing, etc. The most promising approaches developed so far seem to be the use of principal component analysis (PCA), introduced by \\citeN{RLAT+00}, and the expansion in Hermitian functions (EHF) of \\citeN{dTILA03}. In this paper I present a new method which makes use of artificial neural networks (ANNs), and compare it to PCA and EHF in terms of performance and potential usefulness.  The organization of the present paper is as follows. Section~\\ref{sec:expansions} below reviews the PCA and EHF methods and discusses some important ideas that are relevant for our purposes here. Section~\\ref{sec:ANN} gives a detailed introduction to ANNs and puts forward the concept of auto-associative ANNs, which is the base for the feature extraction algorithm. The comparison among the three methods is presented in \\S\\ref{sec:performance}. Finally, the most important conclusions are discussed in \\S\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  The development of new instrumentations and the need to process ever increasing dataflows is turning the attention of solar physicists to pattern recognition techniques (see references in \\S\\ref{sec:intro}). The present work can be considered a follow up to the papers by \\citeN{RLAT+00} and \\citeN{dTILA03}, introducing a new technique for feature extraction and comparing it to the proposed PCA and EHF proposed in those papers.  The AANN expansion is based on a MLP which is trained to reproduce the profile supplied in the input neurons. However, one (or more) of the MLP hidden layers has fewer neurons than the inputs/outputs layers. Upon successful training, the AANN is able to extract a small set of features from a profile. These features are numbers stored in the bottleneck layer, from which the entire profile can be reconstructed. This works effectively as a decomposition/expansion scheme in which the number of coefficients is set by the number of neurons in the bottleneck layer.  The comparison among the three feature extraction techniques does not reveal a ``winner''. Instead, each method has its own advantages and disadvantages. Depending on the particular problem it may be more interesting to use one or another.  Ultimately, the goal is to use the expansion coefficients from whichever method one might choose for the implementation of some inversion procedure. The most straightforward way to do this would be a look-up table, as \\citeN{SNLAL01} did for PCA. Other alternatives include neural networks or support vector machines, which could benefit from the dimensionality reduction offered by these techniques as a pre-processing layer. \\clearpage"
    },
    "0410/astro-ph0410086_arXiv.txt": {
        "abstract": "The spatial distribution of Kuiper belt objects (KBOs) in 2:1 exterior resonance with Neptune constrains that planet's migration history.  Numerical simulations demonstrate that fast planetary migration generates a larger population of KBOs trailing rather than leading Neptune in orbital longitude.  This asymmetry corresponds to a greater proportion of objects caught into asymmetric resonance such that their resonance angles, $\\phi$, librate about values $>\\pi$ (trailing) as opposed to $<\\pi$ (leading).  We provide, for the first time, an explanation of this phenomenon, using physical, analytic, and semi-analytic arguments.  Central to our understanding is how planetary migration shifts the equilibrium points of the superposed direct and indirect potentials.  Symmetric libration, in which $\\phi$ librates about $\\sim$$\\pi$, precedes capture into asymmetric resonance.  As a particle transitions from symmetric to asymmetric libration, if $\\phi$ exceeds its value, $\\psi$, at the unstable point of asymmetric resonance, then the particle is caught into trailing resonance, while if $\\phi<\\psi$, the particle is caught into leading resonance.  The probability that the KBO is caught into trailing resonance is determined by the fraction of time it spends with $\\phi>\\psi$ while in symmetric libration. This fractional time increases with faster migration because migration not only shifts $\\psi$ to values $<\\pi$, but also shifts the stable point of symmetric libration to values $>\\pi$.  Smaller eccentricities prior to capture strengthen the effect of these shifts.  Large capture asymmetries appear for exponential timescales of migration, $\\tau$, shorter than $\\sim$$10^7$ yr. The observed distribution of 2:1 KBOs (2 trailing and 7 leading) excludes $\\tau\\leq10^6$ yr with 99.65\\% confidence. ",
        "introduction": "\\label{sec-intro}  Transfers of energy and angular momentum between disks and bodies embedded within them can explain a variety of observed dynamical architectures. For example, the tight orbits of close-in, extra-solar, Jovian-mass planets (``hot Jupiters'') are thought to be a consequence of orbital migration driven by parent disks composed either of viscous gas (e.g., Ward 1997) or planetesimals (e.g., Murray et al.~1998). Capture of bodies into mean-motion resonances is a celebrated signature of migration (see, e.g., Peale 1986, and references therein), showcased recently by the pair of planets orbiting GJ 876 in a 2:1 resonance (Marcy et al.~2001; Lee \\& Peale 2002). Migration finds application in the solar system as well; it promises to explain the preponderance of highly eccentric Kuiper belt objects (KBOs) trapped in mean-motion resonances with Neptune (Malhotra 1995; Chiang et al.~2003ab). The outermost planet in our solar system might have migrated several AUs outward by gravitationally scattering planetesimals (Fernandez \\& Ip 1984; Hahn \\& Malhotra 1999), thereby sweeping its exterior resonances across the primordial Kuiper belt and filling them with KBOs (Malhotra 1995; Chiang \\& Jordan 2002).  Notwithstanding the numerous appeals to disk-driven migration, it has been difficult to dispel all doubt regarding its relevance. Worries are perhaps most naggingly persistent in the Kuiper belt, where the abundance of resonant KBOs could, in principle, be explained by the following stability argument. One imagines that the entire belt was once indiscriminately, dynamically excited to large eccentricities and inclinations by one or more massive perturbers---perhaps Neptune itself, during that planet's high eccentricity phase (Goldreich, Lithwick, \\& Sari 2004; see also Thommes, Duncan, \\& Levison 1999, 2002)---and that only KBOs fortunate enough to be kicked into resonances enjoyed the phase protection that permits survival for the age of the solar system.  Chiang \\& Jordan (2002, hereafter CJ) propose one test of the migration hypothesis that helps to break degeneracies of interpretation. They find that in scenarios where Neptune's migration is comparatively fast---occurring on timescales shorter than $10^7$ yr, for their chosen initial conditions---more objects are caught into the 2:1 resonance with mean libration angles greater than $\\pi$. This asymmetry in the distribution of libration angles manifests itself as an asymmetry on the sky: at a given epoch, more 2:1 resonant objects will appear at longitudes trailing, rather than leading, Neptune's. This is a prediction of the migration model that can be tested observationally by wide-angle, astrometric surveys such as, e.g., Pan-STARRS (Panoramic Survey Telescope \\& Rapid Response System).  Under either the alternative stability hypothesis or the hypothesis that Neptune's migration occurred over timescales longer than $10^7$ yr, we would expect equal populations trailing and leading Neptune.  The capture asymmetry discovered by CJ---witnessed also in numerical simulations by Wyatt (2003) in the context of debris disks molded by migrating extra-solar planets---arose from a purely numerical orbital integration. In this paper, we explain the capture asymmetry on physical grounds, using a variety of qualitative, analytic, and numerical descriptions. In uncovering the workings of the capture asymmetry on a step-by-step, mechanistic basis, we gain insight into its applicability to observations. Our focus is on the Kuiper belt, but the setting is clearly generalizable.  Terminology in this field can be confusing; we lay down some definitions here. The exterior 2:1 resonance is distinguished in offering the possibility of multiple stable points for the resonance angle, \\beq\\label{eqn-phi} \\phi \\equiv 2\\lambda - \\lambda_{\\rm N}- \\pomega \\,\\, , \\eeq where $\\lambda$ and $\\lambda_{\\rm N}$ are the mean longitudes of the test particle (KBO) and the planet (Neptune), respectively, and $\\pomega$ is the longitude of pericenter of the particle.\\footnote{We label $p$:$q$ resonances such that $p>q$ refers to exterior resonances (outside the perturber) and $p<q$ refers to interior resonances (inside the perturber). This choice runs counter to convention, but seems commonplace in the Kuiper belt literature.} At moderate eccentricities of the particle, the angle $\\phi$ can librate (oscillate with bounded amplitude) about a range of values.  Particles whose libration centers do not equal $\\pi$ are said to be in ``asymmetric resonance'' or ``asymmetric libration'' (Message 1958; Frangakis 1973; Beaug\\'{e} 1994; Malhotra 1996; Winter \\& Murray 1997; Pan \\& Sari 2004, and references therein). The usual libration of $\\phi$ about $\\pi$ is called ``symmetric resonance.'' We explain the physical origins of symmetric and asymmetric resonance in \\textsection\\ref{sec-libration}. We refer to the migration-induced preference for objects to be caught into asymmetric resonance such that they librate about angles greater than $\\pi$ as ``asymmetric capture.''  Note that asymmetric capture does not refer merely to capture into asymmetric resonance. We explain the mechanics underlying asymmetric capture in \\textsection\\ref{sec-capture}. Connections to observations are made in \\textsection\\ref{sec-observations}, and a summary and outlook for future work are presented in \\textsection\\ref{sec-summary}. ",
        "conclusions": "\\label{sec-summary}  Orbital migration of bodies embedded in disks can explain observed structures in systems ranging from planetary rings to extra-solar planets. As a body migrates, it can capture others into mean-motion resonances, leaving a potentially lasting signature of its migration.  In our solar system, numerous Kuiper belt objects (KBOs) exist in mean-motion resonance with Neptune, suggesting that Neptune may have migrated outwards by several AUs (Malhotra 1995; Chiang et al.~2003ab).  Chiang \\& Jordan (2002, CJ) discover by numerical simulation that the subset of KBOs inhabiting the 2:1 resonance furnishes a unique probe of Neptune's migration history.  The resonant angle, $\\phi$, of such KBOs can librate about values equal to $\\pi$ (symmetric resonance), greater than $\\pi$ (trailing asymmetric resonance), or less than $\\pi$ (leading asymmetric resonance). These authors find that if Neptune's migration occurred on timescales shorter than $10^7$ years (for their assumed initial conditions), more 2:1 resonant KBOs would have been captured into trailing resonance than into leading resonance; as a consequence, more KBOs would be discovered today at longitudes trailing Neptune than at longitudes leading it.  For longer timescales of migration, more equal numbers of KBOs would be found leading and trailing Neptune.  Observational confirmation of more trailing than leading KBOs in 2:1 resonance would constitute strong evidence in favor of the (fast) migration hypothesis; while current tallies are more consistent with equal populations, a definitive census should be possible with the upcoming Pan-STARRS survey.  We have explained the physical origin of this capture asymmetry within the context of the circular, restricted, three-body problem.  Asymmetric libration in the 2:1 resonance occurs as a result of the superposition of the direct and indirect components of the planet's perturbing acceleration (see also Frangakis 1973; Pan \\& Sari 2004).  The direct component, produced by the direct gravitational attraction between Neptune and the KBO, serves to accelerate the longitude of conjunction ($\\phi$) toward $\\pi$.  The indirect component arises from changes in the acceleration of the KBO by the Sun brought about by the latter's reflex motion induced by Neptune. This acceleration, whose effect on the KBO's orbit depends only on the difference between the true longitudes of Neptune and of the KBO, accelerates $\\phi$ toward $0$ over a synodic period.  Thus, the indirect and direct perturbations can counter-balance each other to produce asymmetric libration of $\\phi$ about angles intermediate between $0$ and $\\pi$, and intermediate between $\\pi$ and $2\\pi$.  When employing a literal series expansion of the disturbing potential, the inclusion of an inappropriate number of terms can cause spurious asymmetric resonances to appear.  For the 2:1 resonance, asymmetric libration is represented qualitatively well by a literal expansion to second order in $e$ but not by expansions to first, third, or fourth order.  The second-order expansion should be regarded as nothing more than a useful fitting formula for the true potential.  Particles caught by an outwardly migrating 2:1 resonance fall first into symmetric libration before evolving to asymmetric libration.  At the time of transition from symmetric to asymmetric resonance, if the value of $\\phi$ for a particle is larger than a critical angle $\\psi$, then the particle will be caught into the trailing island of asymmetric libration, while if $\\phi<\\psi$, it will be caught into the leading island.  The relative probability that a particle is captured into the trailing rather than the leading island is determined by the fraction of its time spent during symmetric libration at $\\phi$'s greater than and less than $\\psi$.  Three factors determine this fractional time spent: \\begin{enumerate} \\item Migration-induced shifts of the stable and unstable equilibrium points of the resonant potential. Planetary migration shifts the stable points for symmetric and asymmetric resonance to larger values.  For asymmetric resonance, migration shifts the unstable point formerly at $\\pi$, to a smaller value, $\\psi<\\pi$.  Analytic theory informs us that the angular offsets of the stable point of symmetric libration and of the unstable point of asymmetric libration vary inversely with migration timescale, $T_{\\rm m} = a\\subN/\\dot a\\subN$. The greater the shifts, the more likely it is that particles spend more time in symmetric resonance at $\\phi>\\psi$ and are captured into the trailing island.  If the sum of both shifts exceeds the amplitude of symmetric libration---true if $T_{\\rm m} \\lesssim T_{\\rm m, crit} \\sim 10^7$ yr---preference for capture into the trailing island is overwhelming. \\item Initial particle eccentricity. The influence of migration-induced shifts is greatest when the amplitude of symmetric libration is smallest; the latter demands small particle eccentricities in circulation prior to capture. Analytic theory also tells us that the magnitudes of the shifts themselves grow with smaller eccentricities. \\item Uneven libration rates during symmetric libration. The relative time spent is also affected by differences in $\\dot\\phi$ over one period of symmetric libration.  The average libration rate at $\\phi < \\psi$ is lower than at $\\phi > \\psi$ because the close juxtaposition of two equilibrium points at $\\phi < \\psi$ flattens the potential there. The difficulty with which a particle traverses the turning point at $\\phi < \\psi$ counteracts the effect of migration-induced shifts and can even reverse the usual sense of the capture asymmetry for large $T_{\\rm m}$, but not in a way that gives the leading island more than a 2-to-1 advantage over the trailing island in attracting occupants. \\end{enumerate}  We confirm the results of CJ that the migration timescales necessary to capture more particles into the trailing than the leading island are between about $10^6$ and $10^7$ yr. For $T_{\\rm m}$ shorter than $10^6$ yr, the resonance becomes increasingly unable to capture particles at all.  For a realistic migration prescription (e.g., an exponential prescription such as Equation \\ref{eqn-exponential}), Neptune's migration rate decreases with time, so particles beginning at smaller semi-major axes are captured while Neptune is migrating more quickly and are more likely to exhibit asymmetric capture.  The final asymmetry in the longitudes of the particles is sensitive to their initial semi-major axes and, more significantly, their initial eccentricities.  Asymmetric libration is not unique to the 2:1 resonance (e.g., Message 1958; Frangakis 1973; Beaug\\'e 1994; Malhotra 1996; Winter \\& Murray 1997; Pan \\& Sari 2004).  In a $p$:1 exterior resonance, the resonant angle, \\beq \\phi = p\\lambda - \\lambda\\subN - (p-1)\\pomega \\,\\, , \\eeq can librate asymmetrically as well.  For all other resonances, including interior resonances (e.g., 1:2), asymmetric resonance does not occur.  Asymmetric libration of the type exhibited by the 2:1 resonance requires a non-zero indirect acceleration.  We prove below that the indirect acceleration always averages to zero over one synodic period for a $p$:$q$ resonance except when $q=1$, where $p>q$ refers to an exterior resonance and $p<q$ to an interior resonance ($p\\neq q$ are relatively prime, positive integers).  The order of the resonance is $s=|p-q|$, and the resonant angle is $\\phi = p\\lambda -q\\lambda\\subN - (p-q)\\pomega$. Our proof is essentially identical to that of Frangakis (1973), and derived independently; we offer our expanded version for ease of reference.  The azimuthal component of the indirect acceleration of the KBO by Neptune equals $-(Gm\\subN/a\\subN^2)\\sin\\Delta\\theta$, where $\\Delta\\theta \\equiv \\lambda\\subN - f - \\pomega$ is the angle between the true longitudes of the planet and of the particle.  The proof for the radial component reads completely analogously.  Over one synodic period, the tangential acceleration time-integrates to \\begin{align}\\label{eqn-torqueintegral} \\langle T\\rangle &= -\\frac{Gm\\subN}{a\\subN^2} \\oint \\sin(\\lambda\\subN-f-\\pomega) dt \\nonumber \\\\ & = -\\frac{Gm\\subN}{a\\subN^2} \\int_0^{2\\pi s} \\frac{\\sin(\\Delta\\theta)}{n\\subN-\\dot f(\\Delta\\theta)-\\dot\\pomega} d(\\Delta\\theta) \\,\\, . \\end{align} After one full period of the particle, the particle returns to the same true longitude while the planet increases its true longitude by $2\\pi s/q$, increasing $\\Delta\\theta$ by $2\\pi s/q$.  It follows that $\\dot f(\\Delta\\theta)$ is periodic in $\\Delta\\theta$ with period $2\\pi s/q$, as is any function of $\\dot f$, including \\beq F(\\Delta\\theta) \\equiv \\left[n\\subN-\\dot f(\\Delta\\theta)-\\dot\\pomega\\right]^{-1} \\,\\, , \\eeq which may be Fourier-decomposed as \\beq F(\\Delta\\theta) = \\frac{b_0}{2} + \\sum_{j=1}^\\infty b_j\\cos\\left(\\frac{j\\pi}{L}\\Delta\\theta\\right) +  \\sum_{j=1}^\\infty c_j\\sin\\left(\\frac{j\\pi}{L}\\Delta\\theta\\right) \\,\\, , \\eeq where $2L \\equiv 2\\pi s/q$. We have neglected the tiny variation of $\\dot\\pomega \\ll n\\subN$ over a synodic period. Plugging this expression for $F$ into Equation \\ref{eqn-torqueintegral}, we find \\begin{align} \\langle T\\rangle &= -\\frac{Gm\\subN}{a\\subN^2}\\int_0^{2\\pi s}F(\\Delta\\theta)\\sin(\\Delta\\theta)d(\\Delta\\theta) \\nonumber \\\\ &= -\\frac{Gm\\subN}{a\\subN^2}\\left[ s\\sum_{j=1}^\\infty b_j \\int_0^{2\\pi} \\cos(jqx)\\sin(sx)dx + \\right. \\nonumber \\\\ &\\qquad\\qquad\\quad \\left. s\\sum_{j=1}^\\infty c_j \\int_0^{2\\pi} \\sin(jqx)\\sin(sx)dx \\right] \\nonumber \\\\ &= -\\frac{Gm\\subN}{a\\subN^2}s\\pi\\sum_{j=1}^\\infty c_j\\delta_{(jq)s}  \\,\\, . \\end{align} Thus, $\\langle T\\rangle\\neq 0$ only when $jq=s=|p-q|$ for some $j$.  This is satisfied only when $q=1$ and $j=|p-q|$ under our assumption that $p$ and $q$ are relatively prime.\\footnote{When $q$ is even, one can see that the indirect term integrates to zero by appealing to diagrams analogous to Figure \\ref{fig-indirect21} and noting symmetry over the synodic period.} We conclude that asymmetric libration, of the kind exhibited by the 2:1 resonance, can only exist for exterior $p$:1 resonances.\\footnote{The only resonance for which our proof does not apply is the 1:1 resonance. But it is clear that the indirect acceleration does not time-average to zero in tadpole and horseshoe orbits, which may be regarded as asymmetric and symmetric librations, respectively. See Frangakis (1973) and Pan \\& Sari (2004).}  Future theoretical work should focus on determining the degree of stochasticity in migration driven by scattering planetesimals, and on the size spectrum of planetesimals required for resonance capture to proceed smoothly. We have made preliminary estimates which suggest that for a Neptune-mass disk in the vicinity of Neptune, the disk mass must be concentrated in planetesimals whose sizes do not exceed $\\sim$100 km for the theory that we have described in this paper to apply. Reasons to imagine that planetesimal sizes were small can be found in the review of planet formation by Goldreich et al.~(2004)."
    },
    "0410/physics0410134_arXiv.txt": {
        "abstract": "Rapidly rotating spherical kinematic dynamos are computed using the combination of a quasi geostrophic (QG) model for the velocity field and a classical spectral 3D code for the magnetic field. On one hand, the QG flow is computed in the equatorial plane of a sphere and corresponds to Rossby wave instabilities of a geostrophic internal shear layer produced by differential rotation. On the other hand, the induction equation is computed in the full sphere after a continuation of the QG flow along the rotation axis. Differential rotation and Rossby-wave propagation are the key ingredients of the dynamo process which can be interpreted in terms of $\\alpha\\Omega$ dynamo. Taking into account the quasi geostrophy of the velocity field to increase its time and space resolution enables us to exhibit numerical dynamos with very low Ekman (rapidly rotating) and Prandtl numbers (liquid metals) which are asymptotically relevant to model planetary core dynamos. ",
        "introduction": "The magnetic field of the Earth is produced by a dynamo effect in the metallic liquid core of our rotating planet. Many efforts have been made successfully in the last decade to describe the mechanism of a self induced magnetic field either with experimental models \\cite{gai01,sti01,car04} or numerical simulations \\cite{gla95,kag97,kua97,chr99}. Both approaches have limitations. No experiment has been done in rotation while rotation is seen as a key ingredient by geophysicists to explain the geometry and the amplitude of the geomagnetic field \\cite{gub87}. All numerical models \\cite{dor00,kon02} have introduced the Coriolis force in solving the Navier-Stokes equation and the quasi geostrophy (two dimensionality imposed by the Taylor Proudman theorem \\cite{gre68}) of the flow participates in the generation of the magnetic field. Thermal convective vortices aligned with the rotation axis are associated to surface patches of magnetic field \\cite{chr99} and spatio-temporal behaviors of magnetic and vorticity field are similar. This effect is a direct consequence of the prescribed magnetic Prandtl number ($P_m = \\nu/\\eta$, where $\\nu$ is the kinematic viscosity and $\\eta$ the magnetic diffusivity) in the simulations. The current computer capacities limit the computation to magnetic Prandtl number of the order of unity \\cite{dor00} while liquid metals exhibit magnetic Prandtl number lower than $10^{-5}$, even in the planetary core conditions \\cite{poi94}.  In this paper, we propose an approach that aims at computing very low magnetic Prandtl number dynamos taking advantage of the quasi-geostrophic behavior of the velocity field. For very low Ekman number ($ E = \\nu/\\Omega R^2$, where $\\Omega$ is the rotation rate of the spherical container, and $R$ its radius), a quasi-geostrophic (QG) approach models correctly the flow in a rapidly rotating sphere \\cite{bus70,car94}.  It consists of the integration of the flow equations along the rotation axis. Even if the numerical resolution is done with a stream function in the equatorial plane (2D), the top and bottom boundary conditions are taken into account through slope ($\\beta$) and Ekman pumping effects. In the context of the study of thermal convection in rapidly rotating spherical shells, Aubert et al. \\cite{aub03} have compared successfully their QG results with 3D calculations \\cite{dor04} and experimental measurement \\cite{aub01}. Low value of $P_m$ may imply a separation in term of scales and frequencies, between the velocity and magnetic fields in a metallic dynamo. This idea has already been applied to kinematic dynamo computations at low $P_m$ \\cite{pon04}.  In this work, we compute the QG flow in the equatorial plane with a fine spatio-temporal resolution and the velocity is extrapolated to a coarse 3D spherical grid where the induction equation is solved.   In order to demonstrate the validity of this approach, we have decided to apply it to a simple case. Instead of a thermal convective flow for which heat transport has to be modeled, we model the instabilities of an internal geostrophic shear layer. This layer, known as the Stewartson layer, is produced by a differentially rotating inner core in a rotating sphere and consists of two nested viscous shear layers \\cite{ste66,dor98}. For a critical Ro number ($ Ro = \\Delta \\Omega / \\Omega$, where $\\Delta \\Omega$ is the differential rate of rotation of the inner core), the Stewartson layer becomes unstable \\cite{hol03} and generates Rossby waves \\cite{sch04}.  As we will show in this paper, such kind of flow can generate and sustain a magnetic field. The QG-model allows us to compute dynamos at very low Ekman (down to $10^{-8}$) and Prandtl numbers (as low as $3\\,10^{-3}$). ",
        "conclusions": "In summary, we have computed a quasi-geostrophic dynamo based on a Stewartson shear layer flow. The scale separation approach works because the small scales of the flow in our rotating sphere are negligible (very steep kinetic energy spectrum $E(k) \\sim k^{-5}$). Our preliminary results may be interpreted in terms of $\\alpha\\Omega$ dynamo. The $\\Omega$ effect is done by the shear of the Stewartson layer itself whereas the $\\alpha$ effect is produced by vortices associated with the Rossby waves due to the instability of the shear layer. These understandings are very encouraging for our on-going experimental modeling of the geodynamo. As described in Cardin et al. \\cite{car02}, we are building a spherical Couette experiment using liquid sodium which may validate and enlarge our present numerical findings.  For the first time, we have computed a spherical dynamo with a very low magnetic Prandtl number ($<10^{-2}$) and a very low Ekman number ($10^{-8}$) (corresponding to a very high Reynolds number $Re>10^{6}$). Even if our dimensionless parameters stay far away from parameters of planetary cores, our calculations use dimensionless numbers which are in the correct asymptotic regime for the modeling of the geodynamo. The key ingredients of our approach is to take into account a specific property of the rotating fluid (QG) which allows us to use a 2D model to compute the flow evolution, and the separation of scales between the magnetic field and the velocity field, allowing us to use a coarse 3D mesh for the magnetic field.  We also showed that in the case studied in this paper, the mean flow or the static flow fails to produce a dynamo while the fully resolved time-dependent flow succeeds. Indeed, the time evolution of the flow and the $\\beta$ effect are key ingredients for dynamo action in our models, while the Ekman pumping can be neglected without losing the dynamo effect.  The next step will be to add the Lorentz force in the QG equation to compute saturated dynamos. One of the difficulty is to compute the action of the large magnetic field on the small scale motions of the fluid. Preliminary results are encouraging and exhibit saturated dynamos very close to the kinematic dynamos described here.  A quasi geostrophic approach could also be used to build thermal convective dynamos. A zonal geostrophic flow is produced by the Reynolds stress of the thermal columns \\cite{aub01,chr02} but its amplitude is much lower compared to the differential rotation imposed in the Stewartson problem. Would it be enough to start an Stewartson dynamo type? for what forcing? Would it work for very low Ekman and magnetic Prandtl numbers?  \\vspace{1cm} {\\it Aknowledgements}: Calculations were performed at SCCI (Observatoire de Grenoble) and at IDRIS (CNRS). This work has been supported by the programme \"DyETI\" of CNRS/INSU. We wish to thank Dominique Jault and Henri-Claude Nataf for very useful comments."
    },
    "0410/astro-ph0410279_arXiv.txt": {
        "abstract": "Recently it was pointed out that nonmodally (transiently and/or adiabatically) pre-amplified waves in shear flows, undergoing subsequent viscous damping, can ultimately heat the ambient flow. The key ingredient of this process is the ability of waves to grow, by extracting energy from the spatially inhomogeneous mean flow. In this paper we examine this mechanism in the context of the solar coronal plasma flows. ``Self-heating\" (SH) processes are examined when both viscous damping and magnetic resistivity are at work. We show that if the plasma viscosity is in the favorable range of values the asymptotic SH rate in these flows can be quite substantial. ",
        "introduction": "Astrophysical plasma flows are often complex, inhomogeneous, dynamic systems hosting different kinds of collective phenomena. Sometimes interactions between collective phenomena and ambient flows feature a highly sophisticated level of complexity. In general, shear flows (SF) are well-known examples of behavioral complexity - both in neutral and charged fluids, both for terrestrial and astrophysical cases and both on the level of experiments (observations) and theory (simulations) - SF display a wide spectrum of ``shear-induced\" phenomena. In particular, astrophysical SF are gradually becoming a popular subject of research because it becomes increasingly more clear that most of the cosmic plasmas are flowing with spatially inhomogeneous rates and in most of these flows shear-induced processes may leave a considerable imprint on the observational appearance of the related astronomical objects.  A separate class of shear-induced phenomena is related with the non-self-adjointness of the linear operators governing the dynamics of SF systems (Trefethen et al. 1993). Fluctuations in a ``parent\" SF obey ``non-Hermitian\" equations and, therefore, their evolution, their interaction with the ambient flows can not be molded within the standard Hamiltonian formalism. The most striking novelty is the appearance of new, shear-induced nonperiodic modes - so-called Kelvin or shear vortices. Interactions between the modes and between the flow and the modes bring into the game a {\\it nonmodal complexity}: these ``non-normal\" modes show a transient increase in amplitude, they are coupled and are transformed into each other, they may feature beating phenomena and may become the subjects of different shear instabilities.  Recently it was pointed out that SF may host yet another interesting non-modal phenomenon related with the combined presence of nonmodality and dissipative processes in the flows. It can formally be considered as a three-phase process (Rogava 2004):  \\begin{enumerate}  \\item waves and/or vortices get \\emph{excited }within a flow;  \\item they \\emph{amplify} nonmodally, due to the presence of the shear flow, extracting a part of the flow kinetic energy;  \\item high-amplitude waves and/or transiently amplified vortices undergo a viscous decay and/or magnetic diffusion and give, in the form of heat, a part of their energy back to the flow.  \\end{enumerate}  As a result of this three-step process the fluctuation gives back to the flow {\\it more} thermal energy than it had at the moment of its excitation. Repeated continuously, throughout the whole volume of the flow, these elementary processes should lead to a net heating of the flowing plasma. The process was named \\emph{``self-heating\"} (SH), because what actually happens is that nonmodality manages to transfer a part of the flow's regular kinetic energy into a thermal energy; i.e., the heating happens conservatively without any outside source of energy, the flow manages to heat itself.  The possibility of this process was originally surmised in (Rogava 2003a) and later it was examined and specified in detail for the relatively simple case of a neutral fluid SF and sound waves sustained by it (Rogava 2004). It was shown that in the hydrodynamic situation - even when a mixture of sound waves and Kelvin vortices is propagating in a simple, plane-parallel Couette-like SF and, therefore, acoustic waves are able to amplify only linearly, extracting the kinetic energy of the ambient flow - the self-heating rate can be quite high: fluctuations are able to give back to the flow several times more energy than they have initially had.  The purpose of this paper is to explore whether the SH can be efficient in magnetized plasma MHD flows. In this particular study we consider only the case of a plane-parallel flow with a linear velocity profile. However, in order to lay the foundation for a further, more general, consideration of flows with kinematic complexity we derive our equations for the case of swirling (helical) flows, following the approach developed in (Rogava et al. 2003a and 2003b) (henceforth referred to as R1 and R2) and consider the case of the plane-parallel flow as a particular example.  Our results show that both in the incompressible  (when the flow sustains only Alfv\\'en waves) and in the compressible (full spectrum of MHD waves) cases the self-heating is even more efficient than it was in the hydrodynamic case. Our numerical simulations revealed that the values of the dimensionless asymptotic SH rate $\\Xi_{\\infty}$ can be of the order of several tens. This allows us to argue that the SH may easily be one of the robust physical mechanisms contributing to the heating of the solar corona through the viscous and/or resistive damping of nonmodally pre-amplified waves.  It should be noted that nonmodal SH, being efficient \\emph{per se} possesses another attractive feature in the context of its possible relevance for coronal heating: all existing wave heating scenarios face common problem: how is the energy transferred from longer length-scales to shorter ones, where dissipative effects are significant?! SH phenomena, being nonmodal, involve the shear-induced drift of the wave number vector ${\\bf k}(t)$, which provides a natural, flow-related mechanism for the gradual decrease of the mode's length-scale. This process, being linear in plane-parallel SF, may have an exponentially fast nature in geometrically and kinematically more sophisticated cases. Therefore, we envisage that kinematically complex flow patterns, such as solar spicules, macrospicules and tornados, magnetic plumes, might host much more efficient, sometimes even explosive, SH events. ",
        "conclusions": "The purpose of the present study was to clarify whether the shear-induced amplification of MHD waves coupled with viscous and magnetic-resistive  dissipation may lead to a significant \\emph{self-heating} of the SF. In particular, our aim was to see whether this mechanism could be efficient  for solar coronal plasma flows and could, arguably, contribute to the heating of the solar corona.  Since the variety of solar plasma flows is quite wide, both in terms of geometry and kinematics, the \\emph{local} formalism employed and developed within this study is apt for an arbitrary velocity pattern with the condition that the involved MHD waves are having length-scales, ${\\l }$, sufficiently smaller than the linear length-scale,  ${\\cal L}$, of a ``parent\" flow pattern. Basically, our consideration follows Rogava et al. (2003a, 2003b) with addition of dissipative effects related to the presence of viscosity and magnetic resistivity.  In this paper we have examined only the case of the simplest velocity pattern with plane-parallel and linear velocity profile. Both incompressible and compressible limits were investigated. It was found that:  \\begin{enumerate}  \\item Incompressible, Alfv\\'enic, perturbations, in the ideal MHD limit, are known to undergo an algebraic instability (``transient increase\") (Chagelishvili et al. 1993; Rogava et al. 1996, Rogava et al. 2003a), which can be quite strong: under favourable conditions the total energy of a fluctuation may increase several hundred times. We found that when dissipation is taken into account these high-amplitude pre-amplified Alfv\\'en waves get damped and give their energy back to the ambient flow. The resulting SH is quite substantial. As it could be expected the asymptotic self-heating rate in this case is quite large: it increases with the decrease of the Alfv\\'en velocity and in the cases considered here could reach a value of the order of several tens. This happened to be a case for waves with the characteristic periods of the order of few minutes.  \\item In the compressible case we considered the evolution of an initially excited Alfv\\'en wave (as an example we considered Alfv\\'en mode with period of the order of 30 seconds) in a plane-parallel flow pattern with the value of the plasma-$\\beta$ of the order of $10^{-2}$, typical for the coronal environment. In this case Alfv\\'en waves  and fast magnetosonic waves are coupled, which leads to the transformation of initially Alfv\\'enic fluctuations into the fast ones \\footnote{The latter wave, being of the acoustic nature, dissipates quite similarly to plain sound waves, recently shown to be instrumental in the SH of compressible fluid shear flows (Rogava 2004).} (Rogava et al. 2000). We found (see Fig.~\\ref{fig2c}) that this process, coupled with the presence of viscous and resistive dissipation, ensures the SH of the ``parent\" flow with the asymptotic rate of the order of several tens.  \\end{enumerate}  These results are quite encouraging. We see that even the simplest kinds of magnetized plasma SF's are able to heat themselves via the agency of nonmodally pre-amplified Alfv\\'en waves. It is both tempting and reasonable to surmise that in flow patterns with a higher degree of kinematic complexity, where nonmodal processes of energy exchange between flows and waves have a more intense nature (Mahajan and Rogava, 1999), the resulting SH can be considerably higher! In certain cases \\footnote{For instance, in flows with a helical mode of plasma motion (Rogava et al. 2003a, 2003b)} individual flow patterns may tend to self-destruction, undergoing self-imposed, catastrophically fast SH of an eruptive, explosive nature. Further studies in this vein are currently being carried out and the results will be reported elsewhere.  The astrophysically relevant conclusion of the present study is that nonmodal self-heating, even for relatively simple kinds of flows, parallel to the magnetic field and with shear profiles across the field, can pay a significant contribution in the (still poorly understood) phenomenon of solar coronal heating."
    },
    "0410/astro-ph0410035_arXiv.txt": {
        "abstract": "I review results from {\\sl Chandra} observations of nearby normal edge-on galaxies (Sd to Sa types). These galaxies have a broad range of star formation rate, but none of them is dominated by a nuclear starburst. The galaxies are all in directions of low Galactic foreground absorption ($N_{HI} \\la 4 \\times 10^{20} {\\rm~cm^{-2}}$). Extra-Planar diffuse soft X-ray emission is detected unambiguously from all the galaxies, except for N4244 (Sd), which is low in both the stellar mass and the star formation rate. The thermal nature of the X-ray-emitting gas is well established, although its chemical and ionization states remain largely uncertain. The X-ray luminosity of the gas is proportional to the star formation rate and to the stellar mass of the galaxies. But the luminosity accounts for at most a few percent of the expected supernova mechanical energy input. Therefore, there is a ``missing'' energy problem for spiral galaxies. Much of the energy in late-type spirals may be converted and radiated in lower energy bands. But early-type ones most likely have outflows, which are powered primarily by Type Ia supernovae in galactic bulges. These galactic outflows may strongly affect both the dynamics and cooling of the intergalactic gas accretion, hence the evolution of the galaxies. ",
        "introduction": "It has long been theorized that a major, possibly dominant, phase of the interstellar medium (ISM) is gas at temperature $T \\sim 10^6$ K in virtually all galaxies. This hot ISM is thought to be created and maintained primarily by supernova explosions (SNe; e.g., McKee \\& Ostriker 1977). If sufficiently energetic, the hot gas is expected to flow outward, creating a large-scale gaseous corona (e.g., Spitzer 1956; Bregman 1980a, Norman \\& Ikeuchi 1989) or even escaping as a galactic wind (e.g., Bregman 1980b). Therefore, the study of extra-planar hot gas is fundamentally important to our understanding of the mass, energy, and chemical evolution of a disk galaxy such as our own.  Observationally, although an understanding of the large-scale properties of the hot ISM in our own Galaxy is still elusive, the presence of large amounts of diffuse hot gas in nearby edge-on disk galaxies has been established. Much of this progress has been made with recent {\\sl Chandra} observations, which provide high-resolution and panoramic views of extra-planar hot gas and its interaction with other galactic components (e.g., Wang et al. 2001, 2003; Strickland et al. 2004). The arc-second resolution of these observations, in particular, allows for a clean separation of point-like sources from the diffuse emission, a step critical for reliably determining both the content and physical condition of diffuse hot gas. Here I review key results from the observations and discuss their implications. ",
        "conclusions": "Observations}  \\begin{table}[!ht] \\caption{A Sample of Nearby Normal Disk Galaxies Observed with {\\sl Chandra}} \\smallskip {\\scriptsize \\begin{center} \\begin{tabular}{lccccccccc} \\tableline \\noalign{\\smallskip} Galaxy & Hubble & $D$ & Incl. &$L_{IR}$ & $L_{60}/L_{100}$& K$_{tot}$ & Env. & Exp. &$L_x$\\\\ Name   & Type & (Mpc) & (deg)&($10^{10} L_\\odot$)& && & (ks)&($10^{39} L_\\odot$)\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} N4244 & Sd     & 3.6 & 85 &0.02 &0.26 &7.72& I & 60&$< 0.03$ \\\\ N4631 & Sd     & 7.5 & 85 &1.74 &0.40 &6.47& C & 60&2 \\\\ N3556 & Scd    & 14  & 80 &3.70 &0.42 &7.04& I & 60&3 \\\\ N3877 & Sc     & 17  & 76 &3.39 &0.33 &7.75& C &122&1 \\\\ N5775 & Sc     & 25  & 86 &5.25 &0.46 &7.76& C & 46&9 \\\\ N4565 & Sb     & 13  & 87 &0.77 &0.22 &6.06& C & 60&0.3 \\\\ N4594 & Sa     & 8.9 & 84 &0.20 &0.24 &4.96& C & 19&3 \\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} \\end{center} K$_{tot}$  - K band magnitude from the 2MASS survey;\\par Env. -- galaxy environment: I - isolated; C - with galaxy companions;\\par Exp. - {\\sl Chandra} ACIS-S exposure time;\\par $L_x$ - Extra-planar diffuse 0.2-2 keV luminosity. } \\end{table}  Table 1 lists a sample of nearby normal edge-on galaxies observed with the {\\sl Chandra} ACIS-S. They span a broad range of morphological types and star forming properties as judged by the 60$\\micron$ to 100$\\micron$ intensity ratios ($L_{60}/L_{100}$; Table 1). But none of the sample galaxies is dominated by nuclear starburst or AGN activities. Both N4631 and N5775 are clearly disturbed by the interaction with their companions. Others seem to be rather isolated. N4244, an extremely quite and intrinsically low-surface brightness galaxy, presents little sign of diffuse X-ray emission (Strickland et al. 2004). The remaining six galaxies all show  strong evidence for large-scale extra-planar diffuse X-ray emission (Fig. 1). Quantitative results have been reported only for N4631 and N3556 (Wang et al. 2001, 2003). The analysis of the other galaxies is still ongoing, although crude estimates of the extra-planar diffuse X-ray luminosities are included in Table 1.  \\begin{figure}[!bth] \\begin{center} \\hbox{ \\psfig{figure=pfig1.ps,height=7.truein,angle=0.0,clip=} } \\end{center} \\caption{(a) {\\sl Chandra} ACIS-S intensity contours of nearby edge-on galaxies, overlaid on their optical B-band images. } \\end{figure}  \\begin{figure}[!bth] \\begin{center} \\hbox{ \\psfig{figure=pfig2a.ps,height=2.5truein,angle=0.0,clip=} \\psfig{figure=pfig2b.ps,height=2.3truein,angle=0.0,clip=} } \\end{center} \\caption{(a) {\\sl Chandra} ACIS-I intensity contours of the nuclear starburst galaxy M82, overlaid on an optical B-band image. The X-ray image is a mosaic of several existing observations;  CCD readout streaks have been subtracted. (b) VLA 1.4 GHz continuum intensity contours overlaid on the tri-color {\\sl Chandra} ACIS-S image of N4565. The colors represent the intensities in the three energy bands: 0.3-0.7 keV (red), 0.7-1.5 keV (green), and 1.5-7 keV (blue). } \\end{figure}  It is interesting to compare the diffuse X-ray morphology of the normal galaxies with that of nuclear starburst galaxies (e.g., Fig. 2a; see also Strickland et al. 2004). The late-type (Sd-Sc) galaxies tend to show diffuse X-ray emission on scales comparable to, or even larger than, the optical sizes. The distribution of the galaxy-wide diffuse X-ray emission is relatively smooth and is more extended than H$_{\\alpha}$-emitting materials, which indicates that the bulk of the X-ray emission arises in galactic coronae --- hot gas confined around the host galaxies. Interestingly, the nuclear starburst galaxy M82 (Fig. 2a) shows a simiarly large-scale diffuse X-ray-emitting halo, in addition to an apparent bipolar outflow from the galactic central region (Strickland et al. 2004). The X-ray morphology of the halo is more vertically elongated than those in Fig. 1. This elongation is likely caused by the large momentum of the nuclear outflow.  In regions close to the galactic disks of both normal and nuclear starburst galaxies, the X-ray emission appears to be rather filamentary and shows a strong correlation with extra-planar  H$_\\alpha$-emitting clouds, which are probably pushed out from the galactic disk (Wang et al. 2001, 2003; Strickland et al. 2004 and references therein). Much of the X-ray emission may arise from freshly shocked cool gas clouds of small filling factor by fast-moving, low density outflows.  The diffuse X-ray emission seems to be more concentrated toward the inner regions of the early-type (Sb-Sa) galaxies. This concentration may be partly  due to the fact that both N4565 and N4594 are very inactive in star formation. The gas heating by Type Ia SNe from the old stellar population becomes more important than core-collapsed SNe.  The diffuse X-ray spectrum can typically be characterized with a thermal plasma of a few times $10^6$ K. The metal abundances appear to be enhanced in the diffuse hot gas: O-like elements in late-type spirals and Fe-like elements in early-type spirals (Fig. 3). But in general, both the ionization and chemical states of the X-ray-emitting plasma are not well constrained. The estimate of the absolute metal abundances, for example, depends strongly on the assumed plasma emission model, e.g., the temperature distribution. The X-ray luminosity estimate of a low-surface brightness corona is sensitive to the exact background subtraction. Nevertheless, we find that the luminosity appears to be proportional to the star formation rate and to the stellar mass, as traced by the far-IR and K-band luminosities of the galaxies, respectively. But the diffuse X-ray luminosity typically accounts for less than a few percent of the expected SNe energy input in each galaxy.   \\begin{figure}[bth] \\begin{center} \\hbox{ \\psfig{figure=fig3a.eps,height=1.7in,angle=0,clip=} \\psfig{figure=fig3b.eps,height=1.7in,angle=0, clip=} } \\vspace{-1.0cm} \\end{center} \\caption{% (a) The {\\sl Chandra} ACIS-S spectra of the diffuse emission from the central corona region of N4631. The histogram represents the best-fit thermal plasma plus a power law (photon index = 1.5; accounting for $\\sim 10\\%$ counts), constrained with a joint fit to an integrated spectrum of discrete sources. The best-fit temperature is $\\sim 0.4$ keV, while the metal abundances are 0.8 and 0.2 solar for the O-like and Fe-like elements, respectively. (b) An ACIS-S spectrum of the extra-planar diffuse emission from N4594  and the best-fit galactic bulge-outflow model (histogram). Notice the strong and extended Fe-L bump at $\\sim 0.8 $ keV. } \\end{figure}"
    },
    "0410/gr-qc0410072_arXiv.txt": {
        "abstract": "Two mechanisms for nonlinear mode saturation of the r-mode in neutron stars have been suggested: the parametric instability mechanism involving a small number of modes and the formation of a nearly continuous Kolmogorov-type cascade. Using a network of oscillators constructed from the eigenmodes of a perfect fluid incompressible star, we investigate the transition between the two regimes numerically. Our network includes the 4995 inertial modes up to $n\\leq 30$ with 146,998 direct couplings to the r-mode and 1,306,999 couplings with detuning$< 0.002$ (out of a total of approximately $10^9$ possible couplings). The lowest parametric instability thresholds for a range of temperatures are calculated and it is found that the r-mode becomes unstable to modes with $13<n<15$. In the undriven, undamped, Hamiltonian version of the network the rate to achieve equipartition is found to be amplitude dependent, reminiscent of the Fermi-Pasta-Ulam problem. More realistic models driven unstable by gravitational radiation and damped by shear viscosity are explored next.  A range of damping rates, corresponding to temperatures $10^{6}$K to $10^9$K, is considered. Exponential growth of the r-mode is found to cease at small amplitudes $\\approx 10^{-4}$. For strongly damped, low temperature models, a few modes dominate the dynamics. The behavior of the r-mode is complicated, but its amplitude is still no larger than about $10^{-4}$ on average.  For high temperature, weakly damped models the r-mode feeds energy into a sea of oscillators that achieve approximate equipartition. In this case the r-mode amplitude settles to a value for which the rate to achieve equipartition is approximately the linear instability growth rate. ",
        "introduction": "\\label{ThreeMODE} The smallest non-linear system consists of three modes and one coupling. The undamped Hamiltonian system yields an analytic solution in terms of elliptic integrals \\cite{CarlsonCubic} as shown in Appendix~\\ref{AppendThreeHams}. The period of the amplitude oscillations  depends on the initial conditions of the system and can be accurately calculated numerically \\cite{NR}. For the case of a large  amplitude parent coupling to small amplitude daughters the period increases as the parent amplitude decreases. It was this dependence on initial amplitudes of the nonlinear oscillation frequency that was used  \\cite{JdB3} to explain the timescales involved in the catastrophic decay of the ``pumped up'' state of the r-mode observed in hydrodynamical  simulations \\cite{Ster2}.  There exists a minimum parent amplitude below which no oscillations in amplitude occur. This amplitude is the no-damping limit of the  parametric instability threshold discussed in the next paragraph.  Above this threshold very long oscillation periods are realized.  Accurately following these oscillations, which take place over thousands of periods of the oscillators involved, is a challenge to our numerical code, and was one of the tests used to verify the reliability of our integration scheme.    Study of the damped system introduces the concept of the parametric instability, which sets the minimum amplitude  the r-mode  has to reach before it can excite other modes in the system. No energy is transfered out of the driven mode until the instability threshold is reached. The phase space of the damped three mode system was thoroughly explored by Dimant \\cite{Dimant}  and Wersinger et. al \\cite{Ott}. The stationary solutions and parametric instability threshholds are compactly re-derived in Appendix~\\ref{parametricAPPENDIX}. The parametric instability threshold  for an unstable parent $\\alpha$ sharing excitation with two daughters $\\beta$ and $\\gamma$ is \\begin{equation}|c_\\alpha|^2 = \\frac{\\gamma_\\beta \\gamma_\\gamma}{4w_\\gamma w_\\beta \\kappa^2} \\left[1+\\left(\\frac{\\delta w}{\\gamma_\\gamma+\\gamma_\\beta}\\right)^2\\right]~. \\label{ParINSTB2} \\end{equation} Here $\\gamma_\\beta$, $\\gamma_\\gamma$ are the damping rates and $w_\\beta$, $w_\\gamma$ are the frequencies of the daughter modes, while $\\delta w = w_\\alpha - w_\\beta - w_\\gamma$ is the detuning.  Note that the instability threshold does not depend on the linear growth rate of the parent and is subject to the correct relationship between the signs of the mode frequencies (some triplets never go unstable as shown in Appendix~\\ref{parametricAPPENDIX}).   The parametric instability thresholds for the 146,998 direct couplings to the r-mode were calculated for a range of temperatures.  All modes with $n\\leq 30$ were included. The behavior of the lowest three thresholds as a function of temperature is shown in Figure~\\ref{ParametricInstThres} \\begin{figure*} \\includegraphics[width=\\picwidth]{ParametricInst.eps} \\caption{Temperature dependence of lowest 3 parametric instability thresholds of the r-mode to the inertial modes. } \\label{ParametricInstThres} \\end{figure*}  The graph begins with the temperature at which gravitational radiation first destabilizes the r-mode of a neutron star rotating at $\\Omega = 0.37\\sqrt{\\pi G \\rho}$, namely where the balance between the damping due to shear viscosity of the r-mode and the GR driving rate occurs. At this temperature all modes have strong viscous dissipation.   The r-mode has to attain a relatively large amplitude before the nonlinear term dominates the linear damping of the daughters (equation~\\eqref{ParINSTB2}). By contrast in a system where the daughters are less strongly damped  parametric instability occurs more readily.  As the damping is decreased the instability amplitudes decrease monotonically until the limiting value of $|\\delta w/4\\kappa \\sqrt{ w_\\beta w_\\gamma}|$ is reached.  The inset in the graph gives the conversion of the numbers labeling a particular coupling to the eigenmodes labeled according to degree, $n$, and order, $m$, of the Legendre functions that make up the eigenmodes' functional representation. The lowest parametric instability in the network we considered occurred for modes with reasonably high degree, namely $n=13$ and $n=14$. For low enough temperature, the shear damping is large  enough so that the three mode system can dissipate all the energy with which the r-mode is driven, stopping its growth.  However, as the damping decreases the r-mode is driven at a rate that cannot be halted by mere three mode coupling, and more modes become important. ",
        "conclusions": "We have studied the effects of second order nonlinear couplings among inertial modes of an incompressible star in order to understand how they may truncate the CFS instability of the $n=3$, $m=2$ r-mode. Although the system we study is not a realistic version of a neutron star, it illustrates certain general principles that apply to the nonlinear evolution of the instability.  As in our earlier studies \\cite{Katrin1}\\cite{Saturation}\\cite{JdB3}, we have adopted the philosophy that nonlinear effects can slow or stop the growth of the amplitude of the inertial modes of a neutron star so that large amplitudes are {\\it never} attained. Thus, we can employ an expansion of the fully nonlinear hydrodynamical equations of motion, keeping only the lowest order nonlinear terms . This philosophy can be justified in retrospect, since as shown in this paper nonlinear effects do indeed limit the dynamics to small amplitudes.  We may idealize further by restricting attention only to the inertial mode sector of perturbations of a (slowly) rotating neutron star. One reason we can do this is that these modes have frequencies that lie between $\\pm 2\\Omega$, which are generally far below the frequencies of the $p$ and $f$ modes of the star, which have frequencies $\\gtrsim\\sqrt{\\pi G \\rho}$ and, for small mode wavelengths, $\\gg \\sqrt{ \\pi G \\rho}$. Thus, we do not expect any instability involving the r-mode to excite the higher frequency modes efficiently. Moreover, nonlinear effects are magnified by near resonances even at small modal amplitudes. Such resonances are likeliest in the rather dense inertial mode sector. Indeed, one of the general principles governing the importance of nonlinearity in weakly nonlinear perturbation theory of the kind we present is that modes with small relative detunings dominate the dynamics.  In this paper, we have presented a sequence of calculations of increasing complexity in order to illustrate the kinds of nonlinear effects that might arise. We began by treating the simplest system: three nonlinearly interacting modes. Here, a key concept is parametric instability: when the amplitude of the unstable ``parent'' mode exceeds a well-defined threshold (Equation~\\eqref{ParINSTB2}), it will drain energy into the pair of ``daughters'' to which it couples. If the daughters dissipate energy very strongly, then the amplitude of the unstable parent generally cannot grow more than a factor of a few beyond its parametric instability threshold. However, even in such strongly damped cases, the dynamics can be quite complex, as is shown in Figures \\ref{ThreeModSol}(a)-\\ref{ThreeModSol}(c). Characteristically, the amplitude of the parent exhibits a sawtooth evolution and the system as a whole shows multiply periodic temporal behavior. As the dissipation drops, though, the three mode system evolves toward a single period, divergent solution, in which the amplitude of the unstable mode increases secularly. In a realistic system, additional modes would be excited at such low dissipation rates. Indeed, adding the triplet interactions with the second and third lowest parametric instability thresholds slowed and even halted the rate of growth of the instability (Figure \\ref{GrowthRates}).  In order to understand what happens in a system involving a multitude of modes, we began by studying a Hamiltonian system (i.e. no unstable or damped normal modes). For this system, we could track the evolution starting from given initial r-mode amplitudes and noise level. We found that in general the sea of inertial modes tends toward equipartition, but the behavior of the r-mode and even a few other lower order modes is richer. For example, starting with an r-mode amplitude above its second parametric instability threshold leads to equipartition in the sea of modes, but the r-mode itself appears to asymptote toward its first parametric instability amplitude.  We also studied the approach to equipartition as a function of the initial r-mode amplitude at fixed initial noise level. To measure the rate of evolution toward equipartition we introduced the time-dependent entropy (equation~\\eqref{entropyeq} ). In general, the approach to equipartition is faster for larger initial r-mode amplitudes. At smaller amplitudes, the approach to equipartition can become very slow (see Figure~\\ref{HamEntrop} ) and may not occur at all. At small amplitudes the r-mode, as well as a few other low order modes, continue to exhibit phase coherent oscillations up to the end of our simulation at $t\\sim 10^7/2\\Omega$. This maintenance of phase coherence signals a failure of the system to become ergodic, and also warns against assuming the random phase approximation too casually in trying to understand nonlinear truncation of the instability.  Finally, we turned to multimode non-Hamiltonian systems with growth due to the CFS instability, and damping due to (shear) viscous dissipation. There are a number of competing elements that determine how such a complicated system evolves. The two simplest are the growth or decay rates of the modes dictated by linear perturbation theory, but nonlinear interactions introduce several other important timescales. One was already evident in our study of the three mode system: in the absence of dissipation, the modal amplitudes evolve periodically as long as the total energy is not too large. A second was indicated by our study of the multimode Hamiltonian system: the rate of approach to equipartition represents a kind of multi-mode dissipation rate for the phase coherent low lying modes. Here we interpret dissipation broadly as the loss of energy from large scale coherent motion to small scale incoherent excitations.  For the full multimode system, we found that at strong damping only a few modes matter to the dynamics. The saw-tooth, chaotic behavior seen in the three mode system is also evident in strongly damped multimode systems. Many of the modes in the sea simply decay via viscous damping, never having achieved substantial excitation.  At very weak damping, though, numerous modes become involved in the evolution, which comes to resemble more closely what we observed in the approach to equipartition for a Hamiltonian system.  Modal spectra resemble that of the equipartition, Hamiltonian system, except with larger variances of some amplitudes about the mean. In the weak damping regime, the total energy of the system continued to rise throughout our integration time, but the r-mode amplitude appeared to level off. At its apparent asymptotic value, the rate of approach to equipartition -- as indicated by Fig. 6 for our Hamiltonian system -- would be just comparable to the linear growth rate of the CFS instability.  Thus, in this regime, the truncation of the r-mode amplitude itself is due to the ``collective dissipation'' in which energy disappears into a sea of coupling modes, a completely nonlinear effect. We expect that if we were to continue the calculations further in time, the energy growth would continue until the amplitudes of small scale, more highly damped modes became large enough for them to dissipate energy as quickly as the instability causes it to grow.  The behavior outlined above for the weakly damped system also depends on the relative rates of the CFS instability, approach to equipartion, and viscous dissipation. We examined this by increasing the stellar rotation rate, thus increasing the gravitational radiation rate. For this case, we did not observe any leveling of the r-mode amplitude, although we also did not run the simulation long enough (for practical reasons) for the amplitude to reach the value at which the growth rate of the CFS instability would be balanced by the rate of approach to equipartition. Although this does not indicate any inconsistency in the picture we have presented for how the instability is limited by nonlinear effects, it also cautions us that any calculation that enhances the instability growth rate could overshoot the amplitude of the r-mode at which it would level off in reality.  Finally, we note that the saturation amplitudes found here are substantially lower than those estimated in the cascade picture in Arras et al. \\cite{Saturation}, but are larger than but perhaps comparable to those relevant to the spin evolution of a neutron star estimated by Wagoner \\cite{Wagoner}. Our study of nonlinear effects on the development of the r-mode instability also suggests that the temporal evolution of the modal amplitude is not smooth on timescales $\\sim$ millions of stellar rotation periods. The erratic behavior seen in our simulations could, if realized in nature, complicate searches for gravitational radiation emission that depend on simple, phase coherent oscillations of the star.   Because of practical limitations our modal expansion is truncated at $n=30$.  Within this truncation the r-mode is found to go unstable first to  modes with $13\\leq n\\leq 15$. It is possible that some of the dynamics observed in this paper are dependent on our modal truncation.  The most dramatic effect would result if the minimum parametric instablity threshhold were lowered.  In this case the r-mode may go unstable to a very high order mode and dynamics such as equipartition rates may change considerably. It should be noted that this would imply that the maximum amplitude reached is smaller than or equal to those observed in this paper. Should the mode ever grow to amplitudes comparable to those observed here, all mechanisms explored in the current truncation will still limit its growth.  Our truncated expansion thus sets an upper bound for outcomes with larger $n$. The fact that no lower parametric threshholds were found in the region with $30\\leq n <15$ could, however, be an indication that we have gone far enough. Increasing the number of modes will have a secondary effect of increasing the ``specific heat'' or the system since more degrees of freedom have been introduced. This may be particularly relevant to the evolution in  Figure~\\ref{Evo5e8}, causing the noise level to rise more slowly.  In a large enough system, the flat spectrum observed in  Figure~\\ref{Evo5e8} should have a tail that is grounded below the noise level,  indicating a damping dominated region.  For more strongly damped regimes the modal dependence of the damping will play a significant role in determining the accessible phase space.  Sources of damping that are on average  weak, but have a few very strongly damped modes will have very little effect on r-mode dynamics; they will only remove a few degrees of  freedom from a very large phase space. On the other hand damping sources that have a fixed scaling, with large regions of modes damped roughly equally, will have a larger effect in  restricting the phase space.  Thus it is  not the maximum damping rate that is important to the dynamics, but rather the distribution of the most weakly damped modes.   For realistic neutron stars the eigenfunction structure will change somewhat and this will affect the couplings to some extent. Changes in the eigenfrequencies could have a more important effect, since subtle changes could shift various daughter modes into or out of close resonance. Such changes could strongly  influence the parametric instability threshholds.   The modal frequency structure, namely that for each $n$ and $m= 0 \\cdots n-1$ there exist  $n-m$ frequencies distributed between $-2 \\Omega$ and $2 \\Omega$, should persist for more realistic models.  It is this structure that ensures the near resonances among the daughter modes which in turn results in an equipartitioned solution. The rapid increase in equipartition rate with increasing amplitude is also likely to be the result  of the topology of the frequency spectrum and should persist in other models.   Thanks to Larry Kidder for helping with parallelization of the code and numerous other technical details.  This research is supported in part by NSF grants AST-0307273, PHY-9900672 and PHY-0312072 at Cornell University.    \\appendix"
    },
    "0410/astro-ph0410203_arXiv.txt": {
        "abstract": "{ Radio galaxy \\object{PKS~1138$-$262} is a massive galaxy at $z = 2.16$, located in a dense environment. We have found an overdensity of \\lya\\ emitting galaxies in this field, consistent with a proto-cluster structure associated with the radio galaxy.  Recently, we have discovered forty candidate \\ha\\ emitters by their excess near infrared narrow band flux.  Here, we present infrared spectroscopy of nine of the brightest candidate \\ha\\ emitters.  All these candidates show an emission line at the expected wavelength.  The identification of three of these lines with \\ha\\ is confirmed by accompanying [\\ion{N}{ii}] emission. The spectra of the other candidates are consistent with \\ha\\ emission at $z \\sim 2.15$, one being a QSO as indicated by the broadness of its emission line.  The velocity dispersion of the emitters (360 km s$^{-1}$) is significantly smaller than that of the narrow band filter used for their selection (1600 km s$^{-1}$).  We therefore conclude that the emitters are associated with the radio galaxy.  The star formation rates (SFRs) deduced from the \\ha\\ flux are in the range 6 -- 44 M$_\\odot$ yr$^{-1}$ and the SFR density observed is 5--10$\\times$ higher than in the HDF-N at $z = 2.23$.  The properties of the narrow emission lines indicate that the emitters are powered by star formation and contain very young ($< 100$ Myr) stellar populations with moderately high metallicities. ",
        "introduction": "High redshift clusters are prime subjects for the study of galaxy formation and cosmology. The powerful radio galaxy \\rg\\ at $z = 2.156$ appears to be the brightest galaxy in a high redshift cluster. We have discovered an overdensity of \\lya\\ emitters within 1.5 Mpc of \\rgs\\ \\citep[][Paper I]{kur00}. The redshifts of 14 emitters were spectroscopically confirmed to be in the range $2.14 < z < 2.18$ \\citep[][Paper II]{pen00a}. In addition, we have carried out near infrared imaging \\citep[][Paper III]{kur03a}. The number of $K$ band galaxies and extremely red objects in this field is higher than in blank fields.  We found 40 objects with excess narrow band flux, consistent with \\ha\\ emission at $z \\sim 2.16$. The surface density of \\ha\\ emitters increases towards the radio galaxy and their average $K$ magnitude is lower and therefore their inferred stellar mass higher than for the \\lya\\ emitters.  Here, we present infrared spectroscopy of nine candidate \\ha\\ emitters to confirm their redshift and determine the velocity dispersion of the sample.  We assume a flat Universe with $h_0=0.65$ and $\\Omega_{\\rm m}=0.3$. ",
        "conclusions": "Infrared spectroscopy has established the presence of nine line emitters within 0.6 Mpc of the HzRG \\rg. Three emitters show an additional line which confirms the identification with \\ha\\ at $z = 2.15$, while four more have spectra consistent with \\ha\\ at this redshift, one being a QSO as indicated by the broadness of its emission line. One emitter shows only a single strong line, which is possibly \\ha\\ and one emitter exhibits two lines which probably originate from two emission line regions within one galaxy at $z = 2.16$. Additional evidence for identification of all observed lines with \\ha\\ is the small velocity dispersion (360 km s$^{-1}$) as compared with the width of the selection filter. This dispersion is also smaller than the dispersion of the \\lya\\ emitters. The star formation rate density of the observed emitters is a factor ten higher than found at $z = 2.23$ in the HDF-N.  These results support the formerly advocated ideas that \\rg\\ is located in a proto-cluster at $z = 2.16$.  The properties of the narrow emission lines indicate that the emitters are powered by star formation and contain very young ($< 100$ Myr) stellar populations with moderately high metallicities. It seems that we observe these galaxies near the end of their first and major burst of star formation."
    },
    "0410/astro-ph0410729_arXiv.txt": {
        "abstract": "The equations for fully compressible rotating magnetoconvection are numerically solved in a Cartesian box assuming conditions roughly suitable for the geodynamo. The mean electromotive force describing the generation of mean magnetic flux by convective turbulence in the rotating fluid is directly calculated from the simulations, and the corresponding $\\alpha$-coefficients are derived. Due to the very weak density stratification the $\\alpha$-effect changes its sign in the middle of the box. It is positive at the top and negative at the bottom of the convection zone. For strong magnetic fields we also find a clear downward advection of the mean magnetic field.  Both of the simulated effects have been predicted by quasi-linear computations \\citep{1979PEPI...20..134S, 1992A&A...260..494K}. Finally, the  possible connection of the obtained profiles of the EMF with mean-field models of oscillating $\\alpha^2$-dynamos is discussed. ",
        "introduction": "Convective motions in the fluid outer core influenced by rotation and magnetic fields are able to maintain the Earth's magnetic field for long times interrupted by occasionally occurring reversals of the dominating dipole component. The equations describing the physical processes in the fluid outer part of the Earth's core are rather stiff \\citep{1995GAFD...79....1K,2000RvMP...72.1081R} showing a broad range of timescales on which the characteristic behavior is observable. The resulting and dominating advective timescale $\\tau_{\\mathrm{adv}}=d/u'$ leads to a very short coherence length of the physical quantities like velocity or magnetic field which, as a consequence, requires a very high spatial resolution in numerical simulations to include the effects of small-scale turbulence \\citep{2003GDS...31..181}. In case of a conducting and rotating fluid with large-scale density stratification a convectively driven turbulence can generate a mean electromotive force (EMF) parallel to the mean magnetic field --  a process known as $\\alpha$-effect \\citep{1980mfmd.book.....K}. But the Earth's fluid outer core, though a rapid rotator, is rather weakly stratified so that it remains unclear whether the amplitude of the $\\alpha$-effect is sufficient to maintain dynamo action.  Present-date global numerical simulations are unable to include small-scale motions because of computational restrictions. Nevertheless, they show quite satisfying results in the sense that they are able to reproduce several features of the observed geomagnetic  field \\citep{1996PEPI...98..207G, 1999GeoJI.138..393C, 1999JCoPh.153...51K}. In virtue of these results it seems to be justified to ignore small-scale fluctuations but, on the other hand, there still remain many open questions when comparing such simulations with their restricted parameter range with observational data (see e.g. \\citeauthor*{2000GGG...1..62}, \\citeyear{2000GGG...1..62}).  Background of our calculations is the idea to describe a dynamo by an induction equation for the mean magnetic field $\\left<\\vec{B}\\right>$ with a prescribed productive term due to the small-scale turbulence that ensures the existence of dynamo action.  In mean-field dynamo theory the magnetic field and the velocity are split up in a mean part, $\\left<\\vec{B}\\right>$ and $\\left<\\vec{u}\\right>$, and a fluctuating component, $\\vec{B}'$ and $\\vec{u}'$. The time behavior of the mean magnetic field $\\left<\\vec{B}\\right>=\\vec{B}-\\vec{B'}$ is described by \\begin{equation} \\partial_t\\left<\\vec{B}\\right>=\\nabla\\!\\times\\!\\left(\\left<\\vec{u}\\right>\\!\\times\\!\\left<\\vec{B}\\right>\\!+\\!\\vec{\\mathcal{E}}\\!-\\!{\\eta}\\ \\!\\nabla\\!\\times\\!\\left<\\vec{B}\\right>\\right)\\label{mean_field_ind} \\end{equation} with the mean electromotive force $\\vec{\\mathcal{E}}=\\left<\\vec{u}'\\times\\vec{B}'\\right>$ which is usually expressed by \\begin{equation} {\\mathcal{E}}_i = \\left<\\vec{u}'\\times\\vec{B}'\\right>_i= \\alpha_{ij}\\left<{B}_j\\right>+\\beta_{ijk}\\partial_k\\left<B_j\\right>. \\label{eq_alpha} \\end{equation} The tensor  $\\alpha_{ij}$  correlates the turbulent EMF due to small-scale motions with the large-scale magnetic field including the effects of anisotropy. The tensor $\\beta_{ijk}$ is related to the turbulent diffusivity $\\eta_T$ by $\\beta_{ijk}=\\eta_{\\mathrm{T}}  \\epsilon_{ijk}$. The easiest way to look for dynamo action is to solve equation (\\ref{mean_field_ind}) with zero mean flow $\\left<\\vec{u}\\right>$ and with an EMF taken from equation (\\ref{eq_alpha}) without turbulent diffusivity ($\\beta_{ijk}=0$).  Oscillating solutions for so-called $\\alpha^2$-dynamos have recently been presented by \\citet{2003PhRvE..67b7302S}. They determined a  radial profile of a spherically symmetric and isotropic $\\alpha$ under the conditions that the dynamo mode with the lowest eigenvalue is an oscillating solution and all stationary modes (that usually dominate $\\alpha^2$-models) are damped. This constraint leads to an $\\alpha$-effect with characteristic zeros in the  radial profile. Calculations with a uniform radial  $\\alpha$-coefficient have been performed by \\citet{2003A&A...406...15R} who, instead, included a latitudinal variation of $\\alpha$ and effects of anisotropy. They found oscillating $\\alpha^2$-dynamos only for exotic exceptions.  \\citet{2001GApFD..94..263H} analyzed the behavior of a geodynamo-like $\\alpha\\Om$-dynamo model which produces an axial dipole with reversals that are induced by a fluctuating $\\alpha$. The amplitude of the fundamental dipole mode behaves as a damped particle under the influence of a random force in a bistable potential. Here, in contrast to that, a configuration without any differential rotation and, hence, without $\\Om$-effect is considered.  The purpose of the present  paper is to calculate the turbulent EMF directly from numerical solutions of the full set of nonlinear MHD-equations for a convectively driven turbulent flow under the influence of rotation and subject to an imposed magnetic field. Our aim is to get a representation for the EMF based on simulations under conditions characteristic for the geodynamo. The results principally include effects of anisotropy of the $\\alpha$-effect, its radial dependence and its quenching properties. The derived $\\alpha$-coefficients will serve as some input data for future mean-field $\\alpha^2$-dynamo calculations. A simplified geometry of a Cartesian box representing a small part of a rotating spherical shell provides the ability to examine the small-scale behavior of the fluid motions and magnetic field and allows to consider effects of the turbulence that are neglected in global simulations. ",
        "conclusions": "The main result of our computations is that the $y$-component of the EMF describing the $\\alpha$-effect in the azimuthal direction is positive in the upper part of the unstable layer and negative in the lower part on the northern hemisphere. We obtained the $z$-profiles of the $\\alpha$-coefficient for a wide range of imposed magnetic field strength with the amplitude of the $\\alpha$-effect probably sufficient to ensure dynamo action in  mean-field calculations.  The $\\alpha$-effect is quenched under the influence of  strong external  magnetic fields but the suppression is not catastrophic, and significant quenching only sets in for magnetic fields above the equipartition value ${B_{\\mathrm{eq}}}$.  Some weak points in the above model calculations should be mentioned. First, we neglected the radial dependence of the gravitational force and the effects of spherical geometry, thus, the influence of curvature effects remains unknown. Compositional convection is ignored and the adopted parameter values are still far from the realistic values for the Earth. Somewhat unphysical boundary conditions have been used by assuming the conditions of a perfect conductor for both vertical boundaries so it cannot be ruled out that the obtained profiles are profoundly affected by the vertical boundary conditions. Test calculations with rigid boundary conditions for the velocity and/or slightly different boundary conditions for the magnetic field as  used e.g. by \\cite{2002A&A...394..735O} or \\cite{2002A&A...386..331Z} show, however, no major differences in the EMF profiles. It remains, therefore,  a much-promising ansatz for future mean-field calculations of the geodynamo taking into account the special radial profile of the $\\alpha$-effect as it has been obtained in the present paper. \\begin{ack} {This work was supported by the DFG SPP 1097 ``Erdmagnetische Variationen: Raum-Zeitliche Struktur, Prozesse und Wirkungen auf das System Erde''. The computations have been performed using the Hitachi SR8000 and the PC Cluster at the Astrophysikalisches Institut Potsdam (AIP).} \\end{ack}"
    },
    "0410/astro-ph0410173_arXiv.txt": {
        "abstract": "The most successful cosmological models to date envision structure formation as a hierarchical process in which gravity is constantly drawing lumps of matter together to form increasingly larger structures.  Clusters of galaxies currently sit atop this hierarchy as the largest objects that have had time to collapse under the influence of their own gravity.  Thus, their appearance on the cosmic scene is also relatively recent.  Two features of clusters make them uniquely useful tracers of cosmic evolution.  First, clusters are the biggest things whose masses we can reliably measure because they are the largest objects to have undergone gravitational relaxation and entered into virial equilibrium.  Mass measurements of nearby clusters can therefore be used to determine the amount of structure in the universe on scales of $10^{14}$-$10^{15} \\, M_\\odot$, and comparisons of the present-day cluster mass distribution with the mass distribution at earlier times can be used to measure the rate of structure formation, placing important constraints on cosmological models.  Second, clusters are essentially ``closed boxes'' that retain all their gaseous matter, despite the enormous energy input associated with supernovae and active galactic nuclei, because the gravitational potential wells of clusters are so deep. The baryonic component of clusters therefore contains a wealth of information about the processes associated with galaxy formation, including the efficiency with which baryons are converted into stars and the effects of the resulting feedback processes on galaxy formation.  This article reviews our theoretical understanding of both the dark-matter component and the baryonic component of clusters, providing a context for interpreting the flood of new cluster observations that are now arriving from the latest generation of X-ray observatories, large optical surveys, and measurements of cluster-induced distortions in the spectrum of the cosmic microwave background. ",
        "introduction": "\\label{sec:intro}  Cosmology has recently reached an important milestone.  A wide variety of cosmological observations now support a single model for the overall architecture of the observable universe and the development of galaxies and other structures within it.  According to this so-called concordance model, the geometry of the observable universe is indistinguishable from a flat geometry, implying that its total energy density is very close to the critical density needed to close the universe.  The two dominant components of the universe appear to be a non-baryonic form of dark matter, whose gravity is responsible for structure formation, and a mysterious form of dark energy, whose pressure is currently causing the expansion of the universe to accelerate.  The mean density of baryonic matter is about 15\\% of the total amount of matter, and we can observe the baryonic matter  only because the gravitational attraction of non-baryonic dark matter has drawn the baryonic gas into deep potential wells, where a small fraction of it condenses into stars and galaxies.  This model explains many different features of the observable universe, but it is not entirely satisfying because the nature of the dark matter and the provenance of the dark energy remain unknown.  The implications of dark energy for fundamental physics are particularly serious, so we need to be sure that it is absolutely necessary to explain the astronomical observations.  In addition, many aspects of galaxy formation remain poorly understood.  Dark-matter models successfully account for the spatial distribution of mass in the universe, as traced by the galaxies, but they do not explain all the properties of the galaxies themselves.  Dark matter initiates the process of galaxy formation, but once stars begin to form, supernova explosions and disturbances wrought by supermassive black holes can inhibit further star formation by pumping thermal energy into the universe's baryonic gas.  Clusters of galaxies are a particularly rich source of information about the underlying cosmological model, making possible a number of critical tests.  According to the concordance model, clusters are the largest and most recent gravitationally-relaxed objects to form because structure grows hierarchically.  The universe begins in a state of rapid expansion whose current manifestation is Hubble's Law relating a galaxy's recessional velocity $v_r$ to its distance $d$ through Hubble's constant $H_0$: $v_r = H_0 d$.   Generalizing this feature of the local universe to all of observable space links an object's cosmological redshift $z = (\\lambda_{\\rm observed}/\\lambda_{\\rm rest}) - 1$ with a unique time $t(z)$ since the Big Bang, enabling us to probe the evolution of the universe with observations of distant objects.\\footnote{In this definition, $\\lambda_{\\rm rest}$ is the wavelength of a photon emitted by a distant object and $\\lambda_{\\rm observed}$ is the wavelength it is observed to have when it reaches Earth.}  Gravity drives structure formation in this expanding realm because the matter density is nearly equal to the critical density during much of cosmic history.  Regions whose density slightly exceeds the mean density are therefore gravitationally bound and eventually decouple from the expansion, collapse upon themselves, and enter a state of virial equilibrium in which the mean speeds of the component particles are approximately half the escape velocity. Because density perturbations in the concordance model have greater amplitudes on smaller length scales, small sub-galactic objects are the first to decouple, collapse, and virialize.  These small objects then collect into galaxies, and galaxies later collect into clusters of galaxies, whose masses now top out at roughly $10^{15}$ times that of the Sun's ($10^{15} \\, M_\\odot$).  Thus, the growth and development of clusters directly traces the process of structure formation in the universe.  Section~\\ref{sec:observ} outlines the observable properties of galaxy clusters that enable us to measure their masses.  Observables in the optical band include the overall luminosity of a cluster's galaxies, which scales with the overall mass, the velocity dispersion of a cluster's member galaxies, which responds to the depth of the cluster's potential well, and gravitational lensing of background galaxies by the cluster's potential.  Observables in the X-ray band include the overall X-ray luminosity of a cluster, coming from the hot gas trapped in the cluster's gravitational potential, the temperature inferred from the X-ray spectrum of that gas, and the abundances of various elements inferred from the emission lines in that spectrum. This hot gas also leaves an imprint on the microwave sky because its electrons Compton scatter the photons of the cosmic microwave background radiation.  Microwave observations are therefore an alternative source of information about the hot gas and its temperature.  Once we have measured the masses of a sample of clusters, we can use that sample to study cosmology.  Section~\\ref{sec:dcomp} explains how the characteristics of the cluster population relate to cosmological models.  It begins by summarizing the elements of the concordance model and provides a number of useful analytical approximations to the results of numerical simulations of cluster formation based on this model.  Then it covers the dicey middle ground linking those simulations with observations, currently the main source of uncertainty in deriving cosmological parameters from cluster observations. The section concludes with a look at the evolution observed in the cluster population and the constraints that cluster evolution places on cosmological models.  Section~\\ref{sec:baryons} takes up the subject of the baryonic component of clusters, with two purposes in mind.  First, in order to improve the precision of cosmological measurements with clusters, we need to know how the process of galaxy formation affects the relations used to derive cluster masses from observations of a cluster's hot gas and galaxies.  Current numerical simulations accurately reproduce the behavior of the dark component, whose interactions are purely gravitational, but fail to reproduce with similar accuracy the observed behavior of the baryonic component, whose interactions are also hydrodynamical and thermodynamical.  These discrepancies between simulations and observations indicate that galaxy formation alters the state of a cluster's hot gas in a way that preserves information about the poorly understood feedback processes that regulated galaxy formation long before the cluster reached its present state.  Our second purpose is therefore to try to decipher what the state of the hot gas is saying about the process of galaxy formation, so as to gain insight into those feedback processes.  Section~\\ref{sec:concl} concludes the review with some brief remarks about ongoing and future cluster surveys.  Despite this article's length, it falls somewhat short of being a comprehensive review of cluster physics, which would require more pages than this journal is inclined to provide. Instead, I have tried to assemble a readable introduction to cluster evolution for non-experts, concentrating on the middle ground connecting theory to observations and distilling the key theoretical results into a set of simple analytical tools useful to observers.  For more on the subject of clusters and their evolution, readers should consult \\citet{Sarazin88}, \\citet{bmv02}, \\citet{rbn02}, and \\citet{mdo04}. ",
        "conclusions": "\\label{sec:concl}  The next decade of research into cluster evolution promises to be very exciting.  Large optical surveys like the Sloan Digital Sky Survey are greatly increasing the number of well-studied clusters of galaxies in the low-redshift universe.  Deep surveys looking for the Sunyaev-Zeldovich effect will be finding thousands of clusters to distances well beyond a redshift of $z =1$.  The {\\em Chandra} and {\\em XMM-Newton} X-ray observatories are providing our most detailed look yet at the intracluster medium, its thermodynamical state, and some of the feedback processes that regulate condensation of intergalactic gas into galaxies and stars. Also, dedicated X-ray satellite missions to survey a large fraction of the sky for distant clusters are currently being planned.  Making the most of these opportunities will require cooperation between observers in those different wavebands, and theoretical modeling that closely links those cluster observables to cosmological parameters. Optical and infrared followup of S-Z surveys will be critical in order to determine the redshifts of the cluster candidates.  X-ray followup of a subset of the S-Z clusters will also be necessary to establish how the thermodynamics of galaxy formation affects evolution of the mass-observable relations in the microwave band.  Concentrated efforts to observe a calibration set of clusters in all of these wavebands will be very valuable in helping to establish how well the various observables trace mass and the scatter in each of these observables at a given mass.  If the $\\Lambda$CDM concordance model is indeed a good description of the overall architecture of the universe and its initial perturbation spectrum, then the parameters describing the cosmological context in which galaxy formation happens ought to be quite precisely established within this decade.  Studies of cluster evolution will be just one part of this overall effort, which also includes distance determinations to high-redshift supernovae, increasingly sensitive observations of the cosmic microwave background, and the mapping of large-scale structure.   However, consistency between the dynamics of cluster evolution and the geometry of the universe, as measured with supernova and microwave observations, will stand as a particularly critical test of the overall model.  With success, most of the remaining secrets about galaxy formation---other than what dark matter and dark energy actually are---will concern baryons and their complex cooling and feedback processes.  Our understanding of what baryons do is rapidly progressing, thanks in large part to large-scale cosmological simulations on massively parallel computers.  Clusters and their evolution place unique constraints on those models because clusters are the only places in the universe where the majority of the baryons emit detectable radiation, revealing their thermodynamic state and elemental abundances.  Galaxy formation has clearly left its mark in the intracluster medium, but we are just beginning to decipher what it has written there in the gases between the galaxies. Perhaps in ten more years there will be as much optimism about understanding the baryonic side of galaxy formation as there is now about the understanding the darker side."
    },
    "0410/astro-ph0410665.txt": {
        "abstract": "Emission-lines of magnesium at $12\\,\\micron$ have been observed in the spectrum of Procyon. We reproduce the observed, disk-averaged line flux from Procyon (as well as the observed intensity profiles from the Sun) by calculating the line formation, relaxing the assumption of Local Thermodynamic Equilibrium. We find that the lines in Procyon are formed in the photosphere in the same manner as the solar lines. We discuss our modeling of these Rydberg lines and evaluate, among other things, the importance of the ionizing flux and updated model-atom parameters. The lines are of large diagnostic value for measurements of stellar magnetic-fields through their Zeeman splitting. %%%%  changed by mjr %We have, however, not been able to detected any splitting of the \\mgi\\ %lines in Procyon. We have not, however, detected splitting of the \\mgi\\ lines in Procyon. %%%% %We set an upper limit on the disk averaged magnetic %field of 100 Gauss and Using simple arguments, we believe we would have detected %longitudinal(???) a magnetic field, had it been of a strength larger than approximately 800 Gauss covering more than a quarter of the surface. We discuss the prospects for future use of the Zeeman-sensitive, mid-infrared, \\mgi\\ emission lines as a diagnostic tool for stellar magnetic fields. ",
        "introduction": "Although magnetic fields affect, and sometimes dominate, the appearance of cool stars, direct detection and determination of magnetic field strengths are difficult. Methods involving polarized light and magnetograms are problematic for cool stars with assumed solar-like magnetic geometries \\citep{saar}. On the other hand, studies of unpolarized Zeeman-splitting in the near-infrared have given a large scientific return \\citep[see, for example, ][]{muglach92,solanki92,solanki:154,valenti95,johnskrull99}.  \\mgi\\ emission lines at $12~\\,\\micron$ in sunspot penumbrae and plages clearly show Zeeman splitting and so trace the local magnetic field \\citep[see, for example, ][]{brault83,bruls:1,bruls:2}. Zeeman splitting is independent of the field direction and will show up even in a disk-averaged spectrum from a star with a disorganized field structure, which is not the case for polarized light. The wavelength dependence of Zeeman broadening compared to other types of broadening results in the $12\\,\\micron$ lines being more Zeeman-sensitive than near-infrared (NIR) lines. Furthermore, the lines probe other atmospheric layers than the NIR lines. Thus, these lines potentially offer an excellent method for measuring stellar disk-averaged magnetic fields.  %Magnetic fields and activity affect properties of stellar atmospheres of cool stars and %are, therefore, of large interest to determine. %% \\citep{saar}. %%Accurate magnetic-field determinations are, therefore, of large interest. %% stellar work. %However, direct assessments of stellar field-strengths %% in stellar atmospheres including the Sun's %are often rather involved, if possible at all. %Methods measuring and analyzing polarized light and magnetograms, have been shown to be unfavorable, at least for %cool stars with assumed solar-like magnetic geometries \\citep{saar}. %On the other hand, studies of un-polarized light of %%magnetic-field measurements of %Zeeman-splitting of normally degenerate photospheric-lines in the near-infrared have %given a large scientific return \\citep{solanki:154}. %Also, the solar \\mgi\\ emission lines at $12~\\,\\micron$ in sunspot penumbrae and plages %clearly show Zeeman splitting and so trace the local magnetic field \\citep[see, for example, ][]{brault83,bruls:1,bruls:2}. %Zeeman splitting is independent of the field direction and %will show up even in a disk-averaged spectrum from a star having a disorganized field structure, which is not %the case for polarized light.  Solar emission features at $12\\,\\micron$ were first mentioned in 1981 (Murcray et al. 1981), and attracted much attention during the 1980s. In 1983, Chang \\& Noyes \\nocite{chang83} identified the strongest two emission lines as coming from transitions between Rydberg states in \\mgi. \\citet{lemke87} showed that the emission could be formed in the photosphere through stimulated emission and different departures from Local Thermodynamic Equilibrium (LTE) for the lower and upper levels. The first non-LTE modelling giving the lines in emission were by Carlsson et al. (1990, 1992b) and Chang et al. (1991). The line formation process was first explained successfully and in detail by \\citet{mc_mg}. They concluded that the lines were photospheric in origin and accounted very well for their emission strengths and their complicated intensity profiles as a function of the viewing angle on the solar disk.   \\citet{uiten}, observing just one of the two lines, detected emission from the K giant $\\alpha$ Boo and saw absorption from the five M giants and supergiants in their sample.  Their model calculations were unable to fit the data. \\citet{ryde&richter} detected and analysed both Mg I emission lines from Procyon ($\\alpha$~CMi) observed at a resolving power of $R = 86\\,000$. As high-resolution, mid-infrared spectrographs become available on $8-10$~m telescopes, these lines will be observed in many more stars.  A first step towards the sophisticated use of these lines for measuring stellar magnetic fields is the understanding of their formation and response to a broad range of stellar parameters.  In this paper, we discuss our analysis of the magnesium emission lines detected from Procyon. %which was announced by us in \\cite{ryde&richter}. We %investigate our calculation by discuss here the atomic model of \\mgi\\ and the assumptions made in our non-LTE calculation, such as the collisional rates, the atmospheric radiation field used, and the validity of the assumption of atmospheric homogeneity. We conclude with a discussion on using these lines for determining disk-averaged, magnetic-field strengths in stars and an estimate of the maximum field strength and corresponding filling factor allowed by our observations.   %The origin and formation process of the solar \\mgi\\ %lines was first explained successfully and %in detail by \\citet{mc_mg}, by treating the lines in %non Local Thermodynamic Equilibrium (non-LTE). %They concluded the lines were photospheric in origin and %accounted excellently for their emission strengths and their complicated intensity profiles as a %function of viewing angle on the solar disk. %In this paper, we will discuss our analysis of the magnesium emission lines detected from Procyon %($\\alpha$ CMi). %%which was announced by us in \\cite{ryde&richter}. %We will %investigate our calculation by %discussing the atomic model of \\mgi\\ %and the assumptions made in our non-LTE calculation, %such as the collisional rates, atmospheric radiation field %used and the validity of the assumption of atmospheric homogeneity. ",
        "conclusions": "The splitting of the solar $12\\,\\micron$ lines in magnetically active surface regions shows normal Zeeman triplet splitting with one $\\pi$ and a pair of $\\sigma$ components, although all the components actually have tightly grouped substructure. The $\\pi$ and $\\sigma$ components are separated by a width given by an effective Land\\'e factor of approximately 1, which was experimentally found by \\cite{lemoine}. The average separation between the two $\\sigma$ components is given by \\begin{equation} \\label{ekvation} \\Delta\\lambda_B = 2\\times 4.67\\times10^{-9}g\\lambda^2_0 B \\,\\,\\,(\\mu\\mathrm m ), \\end{equation} where $g$ is the Land\\'e factor, $\\lambda_0$ is the unperturbed wavelength of the line in $\\micron$, and $B$ the magnetic field strength in Gauss (G). %{\\it ***  I changed the formula so it would be in microns, since we're %working in the mid-IR.  It would be worth double-checking.  I also changed %$g_\\mathrm{eff}$ to simply $g$ since that is how I have usually seen %it written.  $g_\\mathrm{L}$ would also seem reasonable (but make sure %it gets changed below the equation to agree).  **** } The strength of the solar magnetic fields puts the Zeeman effect in the Paschen-Back regime for these lines \\citep{bruls:2}.   %The line splitting of the solar $12\\,\\micron$ lines in magnetic active regions of the surface, %shows a normal Zeeman triplet splitting, with seemingly  one $\\pi$ and a pair of $\\sigma$ components, but %these components actually have tightly grouped substructure. %The $\\pi$ and $\\sigma$ components are separated by a width given by an effective Land\\'e factor of approximately 1, %which was experimentally found %by \\cite{lemoine}. %The average separation between the two $\\sigma$ components is given by %\\begin{equation} %\\label{ekvation} %     \\Delta\\lambda_B = 2\\times 4.67\\times10^{-13}g_\\mathrm{eff}\\lambda^2_0 B, %\\end{equation} %where where $g$ is the Land\\'e factor, $\\lambda_0$ is the unperturbed wavelength of the line in \\AA , and $B$ the magnetic-filed %strength in Gauss. %The relevant strengths of the magnetic-fields put %the Zeeman effect in the Paschen-Back regime for these lines \\citep{bruls:2}. %  There are advantages to Zeeman magnetic-field studies using infrared lines compared to optical ones. In addition to less line-blending due to the lower line density \\citep[see, for example, ][]{ryde_munchen}, infrared spectral lines have larger Zeeman sensitivity. Optical lines most often do not show large enough Zeeman separation in order to be disentangled from other broadening mechanisms of the line. The ratio of the separation of the Zeeman splitting ($\\Delta\\lambda_\\mathrm{B}\\sim g\\lambda^2B$) and the non-magnetic Doppler width ($\\Delta\\lambda_\\mathrm{D}\\sim\\lambda$), is given by $\\Delta\\lambda_\\mathrm{B}/\\Delta\\lambda_\\mathrm{D}\\sim g\\lambda B$. Thus, the ratio grows with wavelength and, consequently, it is easier to detect Zeeman split lines in the infrared. With larger Zeeman separation between the components, the analysis is simpler. %{\\it *** There are other papers that have made this point before.  Perhaps %we should reference them?  *** }  %There is a general benefit of analysing magnetic-field strengths manifested in IR lines compared to optical %ones. %In addition to less line-blending due to the lower line density \\citep[see, for example, ][]{ryde_munchen}, %infrared spectral lines are also more sensitive to magnetic fields through the Zeeman effect. %Optical lines most often do not show large enough Zeeman separation in order to be disentangled from other %broadening mechanisms of the line. %The ratio of the separation of the Zeeman splitting ($\\Delta\\lambda_\\mathrm{B}\\sim g\\lambda^2B$) and the non-magnetic %Doppler width %($\\Delta\\lambda_\\mathrm{D}\\sim\\lambda$), is given by %$\\Delta\\lambda_\\mathrm{B}/\\Delta\\lambda_\\mathrm{D}\\sim g\\lambda B$. %Thus, the ratio grows with wavelength and, %consequently, it is easier to detect Zeeman %split lines in the infrared. The larger the %Zeeman separation between the components the better for the analysis.   Although the largest Land\\'e factors for optical and near-infrared transitions are roughly 3 while the \\mgi\\ lines at $12\\,\\micron$ have an effective Land\\'e factor of 1, moving to $12\\,\\micron$ from the optical and near-infrared, respectively, increases the sensitivity by at least a factor of 8 and 3, respectively. %{\\it *** These factors take into account Land\\'e factor as well as %wavelength?  *** } According to \\citet{solanki:154}, it should be possible to measure solar magnetic field strengths as low as $150-200$ Gauss from line splitting, and fields down to 100 Gauss are  measurable through (very) careful profile fits. %{\\it  ***  Can we refer to a paper in a refereed journal that is more %generally accessible?  *** } (A large magnetic field would split the line and in the case of low field-strength it would act as an additional broadening.) \\cite{bruls:2} discuss measurements of the Stokes parameters of the solar $12\\,\\micron$ lines affected by solar-magnetic fields, and indicate that field-strengths of approximately 200 G are measurable through \\emph{Stokes V}-line splitting. They also show that line-splitting in the \\emph{Stokes I} light is reliable only for strengths greater than 500 G.  It should be noted that, especially for the broadened lines, the sensitivity for detecting magnetic fields for a disk-averaged stellar, \\mgi\\ spectrum will also depend on the filling factor ($f$) of the field over a stellar surface \\citep[see for example, ][]{valenti95}. In the simplest case, a stellar magnetic field can be modeled by a two-component model, with one magnetic component with a constant field strength, $B$, over a fraction, $f$, of the surface and one non-magnetic component, the latter covering ($1-f$) of the stellar surface \\citep{solanki:invited}. Furthermore, in a field-free approximation \\citep[see, for example, ][]{bruls:2} one can assume that a magnetic field does not affect the departure coefficients significantly. The non-LTE calculations can, therefore, be made in a non-magnetic environment (Piskunov, private communications).  %In spite of the fact that the largest Land\\'e factors in the optical and near-infrared are of the order of 3 whereas the %lines at $12\\,\\micron$ have an effective Land\\'e factor of 1, moving to $12\\,\\micron$ from the %optical and near-infrared, respectively, increases the sensitivity by at least a factor of 8 and 3, respectively. %According to \\citet{solanki:154}, who discusses the case of the Sun, it should, therefore, be possible to measure field %strengths as low as of 150-200 Gauss from line splitting, and fields down to %100 Gauss are  measurable through (very) careful profile fits. (The magnetic field would split the line and %in the case of low field strength it would act as an additional broadening.) \\cite{bruls:2} discuss measurements of the Stokes %parameters of the solar $12\\,\\micron$ lines affected by solar-magnetic fields, %and indicate that field-strengths of approximately 200 G are measurable through \\emph{Stokes V}-line splitting. %They also show that line-splitting in the \\emph{Stokes I} light is reliable only for strengths greater than 500 G. %It should be noted that the sensitivity for a disk-averaged %stellar, \\mgi\\ spectrum also will depend on %the filling factor of the field over a stellar surface.   %As mentioned above, we need to convolve our modeled lines with a Gaussian %having a FWHM of $9.5\\,$km\\,s$^{-1}$. %This is somewhat larger than the FWHM of roughly %$8.5\\,$km\\,s$^{-1}$ found for optical spectra. %If we attribute the extra broadening entirely to Zeeman broadening and thus assume that the %macroturbulence is the same measured at $12\\,\\micron$ as in the optical wavelength %region, Eq. \\ref{ekvation} formally %yields an averaged fields-strength of approximately 100 Gauss. However, in order to measure such low, %averaged field-strengths, a much more careful line fitting is necessary.  %we can set an upper limit on the disk-averaged field %using Eq. \\ref{ekvation} of $B<100$~Gauss. %Unfortunately, it is unknown if the macroturbulence appropriate to the %optical also applies at %$12\\,\\micron$.  {\\it  *** Are there any near-IR measurements of Procyon? %If they give 8.5 km/s for macroturbulence, then we are on more solid ground. %If any molecular lines were present in Procyon, it would be OH absorption %at 814.324, 814.727, 815.402, and 815.953 \\invcm.  If these lines were %detected, we'd be able to determine a width from them. %Can we put in a figure (and accompanying text) that shows %what the model would predict for a 500 Gauss field covering 25\\%\\ of the %star?  I pulled those numbers out of a hat, but the idea would be to see %how low we could go on both parameters and get a firmer limit.  I see that %Bruls and Solanki (1995) suggest modeling the \\mgi\\ lines requires %fairly sophisticated treatment of the flux tubes.  How wrong is the %simple idea of taking a fraction of the line flux (depends on filling %factor) and moving it around (depends on field strength).  We also %need text regarding the likely values of Procyon's magnetic field.  Since %it is a hotter (less convection) and fatter (evolving off main sequence) %star, I would naively expect it to have weaker magnetic fields.  Were %any of the lines Carlos studied with such detail magnetically sensitive? %Did they report any limit on the magnetic field?   *** }  We have not resolved any magnetic components on Procyon and see only single peaked features that are well fitted by Gaussian distributions.  Given the absence of any resolved $\\sigma$ components, we can estimate the field strength and filling factor that would have produced detectable, separate, $\\sigma$ components. We consider the line at 811.578~\\invcm\\ because the signal-to-noise ratio for this line, 130:1, is higher than for the 818.058~\\invcm\\ line. Our estimate considers only the magnetic field strength and filling factor at the photospheric height that dominates the line formation. Pressure equilibrium and magnetic flux conservation suggest that magnetic flux tubes will have a larger filling factor  and weaker field strength higher in the photosphere where the gas pressure is lower.  Given our observations, we believe we would have detected resolved $\\sigma$ components if the magnetic field strength had been $B>830\\,$G and the filling factor $f>26\\%$.  In this simplistic estimate, we assume the observed line, which has a peak height of 0.135 above the normalized continuum, comes from the non-magnetic region with filling factor $(1-f)$.  Given the signal-to-noise ratio, a 3$\\sigma$ detection would require a peak value of 0.024 above the normalized continuum. If the efficiency for producing the two $\\sigma$ components in the magnetic region is the same as for the central peak from the non-magnetic region, we can calculate that the magnetic filling factor must be approximately  26\\%. To determine the maximum field strength allowed by our observations, we consider a separation between the two $\\sigma$ components of 28.5~km\\,s$^{-1}$ (3 times the FWHM of the detected line) to be sufficiently resolved that we would resolve the separate peaks.  Using Eq. \\ref{ekvation} given above, such a  splitting would require a magnetic field strength of 830~G. %Lower field strengths would be possible %to measure through a careful and often involved Zeeman-broadening analysis. For comparison, it can be noted that similar and also stronger fields than this have been detected for other stars. For example, \\cite{johnskrull99} measure a surface-averaged mean-field from IR observations of the T Tauri star BP Tau of $\\Sigma Bf=2.6\\pm0.3$\\,kG, and \\citet{valenti95} find that 9\\%\\ of the deep photosphere of the active star $\\epsilon$ Eri is filled with a 1.4 kG field.  %LLLLLLLLLLLLLLLLLLLL %As mentioned above, we need to convolve our modeled lines with a Gaussian %distribution having a FWHM of $9.5\\,$km\\,s$^{-1}$, which is somewhat larger than the FWHM of roughly %$8.5\\,$km\\,s$^{-1}$ found for optical spectra. %Assuming, that the macroturbulence is the same measured at $12\\,\\micron$ as in the optical wavelength %region, this difference amounts to an extra broadening of the order of $0.01\\,$\\invcm. %Through Eq. \\ref{ekvation} a field of approximately 100 Gauss would separate the $\\sigma$ components %by that amount. However, in order to actually measure such an averaged %field-strength, a much more careful %%%%%  spelling: \"modelling\" -> modeling   - mjr %%line fitting and modelling is necessary. %line fitting and modeling is necessary. %%%%  A further diagnostic value of the $12\\,\\micron$ lines is their different height of formation compared with lines at, for example, $1.5\\,\\micron$. The $1.5\\,\\micron$ lines are commonly used in solar, magnetic-field observations and sample the opacity minimum. For the Sun, the lateral expansion of magnetic flux tubes with height dilutes the field strength \\citep{deming}. Thus fields measured at $12\\,\\micron$, sampling higher regions, may be lower than those measured by lines lying at $1.5\\,\\micron$. Because of the different depth dependence, comparisons of Zeeman splitting from lines at different frequencies can start to trace the vertical structure of magnetic flux tubes \\citep{bruls:2}.  While the $12\\,\\micron$ lines provide high sensitivity to magnetic fields, observing these lines is still difficult. As $8-10$\\,m class telescopes become equipped with high-resolution, mid-infrared spectrographs, many more stars will prove suitable targets. We are confident that high-resolution observations of the $12\\,\\micron$ lines will, certainly, be an important tool for measuring \\emph{stellar} magnetic fields in the future. We have shown that we understand the formation of the $12\\,\\micron$ lines not only for the Sun but also for Procyon and that the non-LTE calculation, based on the one for the Sun, is still up-to-date.   %Thus, the $12 \\,\\micron$ lines %potentially offer an excellent method for measuring stellar disk-averaged %magnetic fields. %% in stars with large averaged fields strengths. %A further diagnostic value of the $12\\,\\micron$ %lines is their different height of formation compared with lines at, for example, $1.5\\,\\micron$. %The $1.5\\,\\micron$ lines are %commonly used in solar, magnetic-field observations and sample the opacity minimum. %However, in spite of the fact that the sensitivity of the $12\\,\\micron$ lines to %magnetic fields is high,  the spectrograph/telescope combinations of today still need increased %sensitivity to make the $12\\,\\micron$ line a readily available tool for measuring \\emph{stellar} magnetic fields, %since only the brightest objects can be observed so far.  %It should also be noted that stars of, for example, spectral class M are brightest at $1-2\\,\\micron$. These stars, and even more %so for warmer stars, will %be fainter in the mid-infrared compared to the near-infrared, which is a drawback. %Furthermore, at least for the Sun, the magnetic-field strength decreases with %height since, due to the decrease %in gas-pressure with height, solar flux tubes expand laterally \\citep{deming}. %Thus fields measured at %$12\\,\\micron$, sampling higher regions, ought to be lower than those measured %by lines lying at $1.5\\,\\micron$, %should such an effect be at play. % %Nevertheless, high-resolution observations of $12\\,\\micron$ lines will, certainly, be %an important tool for measuring \\emph{stellar} magnetic fields in the future. %We believe that we now understand the process of formation of the $12\\,\\micron$ lines %not only for the Sun but also for the first stellar object other than the Sun, namely %$\\alpha$ CMi or Procyon.  %KE: ISM absorption line? of Mg?samma p\u00e5 h\u00f6ger sida om b\u00e5da linjerna.  %"
    },
    "0410/astro-ph0410459_arXiv.txt": {
        "abstract": "We present the CCD photometry, light curve and time series analysis of the classical nova V2275 Cyg (N Cyg 2001 No.2). The source was observed for 14 nights in total using an $R$ filter in 2002 and 2003 with the 1.5 m Russian-Turkish joint telescope (RTT150) at the TUBITAK (The Scientific and Technical Research Council of Turkey) National Observatory in Antalya Turkey, as part of a large program on the CCD photometry of Cataclysmic Variables (CVs). We report the detection of two distinct periodicities in the light curve of the nova : a) P$_1$=0.31449(15) d -- 7.6 h, b) P$_2$=0.017079(17) d -- 24.6 min. The first period is evident in both 2002 and 2003 whereas the second period is only detected in the 2003 data set. We interpret the first period as the orbital period of the system and attribute the orbital variations to aspect changes of the secondary irradiated by the hot WD. We suggest that the nova was a Super Soft X-ray source in 2002 and, perhaps, in 2003. The second period could be a QPO originating from the oscillation of the ionization front (due to a hot WD) below the inner Lagrange point as predicted by King (1989) or a beat frequency in the system as a result of the magnetic nature of the WD if steady accretion has already been re-established. ",
        "introduction": "Classical novae are a subset of cataclysmic variables which are interacting binary systems hosting a main-sequence secondary (sometimes a slightly evolved star) and a collapsed primary component, a white dwarf (Warner 1995). An outburst on the surface of the white dwarf as a result of a thermonuclear runaway in the accreted material causes the ejection of 10$^{-3}$ to 10$^{-7}$ M$_{\\odot}$ of material at velocities up to several thousand kilometers per second (Shara 1989; Warner 1995). V2275 Cyg (Nova Cyg 2001 No.2) was discovered  at a magnitude 7.0-8.8 on 2001 August 18 simultaneously by Nakamura et al. (2001) and Nakano et al. (2001). Early optical spectroscopy showed hydrogen Balmer lines with P Cygni profiles and H$\\alpha$ lines indicating expansion velocities of 1700 km s$^{-1}$ (Ayani 2001). At later stages, high energy coronal lines were found to dominate the spectrum (e.g., [Si X], [Si IX] and [Al IX]). The nova was found to belong to the \"He/N\" subclass of novae defined by Williams (1992), because of the broad lines of H, He and N in its spectrum (Kiss  et al. 2002). In addition, Kiss et al. (2002) measured t$_2$=2.9$\\pm$0.5 d, t$_3$=7$\\pm$1 d and M$_V$= -9$^{m}$.7$\\pm$0.$^{m}$7 which were used to derive a distance of 3-8 kpc for the nova. A USNO star of $R$=18$^{m}$.8 and $B$=19$^{m}$.6 was suggested as a possible progenitor star (Schmeer et al. 2001).  This paper is on the CCD photometry of V2275 Cyg covering years 2002 and 2003 obtained with the 1.5 m Russian-Turkish joint telescope (RTT150) at the TUBITAK National Observatory (TUG) in Antalya, Turkey. We present the discovery of a distinct period reported by Balman et al. (2003) (which we mark as P$_1$) and the detection of a highly coherent QPO (Quasi-periodic oscillations) on 6 different nights in 2003 (using $V$, $R$, $I$ filters). Fast variations from V2275 Cyg similar to  this second period (which we mark as P$_2$) were also reported by Garnavich et al. (2004) derived from a data set of two nights obtained with a $V$ filter.  \\begin{figure*} \\begin{center} \\epsfig{file=balman_f1.ps,width=9.0cm,height=12.7cm,angle=-90} \\end{center} \\caption{The top panel presents the normalized differential light curve of V2275 Cyg observed on 2002 October 01-05. The middle panel displays the light curve obtained on 2003 August 25-27 and the bottom panel shows the light curve obtained on 2003 November 01-02. All data are taken with the TUG 1.5 m telescope using the standard R band filters (Johnson and Cousins). The epoch of the observations are noted on the x axis. The average magnitude errors are $\\pm$0.0083, $\\pm$0.0065 and $\\pm$0.0095 for the top, middle and bottom panels, respectively.} \\end{figure*}  \\begin{table} \\center{ \\caption{The Time Table of the Observations -- (All observations are obtained with standard $V$, $R$, $I$ filters)} \\begin{tabular}{llllll} \\hline UT & Time of Start  & Run &  Number of  & Filter\\\\ Date & (HJD-2452000) & Time (hr) & Frames  & \\\\ \\hline 020610 & 436.484518  & 1.8 & 30 & R \\\\ 020612 & 438.486060  & 1.9 & 29 & R \\\\ 021001 & 549.246533  & 7.1 & 88 & R \\\\ 021002 & 550.248770  & 5.3 & 68 & R \\\\ 021003 & 551.256457  & 5.9 & 99 & R \\\\ 021005 & 553.293507  & 5.4 & 80 & R \\\\ 021201 & 610.238175  & 1.5 & 40 & R \\\\ 021222 & 631.188612  & 2.9 & 68 & R \\\\ 030825 & 877.367836  & 5.8 & 136 & R \\\\ 030826 & 878.376911  & 4.1 & 114 & R \\\\ 030827 & 879.299621  & 7.5 & 204 & R \\\\ 031101 & 945.196586  & 6.2 & 81,82 & R,I \\\\ 031102 & 946.208050  & 5.6 & 65,64 & R,V \\\\ 031228 & 1002.221443  & 2.0 & 23,23 & I,V \\\\ \\hline \\end{tabular} } \\end{table} ",
        "conclusions": ""
    },
    "0410/astro-ph0410345_arXiv.txt": {
        "abstract": "The intense radiation from a gamma-ray burst (GRB) is shown to be capable of melting stony material at distances up to 300 light years which subsequently cool to form chondrules. These conditions were created in the laboratory for the first time when millimeter sized pellets were placed in a vacuum chamber in the white synchrotron beam at the European Synchrotron Radiation Facility (ESRF). The pellets were rapidly heated in the X-ray and gamma-ray furnace to above 1400\\,\\degr C melted and cooled. This process heats from the inside unlike normal furnaces. The melted spherical samples were examined with a range of techniques and found to have microstructural properties similar to the chondrules that come from meteorites. This experiment demonstrates that GRBs can melt precursor material to form chondrules that may subsequently influence the formation of planets. This work extends the field of laboratory astrophysics to include high power synchrotron sources. ",
        "introduction": "Chondrules are millimetre-sized, stony spherules (Fig.\\,\\ref{fig:met}) that constitute the major component of most chondritic meteorites that originate in the region between Mars and Jupiter and which fall to the Earth. They appear to have crystallised rapidly from molten or partially molten drops and were described \\citep{sorby1877} as ``molten drops in a fiery rain''. The properties of the chondrules and chondrites have been exquisitely deduced from an extensive series of experiments \\citep{wasson2003,hewins2004} and two conferences have been devoted completely to chondrules \\citep{king1983,hewins1996}. The mineralogy of chondrules is dominated by olivine ((Mg,Fe)$_2$SiO$_4$) and pyroxene ((Mg,Fe)SiO$_3$) and there is a wide range of compositions for all elements. A compositional map of a chondrule is shown in Fig.\\,\\ref{fig:metxray}. The magnesium is preferentially found in the dark olivine crystals (Fig.\\,\\ref{fig:metxray}(b)) while the aluminium and calcium are located in the glassy material between the crystals (Fig.\\,\\ref{fig:metxray}(d) and Fig.\\,\\ref{fig:metxray}(e)). The diversity of chondrules is consistent with the melting of heterogeneous precursor solids or dust balls. The age of chondrules indicate they formed very early in the solar system. The calcium-, aluminium-rich inclusions (CAIs) are refractory inclusions in carbonaceous and ordinary chondrites that predate the chondrules by several million years and are the oldest known solid materials produced in the nebula \\citep{swindle1996}. The first 10$^7$ years in the complicated development of the solar system has been comprehensively covered in an attempt to understand the CAI to chondrule time interval of several million years \\citep{cameron1995}.  The presence of volatile elements in the chondrules (Fig.\\,\\ref{fig:metxray}(f)) indicate that the high temperature melting period lasted for a matter of seconds to minutes. Experiments based on chemical and textural compositions of chondrules suggest cooling rates that were much slower than radiative cooling of isolated chondrules and imply they were made in some large quantity in relatively opaque nebular domains \\citep{yu1998,brearley1998}. Volatile elements such as alkalis and sulphur occur in chondrule interiors as primary constituents and indicate that some chondrule precursor materials must have reacted with cool nebula gases at ambient temperatures less than 650\\,K.  The identification of the heat source responsible for chondrule formation is important for the understanding of the formation of planets because of the large amount of chondrules present in chondrite meteorites. The heat source remains uncertain and a critical summary of the heating mechanisms was given by \\citet{rubin2000}. These methods include bipolar outflows \\citep{sekiya1998,lee1998} and shock wave heating of the precursor materials \\citep{wood1988,miura2004,ciesla2002}. Almost all proposed heat sources \\citep{rubin2000,boss1996,jones2000} are local to the solar nebula, one exception being the proposal \\citep{mcbreen1999,duggan2001} that the chondrules were flash heated to melting point by a nearby GRB when the iron all through the precursor material efficiently absorbed X-rays and low energy $\\gamma$-rays. The distance to the source was about 300 light years (or 100\\,pc) for a GRB output of 10$^{53}$\\,erg and was estimated using the minimum value of 2\\,$\\times\\,10^{10}$\\,erg\\,g$^{-1}$ required to heat and melt the precursor grains \\citep{grossman1988,wasson1993}. The role of nearby supernovae that preceded the formation of the solar system have been considered \\citep{cameron1995} along with the serious consequences for life on Earth of nearby supernovae \\citep{ruderman1974,clark1977} and GRBs \\citep{thorsett1995,scalo2002}. The consequences of a nearby GRB on the early solar nebula have not been considered elsewhere. ",
        "conclusions": "The experiment demonstrates that GRBs can melt precursor dust balls to form chondrules in nearby planetary forming systems in agreement with  predictions. Once formed, the chondrules can move through the gas more freely and coagulate to form the building blocks of planets \\citep{cuzzi2001}. GRBs with durations greater than 2\\,s are associated with supernovae in massive stars and the formation of Kerr black holes \\citep{popham1999,mcbreen2002}. The discovery of more than 100 extra-solar giant planets has opened a range of questions regarding the mechanisms of planetary formation \\citep{santos2003,laughlin2000}. The probability that any planetary forming system will be blasted by a nearby GRB has been estimated \\citep{mcbreen1999,scalo2002} to be about 0.1\\,\\%. In the solar neighbourhood, 7\\,\\% of stars with high metallicity harbour a planet whereas less than 1\\,\\% of stars with solar metallicity seem to have a planet \\citep{santos2003}. The GRB method seems to be only one of the mechanisms involved because of the high percentage of stars with planets. There is a further difficulty for the model in the case of our solar system because most chondrules were melted more than once requiring a repeating process \\citep{rubin2000,wasson2003}. However the formation of planetary systems may be enhanced by the presence of a nearby GRB which will form chondrules across the whole nebula at essentially the same time \\citep{mcbreen1999} but would require a range of assumptions about the location of the dust in the disk to account for the properties of chondrules in different classes of chondrites. Advances in the methods of detecting remnant GRBs and planetary systems may reveal a link between them in our galaxy.   A GRB can reveal planetary forming systems in other galaxies because there will be short duration ($\\sim$\\,1\\,hour) bursts of infrared radiation from the melted dust when chondrules form across the whole nebula. These infrared bursts can occur for up to several hundred years after the GRB when the expanding shell of radiation melts the dust in planetary forming systems that happen in its way. However the GRB should be in a nearby galaxy to detect the faint infrared bursts with powerful telescopes such as the overwhelmingly large telescope (OWL) and the James Webb Space Telescope (JWST)."
    },
    "0410/astro-ph0410351_arXiv.txt": {
        "abstract": "{ Identifying the processes that determine the initial mass function of stars (IMF) is a fundamental problem in star formation theory. One of the major uncertainties is the exact chemical state of the star forming gas and its influence on the dynamical evolution. Most simulations of star forming clusters use an isothermal equation of state (EOS). However, this might be an oversimplification given the complex interplay between heating and cooling processes in molecular clouds.  Theoretical predictions and observations suggest that the effective polytropic exponent $\\gamma$ in the EOS varies with density. We address these issues and study the effect of a piecewise polytropic EOS on the formation of stellar clusters in turbulent, self-gravitating molecular clouds using three-dimensional, smoothed particle hydrodynamics simulations. In these simulations stars form via a process we call gravoturbulent fragmentation, i.e., gravitational fragmentation of turbulent gas. To approximate the results of published predictions of the thermal behavior of collapsing clouds, we increase the polytropic exponent $\\gamma$ from 0.7 to 1.1 at a critical density $n_{\\rm c}$, which we estimated to be $2.5\\times10^5\\,\\mathrm{cm^{-3}}$. The change of thermodynamic state at $n_{\\rm c}$ selects a characteristic mass scale for fragmentation $M_{\\rm ch}$, which we relate to the peak of the observed IMF. A simple scaling argument based on the Jeans mass $M_\\mathrm{J}$ at the critical density $n_\\mathrm{c}$ leads to $M_{\\mathrm{ch}}\\propto n_{\\mathrm{c}}^{-0.95}$.  We perform simulations with $4.3\\times10^4\\,\\mathrm{cm^{-3}} < n_{\\mathrm{c}} < 4.3\\times10^7\\,\\mathrm{cm^{-3}}$ to test this scaling argument. Our simulations qualitatively support this hypothesis, but we find a weaker density dependence of $M_{\\mathrm{ch}} \\propto n_{\\mathrm{c}}^{-0.5\\pm0.1}$. We also investigate the influence of additional environmental parameters on the IMF. We consider variations in the turbulent driving scheme, and consistently find $M_{\\mathrm{J}}$ is decreasing with increasing $n_{\\mathrm{c}}$. Our investigation generally supports the idea that the distribution of stellar masses depends mainly on the thermodynamic state of the star-forming gas.  The thermodynamic state of interstellar gas is a result of the balance between heating and cooling processes, which in turn are determined by fundamental atomic and molecular physics and by chemical abundances. Given the abundances, the derivation of a characteristic stellar mass can thus be based on universal quantities and constants. ",
        "introduction": "One of the fundamental unsolved problems in astronomy is the origin of the stellar mass spectrum, the so-called initial mass function (IMF). Observations suggest that there is a characteristic mass for stars in the solar vicinity. The IMF peaks at this characteristic mass which is typically a few tenths of a solar mass. The IMF has a nearly power-law form for larger masses and declines rapidly towards smaller masses \\citep{SCA98a, KRO02, CHA03}.  Although the IMF has been derived from vastly different regions, from the solar vicinity to dense clusters of newly formed stars, the basic features seem to be strikingly universal to all determinations \\citep{KRO01b}. Initial conditions in star forming regions can vary considerably. If the IMF depends on the initial conditions, there would thus be no reason for it to be universal. Therefore a derivation of the characteristic stellar mass that is based on fundamental atomic and molecular physics would be more consistent.   There have been analytical models \\citep{JEA02, LAR69, PEN69, LOW76, SHU77, WHI85} and numerical investigations of the effects of various physical processes on collapse and fragmentation. These processes include, for example, magnetic fields \\citep{BAS95a, TOM96, GAL01}, feedback from the stars themselves \\citep{SIL95,NAK95,ADA96} and competitive coagulation or accretion \\citep{SIL79, LEJ86, PRI95, MUR96, BON01a, BON01b, DUR01}. In another group of models, initial and environmental conditions, like the structural properties of molecular clouds determine the IMF \\citep{ELM83, ELM97a, ELM97b, ELM99a, ELM00a, ELM00c, ELM02a}. \\citet{LAR73a} and \\citet{ZIN84, ZIN90} argued in a more statistical approach that the central-limit theorem naturally leads to a log-normal stellar mass spectrum. Moreover, there are models that connect turbulent motions in molecular clouds to the IMF \\citep[e.g.][]{LAR81, FLE82, PAD95, PAD97, KLE98, KLE00b, KLE01c, PAD02}. The interplay between turbulent motion and self-gravity of the cloud leads to a process we call gravoturbulent fragmentation. The supersonic turbulence ubiquitously observed in molecular gas generates strong density fluctuations with gravity taking over in the densest and most massive regions. Once gas clumps become gravitationally unstable, collapse sets in. The central density increases until a protostellar objects form and grow in mass via accretion from the infalling envelope. For more detailed reviews see \\citet{LAR03} and \\citet{MAC04} .  However, current results are generally based on models that do not treat thermal physics in detail. Typically, they use a simple equation of state (EOS) which is isothermal with the polytropic exponent $\\gamma=1$. The true nature of the EOS remains a major theoretical problem in understanding the fragmentation properties of molecular clouds. Some calculations invoke cooling during the collapse \\citep{MON91, TUR95, WHI95}. Others include radiation transport to account for the heating that occurs once the cloud reaches densities of $n(\\mathrm{H}_2) \\geq 10^{10}\\,\\mathrm{cm}^{-3}$ \\citep{MYH92,BOS93}, or simply assume an adiabatic equation of state once that density is exceeded \\citep{BON94, BAT95}. \\citet{SPA00} showed that radiatively cooling gas can be described by a piecewise polytropic EOS, in which the polytropic exponent $\\gamma$ changes with gas density $\\rho$. Considering a polytropic EOS is still a rather crude approximation. In practice the behaviour of $\\gamma$ may be more complicated and important effects like the temperature of the dust, line-trapping and feedback from newly-formed stars should be taken into account \\citep{SCA98b}. Nevertheless a polytropic EOS gives an insight in the differences that a departure from isothermality evokes.   Recently \\citet{LI03} conducted a systematic study of the effects of a varying polytropic exponent~$\\gamma$ on gravoturbulent fragmentation. Their results showed that $\\gamma$ determines how strongly self-gravitating gas fragments. They found that the degree of fragmentation decreases with increasing polytropic exponent $\\gamma$ in the range $0.2 < \\gamma < 1.4$ although the total amount of mass in collapsed cores appears to remain roughly consistent through this range. These findings suggest that the IMF might be quite sensitive to the thermal physics. Earlier, one-dimensional simulations by \\citet{PAS98} already showed that the density probability distribution of supersonic turbulent gas displays a dependence on the polytropic exponent $\\gamma$. However in both computations, $\\gamma$ was left strictly constant in each case. In this study we extend previous work by using a piecewise polytropic equation of state changing $\\gamma$ at some chosen density. We investigate if a change in $\\gamma$ determines the characteristic mass of the gas clump spectrum and thus, possibly, the turn-over mass of the IMF.  In Sect.~\\ref{sec:therm-prop} we review what is currently known about the thermal properties of interstellar gas. In Sect.~\\ref{sec:ana-app} we approach the fragmentation problem analytically, while in Sect.~\\ref{sec:num} we introduce our computational method. In Sect.~\\ref{sec:res} we discuss gravoturbulent fragmentation in non-isothermal gas. In Sect.~\\ref{sec:char} we analyze the resulting mass distribution. We further investigate the influence of different turbulent driving fields and different scale of driving in Sect.~\\ref{sec:dep}. Finally, in Sect.~\\ref{sec:sum} we summarize. ",
        "conclusions": "\\label{sec:sum} Using SPH simulations we investigate the influence of a piecewise polytropic EOS on fragmentation of molecular clouds. We study the case where the polytropic index $\\gamma$ changes from a value below unity to one above at a critical density $n_{\\mathrm{c}}$. We consider a broad range of values of $n_{\\mathrm{c}}$ around the realistic value to determine the dependence of the mass spectrum on $n_{\\mathrm{c}}$.  Observational evidence predicts that dense prestellar cloud cores show rough balance between gravity and thermal pressure \\citep{BEN89, MYE91}. Thus, the thermodynamical properties of the gas play an important role in determining how dense star-forming regions in molecular clouds collapse and fragment. Observational and theoretical studies of the thermal properties of collapsing clouds both indicate that at densities below $10^{-18}\\,\\mathrm{g\\,cm^{-3}}$, roughly corresponding to a number density of $n_{\\mathrm{c}}=2.5\\times10^5\\,\\mathrm{cm^{-3}}$, the temperature decreases with increasing density. This is due to the strong dependence of molecular cooling rates on density \\citep{KOY00}. Therefore, the polytropic exponent $\\gamma$ is below unity in this density regime. At densities above $10^{-18}\\,\\mathrm{g\\,cm^{-3}}$, the gas becomes thermally coupled to the dust grains, which then control the temperature by far-infrared thermal emission. The balance between compressional heating and thermal cooling by dust causes the temperature to increase again slowly with increasing density. Thus the temperature-density relation can be approximated with $\\gamma$ above unity \\citep{LAR85} in this regime. Changing $\\gamma$ from a value below unity to a value above unity results in a minimum temperature at the critical density. \\citet{LI03} showed that gas fragments efficiently for $\\gamma < 1.0$ and less efficiently for higher $\\gamma$. Thus, the Jeans mass at the critical density defines a characteristic mass for fragmentation, which may be related to the peak of the IMF.  We investigate this relation numerically by changing $\\gamma$ from $0.7$ to $1.1$ at different critical densities $n_{\\mathrm{c}}$ varying from $4.3\\times10^4\\,\\mathrm{cm^{-3}}$ to $4.3\\times10^7\\,\\mathrm{cm^{-3}}$. A simple scaling argument based on the Jeans mass $M_\\mathrm{J}$ at the critical density $n_\\mathrm{c}$ leads to $M_{\\mathrm{ch}}\\propto n_{\\mathrm{c}}^{-0.95}$. If there is a close relation between the average Jeans mass and the characteristic mass of a fragment, a similar relation should hold for the expected peak of the mass spectrum. Our simulations qualitatively support this hypothesis, however, with the weaker density dependency $M_{\\mathrm{ch}} \\propto n_{\\mathrm{c}}^{-0.5\\pm0.1}$. The density at which $\\gamma$ changes from below unity to above unity selects a characteristic mass scale. Consequently, the peak of the resulting mass spectrum decreases with increasing critical density. This spectrum not only shows a pronounced peak but also a powerlaw tail towards higher masses. Its behavior is thus similar to the observed IMF.  Altogether, supersonic turbulence in self-gravitating molecular gas generates a complex network of interacting filaments. The overall density distribution is highly inhomogeneous. Turbulent compression sweeps up gas in some parts of the cloud, but other regions become rarefied. The fragmentation behavior of the cloud and its ability to form stars depend strongly on the EOS. If collapse sets in, the final mass of a fragment depends not only on the local Jeans criterion, but also on additional processes. For example, protostars grow in mass by accretion from their surrounding material. In turbulent clouds the properties of the gas reservoir are continuously changing. In a dense cluster environment, furthermore, protostars may interact with each other, leading to ejection or mass exchange. These dynamical factors modify the resulting mass spectrum, and may explain why the characteristic stellar mass depends on the EOS more weakly than expected.  We also studied the effects of different turbulent driving fields and of a smaller driving scale.  For different realizations of statistically identical large-scale turbulent velocity fields we consistently find the characteristic mass is decreasing with increasing critical mass. However, there are considerable variations.  The influence of the natural stochastic fluctuations in the turbulent flow on the resulting median mass is almost as pronounced as the changes of the thermal properties of the gas. Also when inserting turbulent energy at small wavelengths we see the peak of the mass spectrum decreasing with increasing critical density. %  Our investigation supports the idea that the distribution of stellar masses depends, at least in part, on the thermodynamic state of the star-forming gas.  If there is a low-density regime in molecular clouds where temperature $T$ sinks with increasing density $\\rho$, followed by a higher-density phase where $T$ increases with $\\rho$, fragmentation seems likely to be favored at the transition density where the temperature reaches a minimum. This defines a characteristic mass scale. The thermodynamic state of interstellar gas is a result of the balance between heating and cooling processes, which in turn are determined by fundamental atomic and molecular physics and by chemical abundances. The derivation of a characteristic stellar mass can thus be based on quantities and constants that depend solely on the chemical abundances in a molecular cloud. The current study using a piecewise polytropic EOS can only serve as a first step. Future work will need to consider a realistic chemical network and radiation transfer processes in gas of varying abundances."
    },
    "0410/astro-ph0410398_arXiv.txt": {
        "abstract": "We use the eight year light-curve database from the MACHO (MAssive Compact Halo Objects) project together with infrared colors and magnitudes from 2MASS (the Two Micron All Sky Survey) to identify a sample of 22,000 long period variables in the Large Magellanic Cloud (referred to hereafter as LMC LPVs). A period-luminosity diagram of these stars reveals six well-defined sequences, in substantial agreement with previous analyses of samples from OGLE (Optical Gravitational Lensing Experiment). In our analysis we identify analogues to galactic LPVs in the LMC LPV sample. We find that carbon-dominated AGB stars populate only two of the sequences, one of which includes the Mira variables. The high-luminosity end of the same two sequences are also the location of the only stars with $J-K_s > 2$, indicating that they are enshrouded in dust. The unknown mechanism that drives the variability of stars in the longest period produces different morphology in the period-luminosity diagram as compared to the shortest period sequences, which are thought to be caused by pulsation. In particular, the longest period sequence extends to lower luminosity RGB stars and the luminosity function does not peak among the AGB stars. We point out several features which will constrain new models of the period-luminosity sequences. ",
        "introduction": "Long period variables (LPVs) are giant stars that pulsate with periods ranging from weeks to years. The prototypical case is Mira (Omicron) Ceti, from which all large-amplitude ($>2.5$ magnitudes in V) regular LPVs take their name. LPVs are typically asymptotic giant branch (AGB) stars that are in the late stages of stellar evolution, where the pulsation mechanism(s) -- and the relationships between pulsation, mass loss, and the eventual ejection of the stellar envelope as the star becomes a planetary nebula -- are all poorly understood. During this relatively short AGB phase, stars manage to lose a significant fraction of their total mass, returning processed material to the interstellar medium. A qualitative picture of the mass-loss process invokes instabilities in the helium and hydrogen burning shells surrounding the degenerate core of the star. This provides the energy for a variety of stellar pulsation modes, causing shocks in the envelope which lift material high in the atmosphere.  The temperature is then low enough to allow the formation of dust, which is subsequently ejected from the atmosphere by radiation pressure from the luminous core.   Pulsations -- which affect both the radius and temperature of the star -- together with the shocks and dust formation, combine to produce significant photometric variability.  The study of LPVs has undergone a recent revival with the advent of surveys producing large catalogs of variable stars. Gravitational micro-lensing surveys, such as OGLE -- \\cite{1994ApJ...435L.113P} and OGLE II -- \\cite*{1997AcA....47..319U}; EROS -- \\cite{1995A&A...301....1A}; MACHO -- \\cite{macho}; and MOA -- \\cite{2001MNRAS.327..868B} require that millions of stars be imaged nightly. The variable star catalogs that arise as a by-product contain thousands of stars with light curves covering several thousand days.  As noted by \\cite{cook96} the long period variables in the MACHO database form five parallel sequences in period-luminosity space. \\cite{wood99} classified these and suggested possible underlying mechanisms for the five. Our use of infrared magnitudes in the period-luminosity relations shown in Figure \\ref{perlum} reveals six sequences. Wood's second sequence (``B'' in his notation) is split into our Sequences 2 and 3. In terms of our observed sequences, \\cite{wood99} found that Sequence 1 corresponds to Mira type variables, which are theorized to be pulsating in the fundamental mode. Higher-order pulsations are invoked to explain the shorter period sequences (2, 3, and 4). The sequence that lies at longer-periods and lower luminosities than the Miras, E, showed light curves that suggested that all of these variable stars were eclipsing binaries.  Stars in the longest period sequence, Sequence D, proved more mysterious since the long secondary periods of these stars could not be explained by radial pulsations -- they are longer than the periods of Miras, which themselves are theorized to be pulsating in the fundamental mode. Additionally, they all exhibited multiple periods with the short period associated with Wood's Sequence ``B''. This prompted Wood et al. to propose that Sequence D was caused by the eclipse of the AGB star by a cloud of material around an unseen companion.  More recent studies using OGLE-II data [\\cite{2002MNRAS.337L..31I}; \\cite{kiss+bedding}; \\cite{yoshi}; \\cite{2004MNRAS.347L..83K}; \\cite{Soszynski:2004wn}; \\cite{Groenewegen:2004um}] have extended the work on the period-luminosity sequences. These authors showed that the use of infrared magnitudes (e.g. from 2MASS) for luminosity split the original second sequence into two sequences. Further, their data show that Wood's first two sequences (our first three sequences) exhibit a split in luminosity due to the contribution of red giant branch (RGB) stars. Most previous work on LMC LPVs overplot each period of these predominately multiperiodic stars. However, \\cite{2004MNRAS.347L..83K} find there is no noticeable change in the period-luminosity sequences when plotting only the strongest period, or even when plotting the second strongest. For this reason results can be compared between those works and ones which find only a single period for each star: e.g. \\cite{2002MNRAS.337L..31I}, \\cite{yoshi}, and this paper.  In this paper we use the full eight-year MACHO database to re-examine the LPV period-luminosity relations.  MACHO observations have several advantages over previous studies: the time baseline of the light curves is twice as long (eight years compared to four years), and the colors of MACHO light curves yield superior sampling of photometric variability of RGB stars. As with previous studies, we use 2MASS infrared data to supplement our optical lightcurves. ",
        "conclusions": "The combination of observations over a long time baseline, accurate optical photometry extending below the tip of the RGB, and 2MASS infrared photometry has allowed us to place multiple constraints on the stars inhabiting each of the sequences observed in the period-luminosity diagram of the LMC.  \\begin{itemize} \\item Most of the LPVs in this analysis, which comprises the 52\\% of LMC LPVs with well-determined periods, are SRa stars or Miras. The amplitude criterion ($\\delta V > 2.5$ for Miras) used to differentiate Galactic SRa and Mira stars seems arbitrary, as SRa's and Miras overlap in our diagrams. Stars from Sequence D are SRb stars as previous work has found that they are never monoperiodic.  \\item Our period-luminosity Sequences 3 and 4 both show low luminosity extensions comprised of RGB stars, and have AGB stars that are oxygen rich ($J-K_s < 1.4$). As discussed in \\S \\ref{2mass} this implies that these AGB stars haven't yet undergone the third dredge-up. Compared with the significant population of redder (carbon) stars in Sequences 1 and 2 this implies that older stars are segregated to the latter two sequences. The populations of young, massive AGB stars in Sequences 3 and 4 that are identified in Figure \\ref{cmd} (see \\ref{cmd_section}) support this hypothesis.  \\item The shortest period sequence (Sequence 4) is significantly seperated in period from the longer period sequences but it is otherwise undistinguished.  \\item Sequences D and E tend to follow the LMC luminosity function more closely than the others, indicating that they are more nearly drawn from the general population of giant stars in the LMC. In particular, neither sequence shows a ``bump'' in its luminosity function due to AGB stars, unlike Sequences 1, 2, 3, and 4. These characteristics suggest that a different mechanism produces the variability of these sequences. Sequence E is likely populated by eclipsing binary systems.  \\item Although the physical mechanism that causes Sequence D is still elusive, there are numerous constraints on possible models. \\cite{Soszynski:2004wn} show that this sequence is composed of bright high amplitude Mira and SR variable stars \\emph{and} a dimmer population of small amplitude red variables. This latter population is bluer, extends to lower luminosity, and is broader in period than the RGB star components of Sequences 2, 3, and 4.t The AGB population in Sequence D exhibits higher amplitude pulsation on its shorter period side, opposite what is found for Sequence 1.   \\end{itemize}  We are currently analyzing the $\\sim20,000$ multiply periodic LMC LPVs that are not included in the present sample, with the goal of understanding the relationships between the period-luminosity sequences and the positions in period-luminosity space of stars at various stages of their evolution."
    },
    "0410/astro-ph0410484_arXiv.txt": {
        "abstract": "{We present an analysis of the integrated \\ha\\ emission line profiles for the \\hii\\ region population of the spiral galaxies NGC~1530, NGC~6951 and NGC~3359. We show that $\\sim$70\\% of the line profiles show two or three Gaussian components. The relations between the luminosity (${\\rm log~L_{\\scriptsize\\ha}}$) and non--thermal line width (${\\rm log~\\rm\\sigma_{nt}}$) for the \\hii\\ regions of the three galaxies are studied and compared with the relation found taken all the \\hii\\ regions of the three galaxies as a single distribution. In all of these distributions we find a lower envelope in log~$\\rm\\sigma_{nt}$. A clearer envelope in $\\rm\\sigma_{nt}$ is found when only those \\hii\\ regions with $\\rm\\sigma_{nt} >\\sigma_{s}(13~\\kms)$ are considered, where $\\rm\\sigma_{s}$ is a canonical estimate of the sound speed in the interestellar medium. The linear fit for the envelope is ${\\rm log~L_{\\scriptsize\\ha}}=(36.8\\pm0.7)+(2.0\\pm0.5)~{\\rm log~\\rm\\sigma_{nt}}$ where the \\ha\\ luminosity of the region is taken directly from a photometric \\hii\\ region catalogue. When the \\ha\\ luminosity used instead is that fraction of the \\hii\\ region luminosity, corresponding to the principal velocity component, i.e. to the turbulent non--expanding contribution, the linear fit is ${\\rm log~L_{\\scriptsize\\ha}}=(36.8\\pm0.6)+(2.0\\pm0.5)~{\\rm log~\\rm\\sigma_{nt}}$, i.e. unchanged but slightly tighter. The masses of the \\hii\\ regions on the envelope using the virial theorem and the mass estimates from the \\ha\\ luminosity are comparable, which offers evidence that the \\hii\\ regions on the envelope are virialized systems, while the remaining regions, the majority, are not in virial equilibrium. ",
        "introduction": "Supersonic velocity dispersion is a property of the most luminous \\hii\\ regions which have been extensively studied in the li\\-te\\-ra\\-tu\\-re since Smith \\& Weedman (1970) first observed it. In order to find a physical explanation for the supersonic line widths, Terlevich \\& Melnick (1981) proposed that \\hii\\ regions are gra\\-vi\\-ta\\-tio\\-na\\-lly bound systems and that the observed velocity dispersion is produced by motion of discrete ionized gas clouds in the gravitational field created by the mass distribution inside the \\hii\\ region. This conclusion was based on their observational claim that the relations L~$\\sim\\sigma^{4}$ and R~$\\sim\\sigma^{2}$, (where L and R are the luminosity and radius of the \\hii\\ region and $\\sigma$ the velocity dispersion of the line profile), found in the ste\\-llar systems of ellip\\-ti\\-cal galaxies, bulges of spiral galaxies and globular clusters, are also found in the gaseous \\hii\\ regions.  Several authors have tried to confirm these empirical relations but no agreement has been found between the results for the $\\rm L-\\sigma$ relation from the different studies. The variations in the results have been attributed to several effects: 1) limitations of the observations (Gallagher \\& Hunter 1983; Hippelein 1986); 2) comparison between di\\-ffe\\-rent \\hii\\ region sample criteria (Roy, Arsenault \\& Joncas 1986; Arsenault \\& Roy 1988); besides, the different criteria for the estimates of the radii of the \\hii\\ regions do not allow a valid comparison between the relations found by different authors (Sandage \\& Tamman 1974; Melnick 1977; Gallagher \\& Hunter 1983); 3) asymmetries of the integrated \\hii\\ region line profiles (Arsenault, Roy \\& Boulesteix 1990; Hippelein 1986), while others have found secondary components or asymmetries showing that a single Gaussian fitted to the profile may not be a good representation of the inherent velocity dispersion of the gas (Skillman \\& Balick 1984; Rozas \\etal 1998), in some cases a Voigt function has better characterized a significant fraction of the observed line profiles (Arsenault \\& Roy 1986).  The most complete study up to now, which covers the whole \\hii\\ region population for a single galaxy, is that by Rozas \\etal (1998). This study did not find the relation $\\rm L\\propto\\sigma^{4}$, but obtained a lower envelope in $\\sigma$ in the $\\rm L-\\sigma$ diagram, formed by those \\hii\\ regions which Rozas \\etal (1998) suggested are close to virial equilibrium. Terlevich \\& Melnick (1981) found a scatter in the L-$\\sigma$ relation which they suggested might be due to the different metallicities of the \\hii\\ regions in the sample. Hippelein (1986) and Gallagher \\& Hunter (1983) suggested that an apparent dependence of $\\sigma$ on metallicity will be induced when using a constant electron temperature for the co\\-rrec\\-tion of thermal line broadening, since T$_{\\rm e}$ itself is a function of the metal abundance in the ionized gas. The study of abundances in \\hii\\ regions in the discs of spiral galaxies has shown the existence of gradients, with higher abundances towards the centre of a galaxy (e.g. Vila-Costas \\& Edmunds 1992). These gradients are small in the discs of barred galaxies, but in ge\\-ne\\-ral could affect the L--$\\sigma$ distribution of the \\hii\\ regions located across a galaxy disc.  The hypothetical relations of the form L~$\\sim\\sigma^{4}$ and R~$\\sim\\sigma^{2}$ obtained for virialized systems are based on the un\\-suppor\\-ted assumptions for \\hii\\ regions, that the ratios M/L and L/R$^{2}$, where M is the mass of the region, are constant. Thus, departures from these relations do not provide any evidence for or against the gravitational equilibrium model (Melnick \\etal 1987).  As shown by Rozas \\etal (1998) for the \\hii\\ regions in NGC~4321, a fiducial test to prove the virialization of the \\hii\\ regions must rely on the comparison between the total mass inside the \\hii\\ region and the dynamical mass obtained from the velocity dispersion of the observed line profile using the virial theorem. Such a comparison was made individually for nearby extragalactic \\hii\\ regions: for NGC~604 in the SMC by Yang \\etal (1996) and for 30~Dor in the LMC by Chu \\& Kennicutt (1994). Yang \\etal (1996) computed the total mass of NGC~604 and explained using the virial theorem the basic broadening $\\sigma\\sim(9-13)$~\\kms\\ found by them in most positions of the \\hii\\ region. Chu \\& Kennicutt (1994) could not explain the velocity dispersion of the central core of 30~Dor as due to virial motions, even though they took into account not only the io\\-ni\\-zed mass of the \\hii\\ region but also an estimate of the neutral mass. Rozas \\etal (1998) compared the \\hii\\ region masses obtained from the \\ha\\ luminosity with the masses obtained from the application of the virial equilibrium. While the \\hii\\ regions located well above the envelope in $\\sigma$ in the L-$\\sigma$ diagram, cited above, present major differences between the two mass estimates obtained with these procedures, the \\hii\\ regions located on the envelope show comparable values in the masses derived by the two different methods, which was the argument used by Rozas \\etal (1998) in claiming that the envelope was the locus of virialization.   \\vspace{0.5cm}  In this paper we study the L--$\\sigma$ relation for the \\hii\\ regions of three nearby Milky Way sized barred spirals, NGC~1530, NGC~3359 and NGC~6951. We know the \\ha\\ luminosity of each \\hii\\ region from the catalogues obtained from the con\\-ti\\-nuum subtracted \\ha\\ images of the three galaxies as described in Rozas \\etal (1996) and Rozas \\etal (2000a). The Fabry--P\\'erot observations allow us to extract the line profiles for each \\hii\\ region in the galaxy and to obtain the corresponding ve\\-lo\\-ci\\-ty dispersion.  The large number of \\hii\\ regions in a full galaxy disc forms a good sample to study the L--$\\sigma$ relation because it avoids some points of the controversy of previous studies. Firstly, all the \\hii\\ regions are observed with the same instrument, allowing us to extract their integrated spectra and analyze them with the same procedure. Secondly, since all the \\hii\\ regions are within a single galaxy, distance uncertainties in the L--$\\sigma$ relation are eli\\-mi\\-na\\-ted for each galaxy. Thirdly, errors due to using different criteria for estimating the radius of an \\hii\\ region are reduced because the criteria are the same for all the regions and are obtained automatically in the production of the \\hii\\ region ca\\-ta\\-lo\\-gue for all the objects in the sample. And finally, we have tried to overcome the difficulties arising from any asymmetries in the observed line profiles by fitting them with the optimum num\\-ber of Gaussian components, and proposing via specific identification of the minor components, that the component which indicates the correct value of $\\sigma$ is the central most intense component, which in almost all cases contains more than 75\\% of the total \\ha\\ luminosity of the region. ",
        "conclusions": "  \\begin{itemize}  \\item The major fraction of the integrated line profiles for the most luminous \\hii\\ regions in the three galaxies show the presence of se\\-con\\-da\\-ry Gaussian components. A fraction of $\\sim$70\\% have two or three Gaussian components. These se\\-con\\-da\\-ry components may be blended with the central, most intense peak or be located at relatively high velocity with respect to the central peak, and in general have low intensity values.  \\item The Luminosity {\\it versus} non--thermal velocity dispersion of the principal components of the integrated spectra for the most luminous \\hii\\ regions in NGC~1530, NGC~3359 and NGC~6951 shows an envelope that can be fitted by the li\\-near form, ${\\rm log~L_{\\scriptsize\\ha}}=(36.8\\pm0.7)+(2.0\\pm0.5)~{\\rm log~\\rm\\sigma_{nt}}$, where the luminosity of the \\hii\\ region is taken directly from the calibration ca\\-ta\\-lo\\-gues and ${\\rm log~L_{\\scriptsize\\ha}}=(36.8\\pm0.6)+(2.0\\pm0.5)~{\\rm log~\\rm\\sigma_{nt}}$, where the luminosity is a fraction of the \\hii\\ region luminosity, determined as the fraction of the line profile integrated intensity in the main Gaussian peak of the emission spectrum.  \\item We have estimated the stellar mass content within the \\hii\\ regions and found that it represents a small fraction of the total mass of the region. Although a precise determination of the fraction is difficult to obtain due to the uncertainties involved, an upper limit of 10\\% of the total mass of the region in stellar form is a conservatively high value that takes into account all the uncertainties.  \\item We have compared the virially derived mass with the mass for the \\hii\\ region obtained from the observed \\ha\\ lu\\-mi\\-no\\-si\\-ty for two groups of \\hii\\ regions in the log~L$_{\\scriptsize\\ha}$--log~$\\rm\\sigma_{nt}$ diagram: the \\hii\\ regions located on the envelope and \\hii\\ regions located well above the envelope in velocity dispersion (i.e. well below the envelope in luminosity). The vi\\-ria\\-lly estimated mass for the \\hii\\ regions located well above the envelope are significantly bigger than the masses obtained from the \\ha\\ luminosity, indicating that the velocity dispersions for these \\hii\\ regions include contributions from non equilibrium processes that affect the internal ki\\-ne\\-ma\\-tics of these regions.  \\item  For the \\hii\\ regions located on the envelope, the values for virial masses are similar to those for the masses obtained from the \\ha\\ luminosity when the contamination of other mechanisms in the broadening of the central peak is taken into account. This result shows that those \\hii\\ regions in a defined high luminosity range with the lowest velocity dispersion can be taken as virialized systems, while the remaining regions are not in virial equilibrium.  \\end{itemize}   \\appendix"
    },
    "0410/astro-ph0410167_arXiv.txt": {
        "abstract": "Observational studies of cluster evolution over moderate redshift ranges (to $z$ $\\sim$ 1) are a powerful tool for constraining cosmological parameters, yet a comprehensive knowledge of the properties of these clusters has been hitherto unattained.  Using a highly efficient optical selection technique, the Red-Sequence Cluster Survey (RCS) has unearthed a large sample of high redshift cluster candidates.  All six of the clusters from this sample which have been observed with the Chandra X-Ray Observatory were detected in the X-ray.  These Chandra follow-up observations (0.64 $<$ $z$ $<$ 1.0) indicate that the clusters are systematically less luminous than their similarly rich, X-ray selected counterparts at lower redshifts, though they are consistent with standard $L_x-T_x$ relationships. Comparisons with X-ray selected samples suggest that the discrepancy may be due in part to systematic differences in the spatial structure of the X-ray emitting gas.  Our initial results from Chandra follow-up observations of six RCS clusters are presented, including $\\beta$ model parameters and spectral information. ",
        "introduction": "The Red-Sequence Cluster Survey \\citep{gladders} is a 90 square degree optical survey performed at CFHT and CTIO using $R_C$ and $z$' filters.  It utilizes the red sequence of elliptical galaxies to find galactic overdensities on the sky \\citep{gladders2}.  The color information also guards against projection effects, and provides photometric redshift estimates.  The survey has identified between 3500 and 4000 clusters total (in the redshift range 0.2 $<$ $z$ $<$ 1.2), over 1500 of which are at least as optically rich as Abell class 0 clusters \\citep{gladders}.  Simulations indicate that the survey is complete to Abell Class 1 at a redshift of 1 for blue fractions less than 0.45, and that false detection rates are less than 5\\%, which is significantly lower than that of single-passband optical cluster surveys \\citep{postman, donahue2,  gilbank}.  Figure \\ref{figure1} shows an example of a $z$=0.773 RCS cluster.   \\begin{figure} \\includegraphics*[width=3in]{fig1_color.eps} \\label{figure1} \\caption{Top: HST image of RCS 0224-0002 spectroscopically confirmed at $z$=0.773, with overlayed X-ray contours.  The outer gravitational lensing arc has been spectroscopically confirmed at a redshift of $z$=4.87 \\citep{gladders3}, with the others expected to fall within the range 1.8$<$ $z$ $<$3.6  The five linearly spaced contours indicate values between $9.4\\times10^{-6}$ and $2.1\\times10^{-5}$ counts/$\\rm{pix}^2$/$\\rm{cm}^2$ and were created using a Gaussian smoothed (5 pixel FWHM) 0.29-7.0 keV Chandra flux image. Bottom: Adaptively smoothed flux image of RCS 0224-0002 created with the CIAO tool csmooth.  The boxed area indicates the region shown in the upper panel.} \\end{figure} ",
        "conclusions": "The Red-Sequence Cluster Survey has uncovered a large sample of high redshift cluster candidates.  All of the RCS clusters observed with Chandra were detected at greater than 3$\\sigma$ significance in the X-ray.  This is evidence that the RCS method is a reliable way to detect high redshift clusters.  The X-ray properties of RCS clusters are consistent with a standard $L_x-T_x$ relationship.  This leads us to believe that the X-ray gas in these clusters is in a physical state very similar to that found in X-ray selected clusters.  The results of detailed X-ray analysis imply that though these clusters have extended structures of galaxies, they possess relatively small virialized cores.  These findings are similar to those from a number of previously conducted optical surveys \\citep{donahue2, lubin, gilbank}.  Optical, high redshift cluster surveys regularly find candidates with lower X-ray luminosities than those of X-ray selected clusters.  Our study suggests that the RCS survey is detecting a different population of clusters from that found in X-ray selected samples, possibly including ultimately very rich clusters which are currently in the early stages of virialization.  Similar objects have recently been described by \\citet{ford}.  Such systems may be expected to be more common at high redshifts in a low matter density universe.  Optical surveys are much more sensitive to such systems than are X-ray surveys, especially if any related filamentary structures lie along the line of sight (though it is unlikely that all of the RCS cluster candidates that were followed-up with Chandra are associated with such structures).   Upcoming Chandra observations of five new and two previously observed RCS clusters in the next AO, along with ongoing velocity dispersion measurements and weak lensing analysis should help to provide more definitive constraints on the masses, dynamical states, and gas content of high redshift optically selected samples of galaxy clusters."
    },
    "0410/hep-ph0410178_arXiv.txt": {
        "abstract": "{Dark matter may be composed of superWIMPs, superweakly-interacting massive particles produced in the late decays of other particles.  We focus here on the well-motivated supersymmetric example of gravitino LSPs.  Gravitino superWIMPs share several virtues with the well-known case of neutralino dark matter: they are present in the same supersymmetric frameworks (supergravity with $R$-parity conservation) and naturally have the desired relic density.  In contrast to neutralinos, however, gravitino superWIMPs are impossible to detect by conventional dark matter searches, may explain an existing discrepancy in Big Bang nucleosynthesis, predict observable distortions in the cosmic microwave background, and imply spectacular signals at future particle colliders. }  \\normalsize\\baselineskip=15pt ",
        "introduction": "\\label{sec:introduction}  In recent years, there has been tremendous progress in understanding the universe on the largest scales.  In particular, the energy density in non-baryonic dark matter is known to be~\\cite{Spergel:2003cb} \\begin{equation} \\OmegaDM = 0.23 \\pm 0.04 \\end{equation} in units of the critical density.  At the same time, we have no idea what the microscopic identity of non-baryonic dark matter is.  The dark matter problem therefore provides precise, unambiguous evidence for new physics and has motivated new particles such as axions~\\cite{Peccei:1977ur,Wilczek:pj,Weinberg:1977ma}, neutralinos~\\cite{Goldberg:1983nd,Ellis:1983ew}, Q balls~\\cite{Kusenko:1997si}, wimpzillas~\\cite{Chung:1998ua}, axinos~\\cite{Covi:1999ty}, self-interacting dark matter~\\cite{Spergel:1999mh}, annihilating dark matter~\\cite{Kaplinghat:2000vt}, Kaluza-Klein dark matter~\\cite{Servant:2002aq,Cheng:2002ej}, branons~\\cite{Cembranos:2003mr,Cembranos:2003fu}, and many others.  Here we review a new class of dark matter candidates: superWIMPs, superweakly-interacting massive particles produced in the late decays of other particles~\\cite{Feng:2003xh,Feng:2003uy,Feng:2003nr}. SuperWIMPs have several strong motivations: \\begin{itemize} \\item They are present in well-motivated frameworks for new physics, including models with supersymmetry (supergravity with $R$-parity conservation) and extra dimensions (universal extra dimensions with KK-parity conservation). \\item Their relic density is naturally in the right range to be dark matter without the need to introduce and fine-tune new energy scales. \\item They can explain an existing anomaly, namely the observed underabundance of $^7$Li relative to the prediction of standard Big Bang nucleosynthesis. \\item They have rich implications for early universe cosmology and imply spectacular signals at the Large Hadron Collider (LHC) and the International Linear Collider (ILC). \\end{itemize} ",
        "conclusions": "\\label{sec:conclusions}  Although the implications of supergravity for cosmology and particle physics have been considered in great detail for decades, most work has been centered on scenarios in which the LSP is a standard model superpartner.  Here we have explored the gravitino LSP scenario. Recent work has found significant cosmological motivations for this possibility, as the gravitino may explain dark matter, and the scenario may resolve current difficulties in Big Bang nucleosynthesis.  Despite the extremely weak couplings of superWIMPs, this scenario has striking implications for cosmology and particle physics.  In cosmology, the scenario predicts possible $\\mu$ distortions in the CMB spectrum at a level that will be probed by the planned DIMES mission. Such distortions would provide significant corroborating evidence for superWIMP dark matter.  The superWIMP scenario also implies long-lived charged particles, which will provide a spectacular signal at colliders.  As discussed above, such particles may be trapped and their decays studied in detail.  There are many significant implications of such studies.  In the case of supersymmetry and gravitino superWIMPs, these implications have been considered in detail in Refs.~\\cite{Buchmuller:2004rq,Feng:2004gn}.  Briefly, simply by counting the number of slepton decays as a function of time, the slepton lifetime may be determined with high accuracy.  Given thousands of sleptons, we expect a determination at the few percent level.  The slepton decay width of \\eqref{sfermionwidth} is a simple function of the slepton and gravitino masses, and the slepton mass will be constrained by analysis of the collider event kinematics. A measurement of the slepton width therefore implies a high precision measurement of the gravitino mass and, through \\eqref{gravitinomass}, the supersymmetry breaking scale $F$.  Such measurements will provide precision determinations of the relic density of superWIMP gravitino dark matter, the contribution of supersymmetry breaking to vacuum energy, and the opportunity for laboratory studies of late decay phenomena relevant for Big Bang nucleosynthesis and the cosmic microwave background.  The gravitino mass may also be determined, although not necessarily on an event-by-event basis, by measuring the energy of slepton decay products.  This provides a consistency check.  Alternatively, these two methods, when combined, determine not only $m_{\\gravitino}$, but also the Planck mass $\\mstar$.  This then provides a precision measurement of Newton's constant on unprecedentedly small scales, and the opportunity for a quantitative test of supergravity relations."
    },
    "0410/astro-ph0410747_arXiv.txt": {
        "abstract": "We present results from a radial velocity study of two runaway O-type stars, HD~14633 (ON8.5~V) and HD~15137 (O9.5~III(n)). We find that HD~14633 is a single-lined spectroscopic binary with an orbital period of 15.4083 days.   The second target HD~15137 is a radial velocity variable and a possible single-lined spectroscopic binary with a period close to 1 month. Both binaries have large eccentricity, small semiamplitude, and a small mass function.  We show the trajectories of the stars in the sky based upon an integration of motion in the Galactic potential, and we suggest that both stars were ejected from the vicinity of the open cluster NGC~654 in the Perseus spiral arm.  The binary orbital parameters and runaway velocities are consistent with the idea that both these stars were ejected by supernova explosions in binaries and that they host neutron star companions.  We find that the time-of-flight since ejection is longer than the predicted evolutionary timescales for the stars, which may indicate that they have a lower mass than normally associated with their spectral classifications and/or that their lives have been extended through rapid rotation. ",
        "introduction": "%  There are two competing theories to explain the origin of the massive OB-runaway stars.  The model first suggested by \\citet{zwi57} and \\citet{bla61} proposes that these stars were originally the binary companions of a star that exploded in a supernova and that the linear momentum of a runaway star balances the momentum lost in the explosion.  Since mass ratio reversal probably occurs prior to the explosion, many runaways should still be binaries with a neutron star or black hole companion, unless the system was disrupted by an asymmetric kick velocity imparted to the remnant during the supernova \\citep{bp95}. A second model proposes that close gravitational encounters during the young, high stellar number density epoch after cluster formation can lead to ejections through encounters with hard binaries \\citep{len88}. This model predicts that most runaways will be single stars, although some close binaries can be ejected in exceptional circumstances.  \\citet{gie86} made a radial velocity survey of bright, northern sky runaway stars and found that most were indeed radial velocity constant, implying that they were not members of binary systems.   More recently \\citet*{hoo00} explored the motions and origins of runaways using proper motion data from the {\\it Hipparcos Satellite} \\citep{per97}, and they found examples of ejection by both mechanisms.  Here we present new radial velocity measurements for two northern sky runaway stars, HD~14633 and HD~15137. HD~14633 is classified as a nitrogen strong ON~8V star \\citep{wal72}. It appears at Galactic coordinates $l=140\\fdg78$ and $b=-18\\fdg20$, indicating a position nearly 800~pc below the the Galactic plane. The spectral lines have a moderate projected rotational velocity with $V\\sin i$ estimates of 111 km~s$^{-1}$ \\citep{con77a}, 110 km~s$^{-1}$ \\citep{sch88}, and 134 km~s$^{-1}$ \\citep{how97}.  \\citet{rog74} found that the star was a single-lined spectroscopic binary with a period of 15.335 days and an orbital eccentricity of $e=0.68$. However, subsequent analysis by \\citet{bol78} did not confirm the initial orbital parameters, and \\citet{bol78} suggested that the binary might have a nearby third star that modulates the velocity curve.  There is no known visual companion to HD~14633 \\citep{mas98}.  Additional spectroscopic observations by \\citet{sto82} showed little evidence of velocity variability.  The second target is the star HD~15137 that \\citet{gie87} categorized as a field O-star, but we show below (\\S4) that its peculiar velocity is large enough that the star should also be grouped with the runaway stars. It appears in a similar part of the sky as HD~14633 at $l=137\\fdg46$ and $b=-7\\fdg58$. \\citet{wal73} classified HD~15137 as O9.5~II-III~(n), where the suffix (n) indicates broad lines.  \\citet{con77a} reported observing partially resolved double lines in their spectrum.  However, \\citet{how97} analyzed a single high dispersion spectrum from the {\\it International Ultraviolet Explorer Satellite (IUE)} and used a cross-correlation method to find the projected rotational velocity, $V\\sin i = 336$ km~s$^{-1}$. They caution that the cross-correlation function is broad, asymmetric, and difficult to measure.   We show below that the star is indeed broad-lined, and it may display rapid line profile variability normally associated with nonradial pulsation \\citep{how93,kam97}. \\citet{con77} suggest that the stellar radial velocity is variable. There is no evidence of a nascent cluster nearby \\citep{dew04}.  Here we present new radial velocities (\\S2) based upon high S/N CCD spectroscopy of these two runaways.  We give new orbital elements for HD~14633 (\\S3.1) and a tentative binary interpretation for HD~15137 (\\S3.2).  We use radial velocity and proper motion data to calculate the Galactic trajectories of both stars, and we suggest that both originated in or near the open cluster NGC~654 (\\S4).  We argue that both runaways were probably ejected by a supernova in a binary and that their unseen companions are probably neutron stars (\\S5). ",
        "conclusions": "%  OB runaway stars are probably ejected by one of two mechanisms, sudden mass loss during a supernova explosion in a binary or a close gravitational encounter involving binaries \\citep{gie86,hoo00}. The supernova theory predicts that runaways will either be single stars (in which the binary was disrupted due to a large, asymmetric kick velocity imparted during the supernova) or binaries with neutron star or black hole companions (such as the high mass X-ray binaries).  On the other hand, the gravitational encounter theory suggests that most runaways will be single objects, although in rare cases hard binaries of mass ratio near unity are ejected.  Our radial velocity study has demonstrated that HD~14633 and possibly HD~15137 are binary stars with low mass companions. If we suppose that the masses of the primary are $23~M_\\odot$ for HD~14633 \\citep{kee84} and $24~M_\\odot$ for HD~15137 \\citep*{vac96}, then the minimum masses of the companion derived from the orbital mass function (Tables 3 and 4) will be $1.3~M_\\odot$ and $1.5~M_\\odot$, respectively (for an orbital inclination of $90^\\circ$). These masses are close to the $1.35~M_\\odot$ value found for most neutron stars in binaries \\citep{tho99}. These runaways may be the first examples of the long sought ``quiet'' massive X-ray binaries, i.e., those with wide separations in which wind accretion is too weak to power an accretion disk X-ray source \\citep{vdh76}. We searched for evidence of a companion spectrum in both cases using a Doppler tomography algorithm \\citep{bag94}, but no spectral features were found.  A faint, low mass, main sequence star could easily remain hidden in the glare of an O-star (for example, the magnitude difference is $\\triangle V \\approx 8$ between such O-stars and a F3~V companion of mass $1.4~M_\\odot$). Nevertheless, we doubt that these systems are extreme mass ratio binaries containing an O- and F-type star, since no such systems are known among the O-stars and since such systems would probably be disrupted in close gravitational encounters leading to ejection.  Both HD~14633 and HD~15137 have many characteristics in common with the massive X-ray binary and microquasar, LS~5039 \\citep{mcs04}. All are runaway objects with very eccentric orbits and small orbital mass functions.   LS~5039 has a much shorter period (4.4267 days) and the smaller semimajor axis results in a modestly dense wind in the vicinity of the orbiting neutron star, so that LS~5039 is a weak X-ray source.  In contrast, the longer period systems HD~14633 and HD~15137 will have very rarefied winds close to their neutron star companions, and consequently their accretion fluxes are expected to be extremely faint (perhaps also as the result of centrifugal inhibition of accretion; \\citealt*{ste86}). Neither system is listed in the ROSAT All-Sky Survey Faint Source Catalogue \\citep{vog00}.  Furthermore, neither system appears to be associated with an {\\it EGRET} $\\gamma$-ray source \\citep{har99}, nor are they known radio sources \\citep*{val85,wen95,say96}.  Thus, wind accretion onto a neutron star in these systems must be too feeble to produce the high energy phenomena associated with other massive X-ray binaries.  \\citet{mcs04} found that the supernova mass loss prediction for LS~5039 was different depending on whether the calculation was based on orbital eccentricity or runaway velocity, and they argued that both the eccentricity and runaway velocity can be explained if a significant asymmetric kick velocity was imparted to the neutron star during formation. A similar conclusion can be derived for HD~14633 and HD~15137. If we use the expressions for supernova mass loss given by \\citet*{nel99} and adopt the primary masses given above and secondary masses of $1.4~M_\\odot$, then the predicted supernova mass loss is $17~M_\\odot$ and $13~M_\\odot$ for HD~14633 and HD~15137, respectively, based upon their observed eccentricities. On the other hand, the supernova mass loss estimates are $6.9~M_\\odot$ and $6.4~M_\\odot$, respectively, based upon the relative runaway velocities between star and cluster from the models given in \\S4.  These significant differences suggest that both systems suffered kick velocities at birth that substantially altered the eccentricity.  The supernova mass loss estimates from the runaway velocities should be more reliable since the runaway velocities are less affected by kicks \\citep{bp95}.  Two other features of these stars also link them to supernova ejections. First, HD~14633 is a well known nitrogen rich ON star \\citep{wal72,sch88}, and \\citet{mcs04} have shown that massive X-ray binary LS~5039 also shares this trait.  \\citet{mcs04} suggest that the nitrogen enrichment is the result of mass transfer of CNO-processed gas from the supernova progenitor prior to the explosion, although rotationally induced mixing may also play a role.   Second, HD~15137 is a very rapid rotator, a characteristic shared with many other OB runaway stars \\citep{bla93}. Mass transfer prior to the supernova may lead to a spin up of the mass gainer, and this process may be responsible for the largest class of massive X-ray sources, the rapidly rotating, Be X-ray binaries \\citep{coe00}.  Both runaways appear to have been ejected from the Perseus spiral arm, and our analysis of their motions in the Galaxy (\\S4) indicates a probable origin in the open cluster NGC~654 in the Cas~OB8 association. The cluster's age is approximately 14~Myr \\citep{hue93}, and the cluster contains a number of early B-type stars and two massive supergiants (HD~10494, F5~Ia, and BD$+61^\\circ315$, A2~Ib).  \\citet{gar92} include the nearby O-star BD$+60^\\circ261$ (O7.5~III(n)((f)); \\citealt{wal73}) as a cluster member. The Cas~OB8 association has a diameter of approximately 85~pc and contains several other clusters including NGC~581, 659, and 663 \\citep{gar92}, which have slightly greater ages of 22, 35, and 16~Myr, respectively, according to the Webda database\\footnote{Maintained by J.-C. Mermilliod at http://obswww.unige.ch/webda/webda.html} \\citep{mer03}. The time-of-flight for HD~15137 (10~Myr) suggests that the star was ejected from NGC~654 when the cluster was approximately 4~Myr old, which may be consistent with the evolutionary timescale required for a supernova progenitor. However, the main sequence lifetime of a star of $24~M_\\odot$ is approximately 6.7~Myr \\citep{sch92}, which is less than the time-of-flight for HD~15137.  The situation is even more discrepant for HD~14633 which has a time-of-flight of at least 12~Myr (see also the extreme case of the runaway star HD~93521; \\citealt{how93}).  It is difficult to reconcile these long travel times with the expected short lifetimes of O-stars. There are several possible explanations.  First, the runaways may have been rejuvenated by mass transfer just prior to the supernova explosion, which would reset their effective zero-age times to an epoch just prior to ejection.  Second, at least HD~15137 is a rapid rotator, and fast rotation may help to mix gas and extend the main sequence lifetime of massive stars \\citep{heg00,mey00}. Third, these stars may be over-luminous for their mass in same way as found for some massive X-ray binaries \\citep{kap01}, so that their masses are lower and evolutionary lifetimes longer than simple estimates suggest.  The orbital properties of these two runaway binaries, their small mass functions, and their probable origin in a cluster containing evolved, massive stars all indicate that these stars were ejected during a supernova explosion in a binary. They are not known X-ray sources, due presumably to their large semimajor axes and low wind accretion rates, but it is possible that they exhibit transient X-ray emission when their neutron stars pass through the densest stellar wind regions near the periastron orbital phase. It is important to pursue radial velocity studies of other OB runaway stars to search for additional instances of such low amplitude binary systems.  Only then will we determine the relative importance of the supernova and close encounter ejection processes for the kinematics of massive stars."
    },
    "0410/astro-ph0410571_arXiv.txt": {
        "abstract": "{ Herbig Ae stars are young A-type stars in the pre-main sequence evolutionary phase with masses of $\\sim$1.5--3\\,M$_\\odot$. They show rather intense surface activity (Dunkin et~al.\\ \\cite{Du97}) and infrared excess related to the presence of circumstellar disks. Because of their youth, primordial magnetic fields inherited from the parent molecular cloud may be expected, but no direct evidence for the presence of magnetic fields on their surface, except in one case (Donati et~al.\\ \\cite{Do97}), has been found until now. Here we report observations of optical circular polarization with FORS\\,1 at the VLT in the three Herbig Ae stars HD\\,139614, HD\\,144432 and HD\\,144668. A definite longitudinal magnetic field at 4.8\\,$\\sigma$ level, $\\left<{\\cal B}_z\\right>$=$-$$450\\pm93$\\,G, has been detected in the Herbig Ae star HD\\,139614. This is the largest magnetic field ever diagnosed for a Herbig Ae star. A hint of a weak magnetic field is found in the other two Herbig Ae stars, HD\\,144432 and HD\\,144668, for which magnetic fields are measured at the $\\sim$1.6\\,$\\sigma$ and $\\sim$2.5\\,$\\sigma$ level respectively. Further, we report the presence of circular polarization signatures in the \\ion{Ca}{ii}~K line in the V Stokes spectra of HD\\,139614 and HD\\,144432, which appear unresolved at the low spectral resolution achievable with FORS\\,1. We suggest that models involving accretion of matter from the disk to the star along a global stellar magnetic field of a specific geometry can account for the observed Zeeman signatures. ",
        "introduction": "\\label{sec:intro}  It is generally accepted that star formation occurs via accretion. A number of Herbig Ae stars and classical T\\,Tauri stars are surrounded by active accretion disks and, probably, most of the excess emission seen at various wavelength regions can be attributed to the interaction of the disk with a magnetically active star {\\change(e.g.\\ Muzerolle et al.\\ \\cite{Mu04})}. This interaction is generally referred to as magnetospheric accretion. T\\,Tauri stars are solar type pre-main sequence stars, and from detailed magnetohydrodynamic models it is expected that their strong magnetic fields funnel material from the disk onto the star and launch a collimated bipolar outflow {\\change (e.g. Shu et al.\\ \\cite{Sh00})}.  Recent observations of the disk properties of intermediate mass Herbig Ae stars suggest a close parallel to T\\,Tauri stars, revealing the same size range of the disks, similar optical surface brightness and similar structure consisting of inner dark disk and a bright ring. A large number of T\\,Tauri stars are known to drive micro-jets as well as parsec-scale Herbig-Haro outflows (McGroarty \\& Ray \\cite{McGR04}). Similar outflows have been newly discovered also from intermediate mass Herbig Ae stars (McGroarty et~al.\\ \\cite{McG04}). Spectacular observations in the last years have been carried out by the Space Telescope Imaging Spectrograph on board of the Hubble Space Telescope, providing the first detections in Ly$\\alpha$ of micro-jets associated with two of the nearest Herbig Ae stars, HD\\,163296 and HD\\,104237 (Devine et~al.\\ \\cite{De00}; Grady et~al.\\ \\cite{Gr04}).  Although magnetic fields are believed to play a crucial role in controlling accretion onto, and winds from, Herbig Ae stars, contrary to the advance achieved in magnetic studies of T\\,Tauri stars, there is still no observational evidence demonstrating the strength, extent, and geometry of their magnetic fields. There were several attempts in the past to measure magnetic fields in these stars (e.g., Glagolevskij \\& Chountonov \\cite{GC01}; Catala et~al.\\ \\cite{Ca93}), but the only detection has been reported for the optically brightest Herbig Ae star HD\\,104237 viewed with the disk close to face-on (Donati et~al.\\ \\cite{Do97}). Using the Anglo-Australian Telescope with the UCL Echelle Spectrograph and a visitor polarimeter a marginal circular polarization signature has been observed in metallic lines indicating a longitudinal magnetic field of $\\sim$50\\,G.  There are certain reasons why the detection of magnetic fields in Herbig Ae stars is so difficult. The young intermediate mass stars as well as more massive stars tend to be surrounded by large amounts of circumstellar gas and dust, making it hard to measure magnetic fields at visual wavelengths. {\\change A number of spectral lines show circumstellar absorptions with complex profiles, and some lines appear in emission or exhibit P\\,Cyg-type profiles. The optical spectra of Herbig Ae stars contain fewer absorption lines compared to cooler stars, and the rotational velocities of A stars as a class tend to be higher. In addition, Herbig Ae stars are less common than T\\,Tauri stars, since the IMF favors the production of low mass stars, and, therefore, the sample of Herbig Ae stars includes rather faint objects.} Yet, the lack of magnetic field detections demonstrates that magnetic fields in Herbig Ae stars should be rather weak, of the order of a few hundred Gauss and less, and a definite detection of such weak magnetic fields in faint objects certainly requires the collecting areas of large aperture telescopes of 8-10\\,m class coupled with highly efficient spectrographs to achieve a signal-to-noise ratio of at least 1,000. ",
        "conclusions": "\\label{sec:disc} The basic data of the studied Herbig Ae stars are presented in Table~\\ref{tab:overview}. The first eight columns indicate, in order, the HD number of the star, another identifier, the visual magnitude, the spectral type, the stellar parameters and their sources. The other columns list the infrared colour excess $(V-L)_{obs}$, the linear polarization $P$ (Yudin \\cite{Yu00}) and our determination of the mean longitudinal magnetic field $\\left<{\\cal B}_z\\right>$. {\\change The mean longitudinal magnetic field is the average over the stellar hemisphere visible at the time of observation of the component of the magnetic field parallel to the line of sight, weighted by the local emergent spectral line intensity. It is diagnosed from the slope of a linear regression of $V/I$ versus the quantity $-g_{\\rm eff} \\Delta\\lambda_z \\lambda^2 \\frac{1}{I} \\frac{{\\mathrm d}I}{{\\mathrm d}\\lambda} \\left<B_z\\right> + V_0/I_0$ (Fig.\\,\\ref{fig:Fits}). This procedure is described in detail by Bagnulo et al.\\ (\\cite{Ba02}) and Hubrig et al.\\ (\\cite{Hu04a}). Our experience with a study of a large sample of magnetic and non-magnetic Ap and Bp stars revealed  that this regression technique is very robust and that detections with $B_z > 3\\,\\sigma_z$ result only for stars possessing magnetic fields (Hubrig et al.\\ \\cite{Hu04b}).}   The differences in the measured infrared colour excess and linear polarization between these stars may be interpreted in terms of different inclination of their rotation axis to the line of sight, and thus of their circumstellar (CS) disks with respect to the observer. HD\\,139614 and HD\\,144432 are very likely observed close to the rotational pole (face-on for the CS disk; Meeus et~al.\\ \\cite{Me98}) {\\change whereas HD\\,144668 is possibly observed at an intermediate angle of $68^\\circ>i>45^\\circ$ (Natta \\& Whitney \\cite{NW00})}. The pre-main sequence nature of HD\\,139614 and HD\\,144432 has been studied several years ago (Dunkin et~al.\\ \\cite{Du97}). Both stars show emission-line characteristics of H$\\alpha$, Na~I D and He~I 5876\\,\\AA{} indicating accretion activity and winds. No age determination has been found in the literature for HD\\,139614. Although this star has been mentioned in several studies as a Vega-type star (Sylvester et~al.\\ \\cite{Sy96}), it is certainly a member of the Herbig Ae group, considering its spectral energy distribution in the infrared and its emission line characteristics. For HD\\,144432, the age between 1 and 3 million years has been determined quite recently (Perez et~al.\\ \\cite{Pe04}). The star HD\\,144668 is much younger with an age of about 0.5 million years (van den Ancker et~al.\\ \\cite{vdA98}). A search for a magnetic field was previously only performed for HD\\,144432. Using 26 metallic lines Glagolevskij \\& Chountonov (\\cite{gc01}) measured $\\left<{\\cal B}_z\\right>$=$-$$800\\pm600$\\,G. This measurement was severely  hampered by the rather large width of the measured spectral lines as this star is rotating relatively fast.  Spectra of the studied Herbig Ae stars in integral light in the spectral region around the \\ion{Ca}{ii}~K line and close-by H Balmer lines are presented in Fig.\\,\\ref{fig:StokesI2}. From our data, we derive for the star HD\\,139614, viewed with the disk close to face-on, the mean longitudinal field $\\left<{\\cal B}_z\\right>$=$-$$450\\pm93$~G. This is the largest mean longitudinal magnetic field ever diagnosed for a Herbig Ae star. For the stars HD\\,144432 and HD\\,144668 we measure respectively $\\left<{\\cal B}_z\\right>$=$-$$94\\pm60$\\,G and $\\left<{\\cal B}_z\\right>$=$-$$118\\pm48$\\,G.  \\begin{figure} \\centering {\\includegraphics[width=0.4\\textwidth]{Gh093_f2.eps}} \\caption{ Normalized Stokes I spectra of the three Herbig Ae stars. The individual spectra are displaced by 0.65 with respect to each other. } \\label{fig:StokesI2} \\end{figure}  \\begin{figure} \\centering {\\includegraphics[width=0.4\\textwidth]{Gh093_f3.eps}} \\caption{ Stokes V/I spectra of the studied stars around the \\ion{Ca}{ii} doublet. The individual spectra are shifted by 0.8 with respect to each other. } \\label{fig:StokesV2} \\end{figure}   In the Stokes V spectra of HD\\,139614 and HD\\,144432 the interesting but rather unexpected fact is the presence of circular polarization signatures in the \\ion{Ca}{ii}~K line at $\\lambda\\,3933.7$\\,\\AA{}, which appear unresolved at the low spectral resolution achievable with FORS\\,1 (Fig.\\,\\ref{fig:StokesV2}). The line \\ion{Ca}{ii}~H at $\\lambda\\, 3968.5$\\,\\AA{} is blended with the Balmer line H$\\epsilon$. As already noted, Herbig Ae stars are surrounded by circumstellar disks. {\\change Recent magnetospheric accretion models for these stars assume a dipolar magnetic field geometry and accreting gas from a circumstellar disk falling ballistically along the field lines to the stellar surface (Muzerolle et al. 2004). According to a previous study of the \\ion{Ca}{ii} emissivity from dense envelopes near Herbig stars which used photoionization calculations (Hamann \\& Persson \\cite{HP92}), it is very likely that the \\ion{Ca}{ii} lines form very near the photosphere. Probably, models involving accretion of circumstellar matter from the disk to the star along a global stellar magnetic field of a specific geometry can account for the Zeeman signatures observed in the \\ion{Ca}{ii} lines. We note, however, that the actual geometry of magnetic fields in Herbig Ae stars is completely unknown.} Obviously, further spectropolarimetric observations at higher resolution in the spectral region around the \\ion{Ca}{ii}~H and K lines as well as hydrogen Balmer lines would provide an essential clue for the understanding of the accretion mechanism and the geometry of magnetic fields in these stars.  The observations of the Herbig Ae stars were carried out within a framework of the study of evolution of magnetic fields in stars across the upper main sequence. It has been frequently mentioned in the literature that they are potential progenitors of the magnetic Ap stars (e.g., Stepien \\& Landstreet \\cite{SL02}; Catala \\cite{Ca03}). A detection of magnetic fields in Herbig Ae stars is especially important in view of our recent results that Ap magnetic stars of mass below 3\\,$M_\\odot$ are significantly evolved and concentrated towards the centre of the main-sequence band, and practically no magnetic star of mass below 3\\,$M_\\odot$ can be found close to the zero-age main sequence (Hubrig et~al.\\ \\cite{Hu00}; Hubrig et~al.\\ \\cite{Hu04b}). The  search for magnetic fields and the study of their structure in the pre-main sequence counterparts to the magnetic Ap stars is a crucial step towards understanding the origin of the magnetic fields in stars of intermediate mass.  To properly assess the role of magnetic fields in the star formation process it is important to carry out magnetic field studies of a large sample of Herbig Ae stars and to try to establish the magnetic field strength and the magnetic field geometry by monitoring over several rotation cycles. In T\\,Tauri stars, time series of circular polarization measurements of spectral lines formed throughout the photosphere and those in accretion regions show that the photosphere as a whole lacks globally organized magnetic fields, but accretion regions have highly ordered fields of the order of a few kG (Johns-Krull et~al.\\ \\cite{JK03}). Such contrary behaviour is usually reconciled by a notion that accreting material is loaded onto magnetic field lines at several stellar radii, where the dipolar component dominates higher order components of a more complicated surface geometry. The net longitudinal magnetic field in T\\,Tauri stars with an upper limit of 200\\,G is modulated on rotational time scales, implying significant obliquity between the magnetic and rotational axes (Johns-Krull et~al.\\ \\cite{JK03}). The current magnetohydrodynamic models of Herbig Ae stars should be modified by inclusion of magnetic fields of proper strength and geometry and we should look forward towards future magnetohydrodynamic calculations of processes taking place in the accretion regions and the photosphere of these pre-main sequence stars."
    },
    "0410/astro-ph0410601_arXiv.txt": {
        "abstract": "Complex organic molecules have previously been discovered in solar type protostars, raising the questions of where and how they form in the envelope. Possible formation mechanisms include grain mantle evaporation, interaction of the outflow with its surroundings or the impact of UV/X-rays inside the cavities. In this Letter we present the first interferometric observations of two complex molecules, CH$_3$CN and HCOOCH$_3$, towards the solar type protostar IRAS16293--2422. The images show that the emission originates from two compact regions centered on the two components of the binary system. We discuss how these results favor the grain mantle evaporation scenario and we investigate the implications of these observations for the chemical composition and physical and dynamical state of the two components. ",
        "introduction": "Solar type Class 0 protostars are characterized by being observable only at millimeter to far infrared wavelengths. This is because the central forming stars are surrounded and obscured by massive envelopes, making them the coldest known protostars. Yet, the cold envelopes hide inner warm regions where the grain mantles, built during the pre-collapse phase, evaporate. Following gas phase reactions between the evaporated mantle molecules (also called first generation molecules), organic complex molecules (second generation molecules) form in these regions. While this picture has now been widely accepted for massive protostars (see e.g. Kurtz et al. 2000; van Dishoeck \\& Blake 1998), it is only recently that evidences in support of parts of this scenario have been accumulated for low mass protostars. Actually, so far the existence of inner warm regions with evaporated mantles in solar type protostars has been argued indirectly from the sophisticated analysis of molecular multi-frequency single dish observations (Ceccarelli et al. 2000a,b; Sch\\\"oier et al. 2002; Maret et al. 2002, 2004; Doty et al. 2004). A survey on almost a dozen of Class 0 sources has shown that, typically, the inner warm regions have sizes of a few tens of AUs (Maret et al. 2004), i.e the size of the Solar Nebula. Now that the existence of these regions has been demonstrated, the questions of their nature and chemical composition have become relevant because of the link with the formation of our own Solar System. What is the inventory of organic molecules of such regions? What are their origin and evolution? In particular, could these molecules be incorporated into the planet-forming disks surrounding the protostars?  Based on theoretical arguments, no complex, second generation molecules should have the time to form in solar type protostars. The reason is that the gas crossing time of the warm regions is estimated to be much shorter ($\\lesssim 10^3$ yr), than the predicted chemical formation time ($\\gtrsim 10^4$ yr; Sch\\\"oier et al. 2002). But our theories have somewhere flaws, since complex organic molecules have been detected in the two (out of two) sources where they have been searched for: IRAS16293--2422 (Cazaux et al. 2003) and NGC1333-IRAS4A (Bottinelli et al. 2004). We therefore may even doubt the basic prediction that those molecules originate in the warm inner regions of the envelopes! It is, for example, possible that the process which is responsible for the release in the gas phase of the mantle constituents, is not the thermal evaporation, but the interaction of the outflow with the surroundings, or even of UV/X-rays with the cavities excavated by the outflow (Sch\\\"oier et al. 2002, 2004). Alternatively, the emission could come from disk surfaces which would have been heated by accretion shocks (Sch\\\"oier et al. 2004). The proof that the complex organic molecules detected by Cazaux et al. (2003) and Bottinelli et al. (2004) originate from the inner warm regions (envelope or disk), called hot corinos (Ceccarelli 2004; Bottinelli et al. 2004), was still missing until today.  In this Letter we bring the missing proof and report interferometric observations of two complex, second generation molecules, methyl cyanide and methyl formate, towards the prototype of the hot corinos, IRAS16293--2422 (hereafter IRAS16293). ",
        "conclusions": "The most important result of the presented observations is that the complex molecules observed by Cazaux et al. (2003) originate in two compact regions, whose diameters are about 1.5$''$ (A) or less (B), as shown in Figure \\ref{linemaps} (and Table 1). This goes along with the fact that the images do not show any evidence of emission associated with the molecular outflows seen at larger scales (see below for the discussion on the line profiles). Also, following the remark in the Introduction, the Plateau de Bure images do not reveal any evidence of emission from cavities excavated from the outflows either. Indeed, the cavities are typically $\\sim 20\"$ in size, located $\\sim 30\"$ away from the source along the $^{12}$CO outflow (see e.g. Arce \\& Sargent 2004), which would have been easily detected by the Plateau de Bure. The two regions where the molecular emission comes from, are compact, and while source A is barely resolved in the 1mm images\\footnote{It is likely that the emission from source A originates in the disk around this source. A detailed study of this aspect is postponed to a forthcoming paper.}, source B is unresolved. The measured sizes (1.5$''$ at 160 pc correspond to a radius of about 120 AU) are remarkably consistent with the emission coming from a region where the dust temperature exceeds 100 K in source A (150-200 AU: based on multi-frequency single dish observations: Ceccarelli et al. 2000a,b; Sh\\\"oier et al. 2002), and therefore, where the grain mantles evaporate. Thus, these observations support the basic prediction (from the modeling of the single dish observations; Ceccarelli et al. 2000a,b; Cazaux et al. 2003) that a hot corino with a radius of about 150 AU exists inside the cold envelope of IRAS16293, and that in that region, complex molecules are formed because of grain mantle evaporation.  In addition to that, the Plateau de Bure observations confirm that the two sources A and B are different, as noted by previous authors (Wootten 1989; Mundy et al. 1990, 1992). They differ in line intensities and extent (Fig. \\ref{linemaps}), and this corresponds to a difference in their chemical composition. But, before discussing this point, it is necessary to address the second most striking difference in the two sources: their line profiles (Fig. \\ref{spectra1mm}). Source A has clearly broadened spectra (FWHM $\\sim 8$ \\kms), while source B shows apparently much narrower profiles (FWHM $\\sim 2$ \\kms). Furthermore, the lines of source B seem to peak at \\vlsr=2.7 \\kms, whereas the parent cloud velocity is at \\vlsr=3.9 \\kms\\ (compatible with the spectra of source A, although given the broad profiles it is difficult to precisely determine the \\vlsr\\ of source A). Note that the cloud's \\vlsr\\ (3.9 \\kms) is very nicely measured by the two CN absorption lines in the observed band\\footnote{The absorption originates in the foreground (envelope + cloud) cold gas, which absorbs the photons emitted in the hot corino regions (the CN emission component, being extended, is filtered out by the interferometric observations).} (Figure \\ref{spectra1mm}). As said, no evidence of outflowing gas is seen in the images, and also the broad line profiles of source A are consistent with gas {\\it infalling} towards a $\\sim$1 M$_\\odot$ object (Ceccarelli et al. 2000a,b; Sch\\\"oier et al. 2002). Therefore, both the images and line profiles of source A are fully consistent with the hot corino hypothesis.  The case of source B is less obvious: why does this source have narrower lines and why do they peak at 2.7 \\kms ? But is this true? A very careful look at the B spectra rises doubts. Indeed, all the B lines have a second small peak ---sometimes at the limit of the noise--- on the red-shifted side of the spectrum. This second peak could indeed be part of the line itself, which would  be strongly self-absorbed at \\vlsr=3.9 \\kms. If this is the case, the linewidths of source B would be $\\sim$ 4-6 \\kms\\ (Fig.~\\ref{spectra1mm}), similar to the linewidths measured towards source A. Note that the blue peak is expected to be brighter than the red peak in the case of optically thick lines from infalling gas (Leung \\& Brown 1977; Zhou 1992, 1995; Choi et al. 1995), so it would be consistent with the $\\sim 1$ M$_\\odot$ hot corino hypothesis of source B too. This alternative explanation, optically thick lines in source B, is therefore very appealing and worth exploring in some detail. Using the LTE approximation, the required column densities for the \\ch3cn and \\hcooch3 lines to be optically thick are $N\\sim10^{16}$ cm$^{-2}$ and $N\\sim10^{17}$ cm$^{-2}$ respectively. These values are about one order of magnitude larger than the column density derived in B from the emission lines, assuming that the lines are optically thin, LTE populated and that the emission region fills up the Plateau de Bure synthesized beam (see below). Considering that the three adopted assumptions all underestimate the true column density, it is indeed possible, but not firmly established, that the lines in source B are optically thick. Unfortunately, ``physical'' considerations do not help either to distinguish between the two possible interpretations, optically thin or thick lines in source B. In the first case (B has \\vlsr=2.7 \\kms\\ and FWHM $\\sim 2$ \\kms), source B would be less massive than A and would revolve around it at 1.2 \\kms\\ (multiplied by the inclination of the orbit), at a distance of 800 AU, which is fully consistent with $M_{\\rm A} \\sim 1 M_\\odot$ (unless the orbit is in the sky plane). In the second case (B also has \\vlsr=3.9 \\kms\\ and FWHM $\\sim$ 4-6 \\kms), A and B have comparable masses (similar FWHM), but B is more compact. Future high resolution observations of optically thin lines are hence required to definitely settle the question.  As said, the nature of source B affects the determination of the molecular abundances in this source, and hence, how much the chemical composition of the A and B hot corinos differ. Using the relation between column density and observed continuum flux density, and using a dust opacity $\\kappa_\\nu$ of 0.8 cm$^2$ g$^{-1}$ at 1.3 mm and a dust temperature $T_{\\rm d}$ of 40 K (e.g. Walker et al. 1986), we derive molecular hydrogen column densities from the 1.3~mm continuum emission equal to $N({\\rm H_2,A})=3.5\\times10^{24}$~cm$^{-2}$ in A and $N({\\rm H_2,B})=1.7\\times10^{25}$~cm$^{-2}$ in B (consistent with the values of Mundy et al. (1992)\\footnote{We postpone the detailed analysis of the continuum emission to a forthcoming paper.}). Using these values and the \\ch3cn total column densities derived from the rotational diagram method ($N_{\\rm A}=1.7\\times10^{15}$ cm$^{-2}$ and $N_{\\rm B}=3.9\\times10^{14}$ cm$^{-2}$), we get \\ch3cn abundances of $4.8\\times10^{-10}$ and $2.3\\times10^{-11}$ for A and B respectively, i.e. \\ch3cn is $\\sim$ 20 times more abundant in A than in B. A rotational diagram could not be drawn for the \\hcooch3 transitions as they have similar upper energy levels, but assuming $T_{\\rm rot}\\sim$60 K (Cazaux et al. 2003) we get \\hcooch3 column densities of $N_{\\rm A}=2.6\\times10^{16}$ cm$^{-2}$ and $N_{\\rm B}=8.4\\times10^{15}$ cm$^{-2}$, i.e \\hcooch3 abundances of $7.5\\times10^{-9}$ and $4.9\\times10^{-10}$ for A and B respectively, a factor 15 difference\\footnote{Note that the abundances quoted in Cazaux et al. differ from those derived here because of the different estimate of the H$_2$ column density.}. Note, however, that since B is unresolved in the line emission, and the lines could be optically thick (see above discussion), the molecular abundances quoted for B may be underestimated by about an order of magnitude (note also that the region of molecular emission could be more compact than the continuum emission region). Therefore, resolving the problem of the nature of the observed line emission in source B is crucial, not only to determine the dynamical state of this source, but also to correctly assess the difference in the chemical composition of sources A and B.  In summary, these interferometric observations with the Plateau de Bure show unambiguously that the observed complex organic molecules are emitted from two compact ($\\le 1.5''$) regions, which is consistent with grain mantle evaporation. No indication was found in support of other formation mechanisms for these molecules. Two scenarios have been proposed to explain the observed differences in sources A and B. However, the available data are insufficient to fully understand how and why the two components differ in their dynamical and chemical state, and more interferometric observations (e.g. of optically thin lines) are needed in order to solve the issues raised in this Letter. \\clearpage"
    },
    "0410/astro-ph0410437_arXiv.txt": {
        "abstract": "A new radio spectrometer, CALLISTO, is presented. It is a dual-channel frequency-agile receiver based on commercially available consumer electronics. Its major characteristic is the low price for hardware and software, and the short assembly time, both two or more orders of magnitude below existing spectrometers. The instrument is sensitive at the physical limit and extremely stable. The total bandwidth is 825 MHz, and the width of individual channels is 300 kHz. A total of 1000 measurements can be made per second. The spectrometer is well suited for solar low-frequency radio observations pertinent to space weather research. Five instruments of the type were constructed until now and put into operation at several sites, including Bleien (Zurich) and NRAO (USA). First results in the 45 -- 870 MHz range are presented. Some of them were recorded in a preliminary setup during the time of high solar activity in October and November 2003. ",
        "introduction": "The radio emission of solar flares comes in many varieties, probably reflecting different emission processes and physical origins. The emissions are distinct in their spectrum and temporal behavior. They can best be distinguished in spectrograms, as demonstrated already by Wild \\& McCready (1950). In the meter wavelength range, 5 types of radio emissions are distinguished (review by McLean \\& Labrum, 1985) and at decimeter wavelengths there is a similar number (G\\\"udel \\& Benz, 1988; Isliker \\& Benz, 1994). The radiations are caused coherently by electron beams, shocks, possibly trapped electrons, and high-frequency waves in the plasma. The incoherent gyrosynchrotron radiation is the only really broadband radio emission. As the radio emissions carry information on shocks, energetic electrons and their acceleration, but are still largely unknown, research in flare radio emissions will grow in the near future. In view of space weather implications, meter wave bursts have seen a renewed interest, particularly in combination with observations at other wavelengths.  Spectrometers are necessary to identify the nature of coherent solar flare radio emission. Interferometers do not have sufficient frequency channels to replace spectrometers with high frequency resolution, as some decimeter bursts have bandwidths of less than one percent of the center frequency. Spectrometers will be needed in the future to complement interferometers such as the Frequency Agile Solar Radiotelescope (FASR, Bastian, 2004).  The first radio spectrometers recorded optically the analog signal on film. The development of digitally recording spectrometers in the 1970s (Perrenoud, 1982) soon reached the physical limit of sensitivity for frequency-agile receivers given by the radiometer equation. As the quiet Sun is a strong background, increasing the antenna size helps only for field of views smaller than the full disk. In frequency-agile systems frequency channels are measured sequentially. The sensitivity can only be improved by multi-channel spectrometers, which measure multiple frequency channels simultaneously. Such instruments were developed earlier on, but had a limited number of channels and total bandwidth (Dumas, Caroubalos, \\& Bougeret, 1982). This shortcoming was overcome with the development of acousto-optic receivers (Cole, Stewart, \\& Milne, 1978). However, their limited dynamic range makes acousto-optic spectrometers vulnerable to terrestrial interference. For this reason, spectrometers of acousto-optic design have been operated successfully only at centimeter wavelengths where interference is less severe (Lecacheux et al., 1993; Benz, Messmer, \\& Monstein, 2001). The more robust Fast Fourier Transform (FFT) technology is still in development.  While the development of broadband, interference-resistant spectrometers is still a \"desideratum\", we are left at present with receivers operating according to the frequency-agile principle. There is some development in recent years with regard to broader bandwidth, circular polarization, and better software for operation, calibration and data analysis. Frequency-agile or swept-frequency spectrometers are currently in operation in several places, including China (Xu et al., 2003), Hiraiso (Kondo et al. 1995), Ondrejov (Jiricka et al., 1993), Sao Paolo (Sawant et al., 2001), and Tremsdorf (Mann et al., 1992). ETH Zurich also operates broadband spectrometers since many years. The Ikarus-Phoenix line of digital frequency-agile receivers has seen three decades of steady, but costly developments in hardware and software (Benz et al., 1991).  Today the situation has changed dramatically. The hardware and some software have become commercially available on the consumer electronics market. Thus radio spectrometers can be built cheaply and in great numbers. They are small and transportable. This opens new possibilities, in particular concerning the ever increasing interference. Identical spectrometers could be placed far apart to observe different interference-free windows. Alternatively, a spectrometer can be brought temporarily into a remote region with less interference.  Here we describe the new radio spectrometer CALLISTO (Compact Astronomical Low-frequency, Low-cost Instrument for Spectroscopy in Transportable Observatories) built according to these principles. Its main focus is solar flare radio observation. ",
        "conclusions": "We have presented the new solar radio spectrometer CALLISTO now in operation at the ETH Zurich radio observatory in Bleien (Switzerland). In the 45 -- 870 MHz range CALLISTO surveys a spectral region of renewed interest in view of space weather implications. It is a very cheap instrument suitable for surveying the radio spectrum at low frequencies. Copies can easily be made and distributed to other locations. Its principle aims are low-frequency observations for space weather investigations, complementary observations for interferometers, and routine observations for which it could eventually replace aging spectrometers. Five identical instruments have been realized until today. Two are currently operating near Zurich to complement the existing Phoenix-2 spectrometer below its low-frequency limit and in the meter wave band (150 -- 400 MHz), where the small channel width of CALLISTO is exploited to observe between interferences (http://www.astro.phys.ethz.ch/rapp/). A third instrument is being installed at NRAO in Greenbank (USA) to complement the Green Bank Solar Radio Burst Spectrometer in the range 70 to 870 MHz. Two instruments will be used in temporary set-ups at remote locations free of interference."
    },
    "0410/gr-qc0410044.txt": {
        "abstract": "This paper discusses experimental design for the Laser Astrometric Test Of Relativity (LATOR) mission. LATOR is designed to reach unprecedented accuracy of 1 part in 10$^{8}$ in measuring the curvature of the solar gravitational field as given by the value of the key Eddington post-Newtonian parameter $\\gamma$.  This mission will demonstrate the accuracy needed to measure effects of the next post-Newtonian order ($\\propto G^2$) of light deflection resulting from gravity's intrinsic non-linearity.  LATOR will provide the first precise measurement of the solar quadrupole moment parameter, $J_2$, and will improve determination of a variety of relativistic effects including Lense-Thirring precession.  The mission will benefit from the recent progress in the optical communication technologies -- the immediate and natural step above the standard radio-metric techniques. The key element of LATOR is a geometric redundancy provided by the laser ranging and long-baseline optical interferometry.  We discuss the mission and optical designs, as well as the expected performance of this proposed mission. LATOR will lead to very robust advances in the tests of Fundamental physics: this mission could discover a violation or extension of general relativity, or reveal the presence of an additional long range interaction in the physical law.  There are no analogs to the LATOR experiment; it is unique and is a natural culmination of solar system gravity experiments. ",
        "introduction": "\\label{sec:intro}  After almost ninety years since general relativity was born, Einstein's theory has survived every test. Such a longevity, along with the absence of any adjustable parameters, does not mean that this theory is absolutely correct, but it serves to motivate more accurate tests to determine the level of accuracy at which it is violated. Einstein's general theory of relativity (GR) began with its empirical success in 1915 by explaining the anomalous perihelion precession of Mercury's orbit, using no adjustable theoretical parameters.  Shortly thereafter, Eddington's 1919 observations of star lines-of-sight during a solar eclipse confirmed the doubling of the deflection angles predicted by GR as compared to Newtonian-like and Equivalence Principle arguments.  This conformation made the general theory of relativity an instant success.  From these beginnings, the general theory of relativity has been verified at ever higher accuracy. Thus, microwave ranging to the Viking Lander on Mars yielded accuracy  $\\sim$0.2\\% in the tests of GR \\cite{viking_shapiro1,viking_reasen,viking_shapiro2}. Spacecraft and planetary radar observations reached an accuracy of $\\sim$0.15\\% \\cite{anderson02}.  The astrometric observations of quasars on the solar background performed with Very-Long Baseline Interferometry (VLBI) improved the accuracy of the tests of gravity to $\\sim$0.045\\% \\cite{RoberstonCarter91,Lebach95,Shapiro_SS_etal_2004}. Lunar laser ranging,  a continuing legacy of the Apollo program, provided $\\sim$0.011\\% verification of GR via precision measurements of the lunar orbit \\cite{Ken_LLR68,Ken_LLR91,Ken_LLR30years99,Ken_LLR_PPNprobe03,JimSkipJean96,Williams_etal_2001,LLR_beta_2004}. Finally, the recent experiments with the Cassini spacecraft improved the accuracy of the tests to $\\sim$0.0023\\% \\cite{cassini_ber}. As a result general relativity became the standard theory of gravity when astrometry and spacecraft navigation are concerned.  However, the tensor-scalar theories of gravity, where the usual general relativity tensor field coexists with one or several long-range scalar fields, are believed to be the most promising extension of the theoretical foundation of modern gravitational theory. The superstring, many-dimensional Kaluza-Klein, and inflationary cosmology theories have revived interest in the so-called `dilaton fields', i.e. neutral scalar fields whose background values determine the strength of the coupling constants in the effective four-dimensional theory. The importance of such theories is that they provide a possible route to the quantization of gravity and unification of physical law.  Recent theoretical findings suggest that the present agreement between Einstein's theory and experiment might be naturally compatible with the existence of a scalar contribution to gravity. In particular, Damour and Nordtvedt \\cite{damour_nordtvedt} (see also \\cite{DamourPolyakov94} for non-metric versions of this mechanism and \\cite{DPV02} for the recent summary of a dilaton-runaway scenario) have found that a scalar-tensor theory of gravity may contain a `built-in' cosmological attractor mechanism towards GR.  A possible scenario for cosmological evolution of the scalar field was given in \\cite{Ken_LLR_PPNprobe03,damour_nordtvedt}. Their speculation assumes that the parameter  $\\frac{1}{2}(1-\\gamma)$  was of order of 1 in the early universe, at the time of inflation, and has evolved to be close to, but not exactly equal to, zero at the present time. In fact, the analyzes discussed above not only motivate new searches for very small deviations of relativistic gravity in the solar system, they also predict that such deviations are currently present in the range from 10$^{-5}$ to $\\sim 5\\times 10^{-8}$ of the post-Newtonian effects.  This would require measurement of the effects of the next post-Newtonian order ($\\propto G^2$) of light deflection resulting from gravity's intrinsic non-linearity. An ability to measure the first order light deflection term at the accuracy comparable with the effects of the second order is of the utmost importance for the gravitational theory and is the challenge for the 21st century fundamental physics.  The Eddington parameter $\\gamma$, whose value in general relativity is unity, is perhaps the most fundamental PPN parameter, in that $(1-\\gamma)$ is a measure, for example, of the fractional strength of the scalar gravity interaction in scalar-tensor theories of gravity \\cite{Damour_EFarese96}.  Within perturbation theory for such theories, all other PPN parameters to all relativistic orders collapse to their general relativistic values in proportion to $(1-\\gamma)$. This is why measurement of the first order light deflection effect at the level of accuracy comparable with the second-order contribution would provide the crucial information separating alternative scalar-tensor theories of gravity from GR \\cite{Ken_2PPN_87} and also to probe possible ways for gravity quantization and to test modern theories of cosmological evolution \\cite{damour_nordtvedt,DamourPolyakov94,DPV02}.  The LATOR mission is designed to directly address the issue above with an unprecedented accuracy \\cite{lator_cqg_2004}. The test will be performed in the solar gravity field using optical interferometry between two micro-spacecraft.  Precise measurements of the angular position of the spacecraft will be made using a fiber coupled multi-chanelled optical interferometer on the ISS with a 100 m baseline. The primary objective of the LATOR mission will be to measure the gravitational deflection of light by the solar gravity to accuracy of 0.1 picoradians (prad) ($\\sim0.02 ~\\mu$as), which corresponds to $\\sim$10 picometers (pm) on a 100 m interferometric baseline.  A combination of laser ranging among the spacecraft and direct interferometric measurements will allow LATOR to measure deflection of light in the solar gravity by a factor of more than 3,000 better than had recently been accomplished with the Cassini spacecraft. In particular, this mission will not only measure the key PPN parameter $\\gamma$ to unprecedented levels of accuracy of one part in 10$^8$. As a result, this experiment will measure values of other PPN parameters such as parameter $\\delta$ to 1 part in $10^3$ and discussion thereafter), the solar quadrupole moment parameter $J_2$ to 1 part in 20, and the frame dragging effects on light due to the solar angular momentum to precision of 1 parts in $10^2$.  The LATOR mission technologically is a very sound concept; all technologies that are needed for its success have been already demonstrated as a part of the JPL's Space Interferometry Mission (SIM) development. Technology that has become available in the last several years such as low cost microspacecraft, medium power highly efficient solid state and fiber lasers, and the development of long range interferometric techniques make possible an unprecedented factor of 3,000 improvement in this test of general relativity possible. This mission is unique and is the natural next step in solar system gravity experiments which fully exploits modern technologies.  This paper organized as follows:  Section \\ref{sec:lator_description} provides the overview for the LATOR experiment including the preliminary mission design. In Section \\ref{sec:lator_current} we discuss the current design for the LATOR flight system. In Section \\ref{sec:error_bud} we will discuss the expected performance for the LATOR instrument. Section \\ref{sec:conc} discusses the next steps that will taken in the development of the LATOR mission. ",
        "conclusions": "\\label{sec:conc}  Concluding, we would like to summarize the most significant results of our LATOR mission study. The most natural question is ``Why is LATOR potentially orders of magnitude more sensitive and less expensive?''  First of all, there is a significant advantage in using optical wavelengths as opposed to the microwaves -- the present navigational standard. This is based on the fact that solar plasma effects decrease as $\\lambda^2$ and, in the case of LATOR, we gain a factor of $10^{10}$ reduction in the solar plasma optical path fluctuations by simply moving from $\\lambda=10$~cm to $\\lambda=1 ~\\mu$m. Another LATOR's advantage is its independence of a drag-free technology. In addition, the use of a redundant optical truss offers an excellent alternative to an ultra-precise orbit determination. This feature also makes LATOR insensitive to spacecraft buffeting from solar wind and solar radiation pressure.  Furthermore, the use of existing technologies, laser components and spacecraft make this mission a low cost experiment. Thus, 1~W lasers with sufficient frequency stability and $>10$ years lifetime already developed for optical telecom and also are flight qualified for SIM. Additionally, small optical apertures $\\sim$10-20cm are sufficient and provide this experiment with a high signal-to-noise ratio. There also a significant advantage in using components with no motorized moving parts. This all makes LATOR an excellent candidate for the next flight experiment in fundamental physics. Table \\ref{tab:summ_science} summarizes the science objectives for this mission.  The LATOR mission aims to carry out a test of the curvature of the solar system's gravity  field with an accuracy better than 1 part in 10$^{8}$. In spite of the previous space missions exploiting radio waves for tracking the spacecraft, this mission manifests an actual breakthrough in the relativistic gravity experiments as it allows to take full advantage of the optical techniques that recently became available.  The LATOR experiment has a number of advantages over techniques that use radio waves to measure gravitational light deflection. The optical technologies allows low bandwidth telecommunications with the LATOR spacecraft. The use of the monochromatic light enables the observation of the spacecraft almost at the limb of the Sun. The use of narrowband filters, coronagraph optics and heterodyne detection will suppress background light to a level where the solar background is no longer the dominant noise source. The short wavelength allows much more efficient links with smaller apertures, thereby eliminating the need for a deployable antenna. Finally, the use of the ISS enables the test above the Earth's atmosphere -- the major source of astrometric noise for any ground based interferometer. This fact justifies LATOR as a space mission.  The LATOR mission will utilize several technology solutions that recently became available. In particular, signal acquisition on the solar background will be done with a full-aperture narrow band-pass filer and coronagraph. The issue of the extended structure vibrations of the will be addressed by using $\\mu$-g accelerometers. (The use of the accelerometers was first devised for SIM, but at the end their utilization is not needed. The Keck Interferometer uses accelerometers extensively.) Finally, the problem of monochromatic fringe ambiguity that complicated the design of the previous version of the experiment \\cite{yu94} and led to the use of variable baselines lengths -- is not an issue for LATOR. This is because the orbital motion of the ISS provides variable baseline projection that eliminates this problem for LATOR.  The LATOR experiment technologically is a very sound concept; all technologies that are needed for its success have been already demonstrated as a part of the JPL's interferometry program.  The LATOR experiment does not need a drag-free system, but uses a geometric redundant optical truss to achieve a very precise determination of the interplanetary distances between the two micro-spacecraft and a beacon station on the ISS. The interest of the approach is to take advantage of the existing space-qualified optical technologies leading to an outstanding performance in a reasonable mission development time.  The  availability of the ISS makes this mission concept realizable in the very near future; the current mission concept calls for a launch as early as in 2011 at a cost of a NASA MIDEX mission.  This mission may become a 21st century version of Michelson-Morley experiment in the search for a  cosmologically evolved scalar field in the solar system. As such, LATOR will lead to very robust advances in the tests of fundamental physics: it could discover a violation or extension of GR, or reveal the presence of an additional long range interaction in the physical law.  There are no analogs to the LATOR experiment; it is unique and is a natural culmination of solar system gravity experiments.  \\vfill %**************************************"
    },
    "0410/astro-ph0410711_arXiv.txt": {
        "abstract": "Galaxy-galaxy lensing is a powerful probe of the relation between galaxies and dark matter halos, but its theoretical interpretation requires a careful modeling of various contributions, such as the contribution from central and satellite galaxies. For this purpose, a phenomenological approach based on the halo model has been developed, allowing for fast exploration of the parameter space of models. In this paper, we investigate the ability of the halo model to extract information from the g-g weak lensing signal by comparing it to high-resolution dissipationless simulations that resolve subhalos. We find that the halo model reliably determines parameters such as the host halo mass of central galaxies, the fraction of galaxies that are satellites, and their radial distribution inside larger halos. If there is a significant scatter present in the central galaxy host halo mass distribution, then the mean and median mass of that distribution can differ significantly from one another, and the halo model mass determination lies between the two. This result suggests that when analyzing the data, galaxy subsamples with a narrow central galaxy halo mass distribution, such as those based on stellar mass, should be chosen for a simpler interpretation of the results. ",
        "introduction": "Understanding the connection between the spatial distribution of galaxies and dark matter (DM) is one of the most important problems in modern cosmology.  From the perspective of fundamental physics, this connection needs to be understood if we wish to use galaxy clustering to determine dark matter correlations, a potentially powerful discriminator between cosmological models. This connection can also be a powerful test of the tenets of the standard cosmological paradigm, such as the collisionless nature of cold dark matter.  From the astrophysics perspective, this connection is an essential ingredient in the physics of galaxy formation. Current models based on pure N-body simulations, semi-analytic models or hydrodynamic simulations can accomodate many observational aspects of this connection, but several ingredients of these models remain uncertain and need to be inserted ad-hoc, so it is not clear how much of the success is a result of the allowed freedom within these models. It is important to look for new ways to test this connection, to confirm and improve the existing models of structure formation, and to distinguish astrophysical effects from cosmological ones.  One of the probes of the galaxy-DM connection that recently became available is weak lensing around galaxies, or galaxy-galaxy (hereinafter g-g) lensing \\citep{1984ApJ...281L..59T,1996ApJ...466..623B,1998ApJ...503..531H,2000AJ....120.1198F,2001ApJ...551..643S,2001astro.ph..8013M,2003MNRAS.340..609H,2004AJ....127.2544S,2004ApJ...606...67H,2005PhRvD..71d3511S}. Gravitational lensing induces tangential shear distortions of background galaxies around foreground galaxies, allowing direct measurement of the galaxy-dark matter correlation function around galaxies. The individual distortions are small (of order 0.1\\%), but by averaging over all foreground galaxies within a given subsample, we obtain high signal to noise in the shear as a function of angular separation from the galaxy. If we know the galaxy redshifts, the shear signal can be related to the projected mass density as a function of proper distance from the galaxy. This allows us to determine statistically the dark matter distribution around any given galaxy sample.  In recent years, the progress on the observational side of g-g lensing has been remarkable. In the latest Sloan Digital Sky Survey (SDSS) analyses \\citep{2004AJ....127.2544S,2005PhRvD..71d3511S}, 20-30 sigma detections of the signal as a function of physical separation have been obtained. Similarly high S/N detections have also been observed as a function of angular separation with other surveys \\citep{2004ApJ...606...67H}. This increased statistical power has been accompanied by a more careful investigation of systematic errors, such as calibration biases and intrinsic correlations, which for the SDSS are currently around 10\\% and therefore already dominate the error budget \\citep{2004MNRAS.353..529H}. As the quality and quantity of the data and its analysis improves, the reliability of theoretical interpretation must be improved as well. The goal of the present paper is to compare the various theoretical analyses among themselves, and to discuss how they should be applied to the data.  Theoretical analysis of g-g lensing falls into two categories. The first approach is a direct comparison of simulations to the data \\citep{2001MNRAS.321..439G,2003MNRAS.339..387Y,2004ApJ...601....1W,2004ApJ...614..533T}. This approach is direct, but rather expensive, since the process of galaxy formation is not understood sufficiently well to result in unique predictions for a given cosmological model, nor is the cosmological model  itself determined yet. Moreover, current simulations still suffer from a limited dynamical range, in the sense that they require a high mass and force resolution to resolve individual galaxies and their associated dark matter halos, while at the same time they must also have sufficiently large volume to simulate a representative region of the universe. Several simulations of varying box size are needed to cover the whole observational range in luminosity and scale. As a result, many different simulations would be needed to explore the whole range of parameter space; at present, only a handful of simulations have been used for this application.  The second approach is to use the halo model \\citep[for a review, see][]{2002PhR...372....1C} to describe the connection between galaxies and dark matter \\citep[][hereinafter GS02]{2002MNRAS.335..311G}. This approach is more phenomenological, and for a given galaxy class leads to determination of quantities such as the virial mass distribution and the fraction of these galaxies that are satellites, which are useful quantities for constraining the galaxy formation models and cosmological models. While the information extracted from the data using the halo model may be all that is needed to quantify the galaxy-dark matter connection, halo models are reliable only to the extent that they are able to reproduce the simulations, so they must be tested and calibrated on the simulations before applications to the real data are reliable.  In this paper, we attempt to understand how well the phenomenological halo model developed in~\\cite{2000MNRAS.318..203S} and studied in GS02 can reproduce the simulations. We are interested in its ability to reproduce the g-g lensing signal, which in turn can be used to determine the halo mass probability distribution, the radial profile of satellite galaxies inside larger halos, and other quantities of interest. ",
        "conclusions": "The goal of this paper is to present a detailed comparison of halo model and simulation predictions of galaxy-mass correlations as observed in weak lensing.  In its simplest form, and for a given cosmological model, the halo model has 3 free parameters, which depend on the physics of galaxy formation.  These are the typical halo mass of central galaxies (where ``typical'' is between the median and mean halo mass), the fraction of galaxies that are satellites, and the radial distribution of satellites inside the larger halo.  The halo model provides a fairly accurate description of the g-g lensing signal from simulations, in the sense that it can determine the virial mass distribution of central and non-central galaxies and their relative ratios, as well as the radial distribution of both galaxies and dark matter.  One must be careful when interpreting the halo masses from NFW fits if the halo mass distribution is broad as is expected for luminous galaxies ($L>$ few $L_{\\ast}$). In this case, g-g lensing is determining something between the mean and the median mass and may not recover the correct radial distribution of the satellite galaxies in their host halos.  If one is interested in the mean halo mass, as in the case of application to bias studies \\citep{2005PhRvD..71d3511S}, then an upward correction needs to be applied; this correction is small at the faint end but can be quite significant at the bright end, especially if there is a significant scatter in the luminosity-halo mass relation.  There are additional parameters that can be added to the description, such as the generalization of the assumed halo mass probability distribution, but they do not improve the fit and are strongly correlated with these 3 parameters.  Galaxy-galaxy lensing also allows a determination of the halo mass profile of both central galaxies and satellites. In the latter case, one can constrain the amount of satellite mass stripping inside larger halos, which determined how much dark matter remains attached to the satellites.  To investigate the radial profiles of dark matter it is best if one identifies samples that consist predominantly of central or satellite galaxies, respectively.  Galaxy-galaxy lensing is also quite sensitive to the dark matter concentration parameter $c_{\\rm dm}$, which (in standard cold dark matter models) is not a free parameter but is fixed for any given cosmological model. It varies modestly as a function of cosmology, so in principle one could use this feature to determine the cosmological model. For example, for a typical $L_*$ galaxy with mass around $(1-2)\\times 10^{12}h^{-1}M_{\\sun}$, we have $c_{\\rm dm}=12$ in a $\\sigma_8=0.9$, $\\Omega_m=0.3$ cosmology, which drops to $c_{\\rm dm}=10$ if $\\sigma_8=0.7$ is used instead. This change is small and we find that it is degenerate with the amount of scatter, which also leads to a lower value of concentration parameter. With current samples, one cannot yet reliably distinguish between these models, but one can test the overall consistency. The accuracy and reliability with which the concentration can be extracted from the data may be improved if galaxies in underdense regions, for which the satellite fraction is lower, are selected. In this case, one can use the signal to a larger distance since it is not swamped by the noncentral component. Concentration can also be determined more accurately with deeper surveys, which can probe dark matter halo at smaller separations \\citep{2004ApJ...606...67H}, although modeling angular projection effects in the absence of reliable redshift information for lens galaxies may be difficult.  We find that the data are rather insensitive to the truncation radius of the satellites inside the larger halo as long as it is not very small. We can, however, test the extreme possibility that satellites have no dark matter attached to them. This is unlikely to be the case for all satellites even though it may be valid for some fraction of those near the center. For example, \\cite{2004MNRAS.352L...1G} argue that up to 40\\% of galaxies in clusters may not have associated DM halos, a conjecture which may be testable with the future weak lensing measurements (see Fig.~\\ref{F:trunc}).  Additional information about the galaxy-dark matter halo connection may be more challenging to extract from the data. For example, we have assumed that the number of satellites inside the halo scales linearly with halo mass. If this assumption is dropped and a more general power law relation is allowed, a degeneracy is developed between the satellite fraction and the power law slope, but the overall fraction of galaxies in $(10^{13}-10^{14})h^{-1}M_{\\sun}$ halos is preserved.  We note here that the power law slope and other parameters may be constrained by other observations such as the galaxy auto-correlation function analysis, velocity information in redshift surveys and direct counting of galaxies in groups and clusters. In fact, a halo model similar to the one used here was applied to the galaxy auto-correlation analysis \\citep{2004astro.ph..8569Z}. Combining and comparing the two data sets should provide important consistency checks and will further constrain the galaxy-dark matter connection."
    },
    "0410/astro-ph0410527_arXiv.txt": {
        "abstract": "We use highly spectroscopically complete deep and wide-area {\\em Chandra\\/} surveys to determine the cosmic evolution of hard X-ray--selected AGNs. For the deep fields, we supplement the spectroscopic redshifts with photometric redshifts to assess where the unidentified sources are likely to lie. We find that the median redshifts are fairly constant with X-ray flux at $z\\sim 1$. We classify the optical spectra and measure the FWHM line widths. Most of the broad-line AGNs show essentially no visible absorption in X-rays, while the sources without broad lines (${\\rm FWHM} < 2000$~km~s$^{-1}$; ``optically-narrow'' AGNs) show a wide range of absorbing column densities. We determine hard X-ray luminosity functions for all spectral types with $L_X\\ge 10^{42}$~ergs~s$^{-1}$ and for broad-line AGNs alone. At $z<1.2$, both are well described by pure luminosity evolution, with $L_\\ast$ evolving as $(1+z)^{3.2\\pm 0.8}$ for all spectral types and as $(1+z)^{3.0\\pm 1.0}$ for broad-line AGNs alone. Thus, all AGNs drop in luminosity by almost an order of magnitude over this redshift range. We show that this observed drop is due to AGN downsizing rather than to an evolution in the accretion rates onto the supermassive black holes.  We directly compare our broad-line AGN hard X-ray luminosity functions with the optical QSO luminosity functions and find that at the bright end they agree extremely well at all redshifts. However, the optical QSO luminosity functions do not probe faint enough to see the downturn in the broad-line AGN hard X-ray luminosity functions and even appear to be missing some sources at the lowest luminosities they probe.  We find that broad-line AGNs dominate the number densities at the higher X-ray luminosities, while optically-narrow AGNs dominate at the lower X-ray luminosities. We rule out galaxy dilution as a partial explanation for this effect by measuring the nuclear UV/optical properties of the {\\em Chandra\\/} sources using the {\\em HST\\/} ACS GOODS-North data. The UV/optical nuclei of the optically-narrow AGNs are much weaker than expected if the optically-narrow AGNs were similar to the broad-line AGNs. We therefore postulate the need for a luminosity dependent unified model. An alternative possibility is that the broad-line AGNs and the optically-narrow AGNs are intrinsically different source populations. We cover both interpretations by constructing composite spectral energy distributions---including long-wavelength data from the mid-infrared to the submillimeter---by spectral type and by X-ray luminosity. We use these spectral energy distributions to infer the bolometric corrections (from hard X-ray luminosities to bolometric luminosities) needed to map the accretion history.  We determine the accreted supermassive black hole mass density for all spectral types and for broad-line AGNs alone using the observed evolution of the hard X-ray energy density production rate and our inferred bolometric corrections. We find that only about one-half to one-quarter of the supermassive black hole mass density was fabricated in broad-line AGNs. Using either recent optical QSO luminosity function determinations or our broad-line AGN hard X-ray luminosity function determinations, we measure an accreted supermassive black hole mass density that is a factor of almost two lower than that measured by previous work, assuming $\\epsilon=0.1$. This leaves room for the obscured accretion when compared with the local supermassive black hole mass density. In fact, we find reasonable agreement between the accreted supermassive black hole mass density from all spectral types and the local supermassive black hole mass density, assuming $\\epsilon\\approx 0.1-0.2$. However, there is very little room for further obscured sources or for any low efficiency accretion periods. ",
        "introduction": "\\label{secintro}  The determination of the time-history of accretion is crucial to our understanding of how supermassive black holes form and evolve. However, much of the accretion power in the universe is absorbed (e.g., Almaini, Lawrence, \\& Boyle 1999), making it difficult to measure at optical wavelengths. The {\\em Chandra\\/} (Weisskopf et al.\\ 2002) and {\\em XMM-Newton\\/} (Jansen et al.\\ 2001) {\\em X-ray Observatories\\/} have revolutionized distant active galactic nucleus (AGN) studies by making it possible to map the history of the AGN population using hard ($2-8$~keV) X-ray surveys. Hard X-rays can directly probe AGN activity, are uncontaminated by star formation processes at the X-ray luminosities of interest, and detect all but the most absorbed sources. Thus, hard X-ray surveys provide as complete and unbiased a sample of AGNs as is presently possible (e.g., Mushotzky 2004), though they will still miss Compton-thick sources.  Not surprisingly, hard X-ray luminosity functions constructed from {\\em Chandra\\/} and {\\em XMM-Newton\\/} samples have revealed that previous optically-selected and soft X-ray--selected AGN samples substantially undercount the AGN population, at least at low and intermediate X-ray luminosities ($L_{2-8~{\\rm keV}}<10^{44}$~ergs~s$^{-1}$; Cowie et al.\\ 2003; Hasinger 2003; Steffen et al.\\ 2003; Fiore et al.\\ 2003; Ueda et al.\\ 2003). In fact, optically-selected broad-line AGNs comprise only about a third (e.g., Barger et al.\\ 2003b) of the X-ray background (Giacconi et al.\\ 1962), and many X-ray sources show no signs of AGN activity at all in their optical spectra (e.g., Barger et al.\\ 2001b; Tozzi et al.\\ 2001; Hornschemeier et al.\\ 2001). Theoretical models of supermassive black hole formation (e.g., Haehnelt, Natarajan, \\& Rees 1998; Kauffmann \\& Haehnelt 2000, 2002) have historically relied on comparisons with the optical quasar luminosity function, so these models now need to be reworked.  In addition, the hard X-ray luminosity functions have revealed that broad-line AGNs dominate the number densities at the higher X-ray luminosities, while non--broad-line AGNs dominate at the lower X-ray luminosities (Steffen et al.\\ 2003; we hereafter refer to this as the ``Steffen effect''). Although we do not yet have a physical explanation for the Steffen effect, there are two simple possibilities to consider. One possibility is that the simple unified model for AGNs, where the differences between broad-line AGNs and non--broad-line AGNs are only due to orientation effects, needs to be modified to include an X-ray luminosity dependent covering factor. The second, more speculative possibility is that the Steffen effect is a consequence of the broad-line AGNs being intrinsically different than the non--broad-line AGNs.  We are now in a position to do a very thorough study of the nature and evolution of hard X-ray--selected AGNs. The primary aims of this paper are to explore the origin of the Steffen effect and to determine the bolometric luminosities of all AGNs in order to understand the energy release history of supermassive black hole accretion. The recent advances that make this comprehensive study possible are the high-resolution, multicolor observations that have been obtained with the ACS camera on the {\\em Hubble Space Telescope (HST)\\/} of fields with deep {\\em Chandra\\/} X-ray data, and the extensive, high-quality optical spectroscopic follow-up observations that have been made of X-ray sources detected in both deep and wide-area {\\em Chandra\\/} surveys.  The structure of the paper is as follows. In \\S\\ref{secsamp}, we describe the X-ray samples used in our analysis and the completeness of the spectroscopic and photometric redshift identifications. We then determine the X-ray luminosities and median redshifts versus X-ray flux. In \\S\\ref{secclass}, we spectrally classify the optical counterparts to the X-ray sources and investigate the dependence of the optical spectral types on X-ray obscuration. In \\S\\ref{secnuc}, we use multicolor {\\em HST\\/} ACS Great Observatories Origins Deep Survey (GOODS-North; Giavalisco et al.\\ 2004) observations of the CDF-N to separate the nuclear component of each source from the host galaxy light, and we then compare the (nuclear -- galaxy) colors with the optical spectral types.  In \\S\\ref{sechxlf}--\\S\\ref{sechizhxlf}, we construct up-to-date low and high-redshift hard X-ray luminosity functions (all spectral types and broad-line AGNs alone) for our highly spectroscopically complete samples, do maximum likelihood fits over the redshift range $z=0-1.2$, and examine the evolution of the hard X-ray luminosity functions. In \\S\\ref{seclf}, we directly compare the broad-line AGN hard X-ray luminosity functions with the optical QSO luminosity functions. In \\S\\ref{secenergy}, we determine the evolution of the rest-frame hard X-ray comoving energy density production rate for all spectral types together as well as separated by spectral type. In \\S\\ref{secsmbh}, we estimate black hole masses for a small sample of broad-line AGNs using the measured MgII 2800~\\AA\\ line widths and nuclear optical luminosities. We then explore ``mass starvation'' versus ``AGN downsizing'' as explanations for the observed rapid decline in the energy density production rates between $z=1$ and $z=0$.  We show in \\S\\ref{secfailure} that the simple unified model cannot explain the Steffen effect. In \\S\\ref{seclum}, we postulate that the simplest interpretation of the Steffen effect is a luminosity dependent unified model, although another interpretation might be intrinsic differences in the source populations. In \\S\\ref{sechow}, we discuss how the bolometric corrections to go from X-ray luminosities to bolometric luminosities can be determined. To infer what these corrections are, we need long-wavelength data. In \\S\\ref{seclongwave}, we use mid-infrared (MIR) and far-infrared (FIR)/submillimeter data obtained with the ISOPHOT and ISOCAM instruments on the {\\em Infrared Space Observatory (ISO)\\/} and the SCUBA bolometer array on the James Clerk Maxwell Telescope (JCMT) to observe directly any enhancements at these wavelengths due to absorption and reradiation by gas and dust. We note that the observational situation in the MIR/FIR may be expected to improve with the {\\em Spitzer Space Telescope\\/}. Since we do not know for sure what the origin of the Steffen effect is, in \\S\\ref{secspectype} and \\S\\ref{sechxlum}, we cover both possibilities (intrinsic differences in the source populations or a luminosity dependent unified model) by constructing composite spectral energy distributions of the sources first by optical spectral type and then by X-ray luminosity. In \\S\\ref{secbolcorrtype} and \\S\\ref{secbolcorrlum}, we infer the bolometric corrections by spectral type and by X-ray luminosity. In \\S\\ref{secacc}, we use our bolometric corrections and the rest-frame hard X-ray comoving energy density production rate to determine the accretion history of the universe. We summarize our results in \\S\\ref{secsummary}.  We assume $\\Omega_M=1/3$, $\\Omega_\\Lambda=2/3$, and $H_o=65$~km~s$^{-1}$~Mpc$^{-1}$. All magnitudes are in the AB magnitude system. ",
        "conclusions": "\\label{secsummary}  In this paper, we presented a thorough analysis of the nature and evolution of hard X-ray--selected AGNs using both deep and wide-area {\\em Chandra\\/} surveys with nearly complete optical spectroscopic follow-up observations. Our results are as follows.  $\\bullet$ We determined median redshifts and $1\\sigma$ median redshift ranges with X-ray flux for the spectroscopically and spectroscopically plus photometrically identified X-ray samples. We found that the median redshifts are fairly constant with X-ray flux at $z\\sim 1$ and that the lower X-ray flux sources, where the surveys are more spectroscopically incomplete, are dominated by sources with hard X-ray luminosities below $10^{44}$~ergs~s$^{-1}$ and redshifts near one.  $\\bullet$ We spectrally classified the optical counterparts to the X-ray sources and measured the FWHM line widths of the spectra. We found that the X-ray--selected broad-line AGNs recover essentially all of the optically-selected type 1 AGNs and that the only distinction that can be made from the X-ray colors is between the broad-line AGNs and the sources that do not show broad lines (${\\rm FWHM} < 2000$~km~s$^{-1}$; the ``optically-narrow'' AGNs). Most of the broad-line AGNs show essentially no visible absorption in X-rays, while the optically-narrow AGNs show a wide range of absorbing column densities. However, these absorbing column densities show little dependence on the optical spectral type or on the line widths of the optically-narrow AGNs.  $\\bullet$ We constructed up-to-date low and high-redshift hard X-ray luminosity functions for all spectral types and for broad-line AGNs alone using our highly spectroscopically complete observed samples. We did maximum likelihood fits over the redshift range $z=0-1.2$. We found that all of the hard X-ray luminosity functions are consistent with pure luminosity evolution, with $L_\\ast=(1+z)^{3.2\\pm 0.8}$ for all spectral types and $L_\\ast=(1+z)^{3.0\\pm 1.0}$ for broad-line AGNs. Thus, all AGNs drop in luminosity by almost an order of magnitude over this redshift range.  $\\bullet$ We directly compared our broad-line AGN hard X-ray luminosity functions with the optical QSO luminosity functions and found that at the bright end, the luminosity functions agreed extremely well at all redshifts. However, we found that the optical QSO luminosity functions do not probe faint enough to see the downturn in the broad-line AGN hard X-ray luminosity functions and that they may even be missing some sources at the very lowest luminosities to which they probe.  $\\bullet$ We determined the evolution with redshift of the hard X-ray energy density production rate and found that at $z<1$, it bears a striking resemblance to the overall redshift evolution of the star formation rate density. We also determined the evolution according to spectral type, applying incompleteness corrections to the production rates of the optically-narrow AGNs. We found that at $z<1.5$, most of the hard X-ray energy density production is due to optically-narrow AGNs, while at higher redshifts, the broad-line AGNs dominate the production.  $\\bullet$ We estimated black hole masses for the broad-line AGNs using our measured MgII 2800~\\AA\\ line widths and nuclear optical luminosities of sources in the {\\em HST\\/} ACS GOODS-North region of the CDF-N. For the resulting small sample, we found a surprisingly tight correlation between hard X-ray luminosity and black hole mass at $z=1$, with the black holes accreting at a fairly constant rate of 16\\% of the Eddington luminosity.  $\\bullet$ From the X-ray sources in the ACS GOODS-North region of the CDF-N, we found that the most luminous X-ray sources lie in the most optically-luminous host galaxies, as would be expected from the bulge luminosity--black hole mass relation. We also found that the host galaxy optical luminosities do not change much with redshift for a given interval in hard X-ray luminosity, but the most optically-luminous host galaxies that existed at high redshifts have completely disappered at low redshifts. We therefore concluded that the observed drop in the hard X-ray luminosities is due to AGN downsizing rather than to an evolution in the accretion rates onto the supermassive black holes. That is, the highest mass supermassive black holes in the most optically-luminous host galaxies are switching off between $z=1$ and the present time, leaving only the lower mass supermassive black holes in the less optically-luminous galaxies active.  $\\bullet$ A primary aim of the paper was to explore the origin of the Steffen effect, which is the observation that broad-line AGNs dominate the number densities at the higher X-ray luminosities, while non--broad-line AGNs dominate at the lower X-ray luminosities. We ruled out galaxy dilution as being a partial explanation of the Steffen effect by seeing how the nuclear UV/optical properties of the X-ray sources in the CDF-N (measured from the ACS GOODS-North data) relate to the spectral characteristics of the sources. We found that the UV/optical nuclei of the optically-narrow AGNs are much weaker than expected if the optically-narrow AGNs were similar to the broad-line AGNs. Since the simple unified model cannot explain the Steffen effect, we postulated that the simplest interpretation of the Steffen effect is a luminosity dependent unified model. Another possible interpretation is that the broad-line AGNs and the optically-narrow AGNs are intrinsically different. Regardless of which interpretation is correct, a large fraction of the low-redshift, hard X-ray energy density production is in optically-narrow AGNs; thus, bolometric corrections for these sources are needed if we are to map the accretion history with redshift.  $\\bullet$ To infer the bolometric corrections, we constructed composite spectral energy distributions (including long-wavelength data from the mid-infrared to the submillimeter) as a function of both spectral type and X-ray luminosity (to cover both interpretations of the Steffen effect). In contrast to the broad-line AGNs, we found that the optically-narrow AGNs are, on average, highly suppressed in the frequency range between the UV and soft X-rays and have a strong far-infrared component. We inferred a bolometric correction of about 85 for these sources. Likewise, we found that the low X-ray luminosity sources have much fainter UV/optical nuclei relative to their $2-8$~keV fluxes than do high X-ray luminosity sources. The far-infrared luminosities of these low X-ray luminosity sources are also consistent with a bolometric correction of 85.  $\\bullet$ Using our bolometric corrections and our parameterizations of the evolution of the rest-frame hard X-ray energy density production rates, we inferred the accreted supermassive black hole mass densities for all spectral types and for broad-line AGNs alone. We found that only about one-half to one-quarter of the supermassive black hole mass density was fabricated in broad-line AGNs.  $\\bullet$ We compared our accreted supermassive black hole mass density for broad-line AGNs alone with the optical QSO determination made by Yu \\& Tremaine (2002) using the Boyle et al.\\ (2000) QSO luminosity function, assuming $\\epsilon=0.1$. Our value is 1.75 times smaller than their value, but this difference is very sensitive to their extrapolation of the optical QSO luminosity function outside the optically observed luminosity range. The new Croom et al.\\ (2004) QSO luminosity function has a shallower faint-end slope than Boyle et al.\\ (2000). We recalculated the accreted supermassive black hole mass densities for optical QSOs using the Croom et al.\\ fit, first by restricting to the measured absolute magnitude range, and second by extrapolating the fit to higher and lower luminosities, and we found that the two values bracketed our hard X-ray luminosity function determination. Thus, both the new optical QSO luminosity function and the hard X-ray luminosity function determinations of the supermassive black hole mass densities are a factor of almost two lower than the Yu \\& Tremaine (2002) determination. This leaves room for obscured accretion by the optically-narrow AGNs in a comparison with the local supermassive black hole mass density.  $\\bullet$ For $\\epsilon\\approx 0.1-0.2$, we found reasonable agreement between our hard X-ray luminosity function supermassive black hole mass density for all spectral types and the local supermassive black hole mass density; however, there is little room for further obscured sources or for any low efficiency accretion periods."
    },
    "0410/astro-ph0410657_arXiv.txt": {
        "abstract": "We report the detection of the linear radical HC$_4$N in the C-rich envelope of IRC+10216. After HCCN \\citep{Gue91},  HC$_4$N is the second member of the allenic chain family HC$_{2n}$N observed in space. The column density of HC$_4$N is found to be 1.5 10$^{12}$ cm$^{-2}$. The abundance ratio HC$_2$N/HC$_4$N is 9, a factor of two larger than the decrement observed for the cyanopolyynes (HC$_{2n+1}$N/HC$_{2n+3}$N). Linear HC$_4$N has a $^3\\Sigma$ electronic ground state and is one of the 3 low-energy isomeric forms of this molecule. We have searched for the bent and ringed HC$_4$N isomers, but could only derive an upper limit to their column densities $\\le$ 3\\, 10$^{12}$  cm$^{-2}$ (at 3-$\\sigma$). ",
        "introduction": "The detection of long linear carbon-chain molecules in interstellar and circumstellar clouds has come as a surprise. Ab initio calculation show that long linear species are usually more energetic than their ring or ring-chain isomers and are observed to be less stable in the earthbound environment. Yet many linear chains, such as the cyanopolyynes HC$_{2n+1}$N, are found to be widespread and abundant in the cold circumstellar and interstellar clouds. This abundance of long chains and the scarcity of rings in the cold UV-shielded clouds yield keys to the formation of large molecules in interstellar space and, may be, can help us to understand to the origin of the diffuse interstellar bands \\citep{Tul98,Mai04}.  Carbon chain molecules can be classified into several families, depending on the structure of their linear backbone (acetylenic backbone with alternating single and triple carbon-carbon bonds, or cumulenic backbone with double bonds), on their electronic ground state (open shell or closed shell) and on their end groups (H, CN, CH$_2$ or CH$_3$). The two most widespread families are the polyynes (HC$_{2n+1}$N and CH$_3$C$_{2n+1}$H) and the polyacetylenic radicals (C$_{n}$H). The most abundant is presumably the family of polyacetylenes (HC$_{2n}$H), which, unfortunately are non-polar, hence not detectable at radio wavelengths. Note, however, that these species have been detected in the mid infrared with the Infrared Space Observatory \\citep{Cer01a,Cer01b}.  It has been proposed that the polyynes and carbon-chain radicals form directly in the gas phase through reactions of CCH with polyynes, polyacetylenes and/or polyacetylenic ions (e.g. HC$_{2n+1}$N + CCH  $\\rightarrow$ HC$_{2n+3}$N + H) \\citep{Her89,Che93,Mil00}. Such reactions tend to insert two triply bonded carbon atoms into the carbon backbone, uncoupling the formation of chains with an odd number of C-atoms from those with an even number, which could explain the alternance between high and low abundances observed as the length of the C backbone increases. For example, in the molecular shell around IRC+10216 and in the dark cloud TMC1, the two main astronomical sources of carbon chain molecules, the abundance of the C$_n$H radicals with an {\\it even} number of C-atoms is $\\simeq30$ times larger than that of radicals of similar size with an {\\it odd} number of carbon atoms. In contrast, the abundance decrement between successive species within the even (or odd) number of C-atoms families, C$_n$H/C$_{n+2}$H, is found to be only 4-6 \\citep{Cer87, Gue91}.  The same seems to apply to the chains terminated by CN: the abundance decrement between cyanopolyynes, HC$_{2n+1}$N/HC$_{2n+3}$N, is between 3 and 5 (for n=1-4), whereas HCCN, the first member of the `even' chain family, is found to be $\\simeq 200$ times less abundant than HC$_3$N \\citep{Gue91}. So far, this similarity between polyynes and carbon chains, could not be further explored as no other member of the HC$_{2n}$N was observed in space. In this Letter, we report the detection in IRC+10216 of the next member of this family, HC$_4$N, and show that the HCCN/HC$_4$N abundance ratio is 20 times smaller than the HC$_3$N/HCCN ratio. ",
        "conclusions": "Figure 2 shows the spectra covering the 4 triplets, which correspond to 4 successive rotational transitions of HC$_4$N (from N=$18 \\rightarrow 17$ to $21\\rightarrow 20$). The spectrum centered at 82.9 GHz, which was observed while searching for SiCN \\citep{Gue00}, SiNC \\citep{Gue04} and HCCN, is shown twice in this figure: full scale, i.e. from 82650 MHz to 83150 MHz (Fig. 2a) and half-scale, around the HC$_4$N N=18$\\rightarrow$17 triplet (Fig. 2b). The r.m.s. noise in this 1 MHz resolution spectrum is only 0.5 mK, which makes it one of the the most sensitive spectra ever obtained with the 30-m telescope.  Forty spectral lines are detected in Fig. 2a. Most are weak and partly blended. The blended lines are easily resolved, thanks to their characteristic cusps, sharp edges and constant width in velocity (this shape results from the uniform expansion of the outer circumstellar envelope). We can identify almost 60\\% of these lines with the help of the molecular line catalog maintained by one of us (JC) and described in \\citet{Cer00}. The catalog contains the mm transitions of some 1200 different molecular species. Of the remaining 40\\% unidentified lines, at least half could be tentatively assigned to some heavy species (see below). The carriers of the unambiguously identified lines are mostly $^{12}$C and $^{13}$C isotopomers of long carbon chain molecules (C$_4$H, C$_7$H, C$_8$H, H$_2$C$_3$, H$_2$C$_4$, and HC$_5$N), as well as silicon compounds, such as SiC$_4$, SiCN, and SiNC. For the first time, we detect a doubly-substituted $^{13}$C isotopomer of cyanoacetylene: H$^{13}$C$^{13}$CCN (see Fig2e, 96623 MHz; several lines of the three doubly-substituted $^{13}$C isotopomers of HC$_3$N have been detected in our 3mm line survey of IRC+10216).  The number of unidentified lines at the 3 mK level detected in the 4 frequency bands is $\\simeq$15 lines/GHz. Several of these lines probably arise from vibrationally excited C$_5$H and C$_6$H and their $^{13}$C isotopomers, species for which we have no accurate laboratory frequencies (for example, U82924 and/or U82938 and U82995 could well correspond to the $J=61/2\\rightarrow59/2$ lines of $^2\\Pi_{1/2}$ CCCCC$^{13}$CH; U82677, U82737, U92034 and U92045 could correspond to different $\\Sigma$ and $\\Pi$ vibronic states of the lowest energy bending mode of C$_6$H). Others probably come from $^{13}$C and $^{29}$Si isotopomers of of SiC$_4$ (for example the J=29-28 and J=32-31 lines of $^{29}$SiC$_4$ at 87588 and 96648 MHz respectively and the J=27-26 line of SiC$^{13}$CCC at 82704 MHz). Finally, some unidentified lines are close to transition frequencies predicted for SiC$_5$ (U82746, U82876 and U82995), C$_5$S (U83038) and KCN (for which we have detected several additional lines that will be published elsewhere), and C$_9$H (U83011). Obviously, the U-line density is too high at present to claim the detection of any of these species, which all have a rich mm spectrum, on the basis of just one, two, or even three weak lines.  The HC$_4$N lines shown in Fig. 2 are cusped (see in particular Fig. 2a, where the lines have the highest signal-to-noise ratio). This implies that this molecule is mostly confined into the outer part of the circumstellar envelope, presumably in the 15\"-radius hollow shell where mos t of the free radicals and carbon chains are found \\citep{Gue93}. From the intensities of the four triplets we derive a rotational temperature of 25$\\pm$4 K, close to that derived for HC$_3$N.  The permanent dipole moment of HC$_4$N has been calculated by \\citet{Iku00} to be 4.3 D. Using this value and assuming thermal equilibrium at 25 K, we derive a column density for HC$_4$N of 1.5 10$^{12}$ cm$^{-2}$.  We have re-calculated the rotation temperature and the abundance of HCCN by adding to the 2-mm lines reported by \\citet{Gue91} the lines observed in our new 3-mm survey. We find a rotational temperature of 15$\\pm$2 K, slightly higher than our previous estimate (12$\\pm$2 K). The HCCN lines are also cusped and the HCCN column density, assuming a dipole moment of 3 D \\citep{Gue91}, is 1.4 10$^{13}$ cm$^{-2}$. We note that the increase in rotation temperature between HCCN and HC$_4$N is not surprising as a similar increase is observed between the cyanopolyynes HC$_3$N and HC$_5$N.  The HCCN/HC$_4$N abundance ratio is found equal to 9. This is about twice the decrement observed between the successive members of the cyanopolyynes family, HC$_3$N/HC$_5$N, but $20-40$ times less than the HC$_3$N/HCCN ($\\simeq 200$) and HC$_5$N/HC$_4$N (400) abundance ratios, which confirms that, like the carbon chain radicals, the even and odd HC$_n$N families have parallel, but distinct formation paths. We have also searched in our 3mm line survey for the linear triplet HC$_6$N, but without success. The 3$\\sigma$ upper limit to the column density of this species, assuming a rotation temperature of 25 K, is 1 10$^{13}$ cm$^{-2}$, not low enough to constrain significantly the HC$_4$N/HC$_6$N ratio, but already 25 times lower than the HC$_7$N column density \\citep{Cer87}.  As noted above, the carbon chain molecules C$_{2n+1}$N and HC$_{2n+1}$N are thought to be formed in IRC+10216 from the reaction of radicals such as CCH and CN with carbon-chain molecules or radicals \\citep{Che93,Mil00}. Similarly, we could consider that HC$_4$N may form through the reaction of C$_3$N and C$_3$H with HCN and CH$_2$, e.g., C$_3$H + HCN $\\rightarrow$ HC$_4$N + H, or may be C$_3$N + CH$_2 \\rightarrow $HC$_4$N + H. All the reactants are relatively abundant in the outer layers of the circumstellar envelope and we expect, by analogy with similar reactions studied in the laboratory, that at least one of these reactions could proceed rapidly.  Another path proposed for the formation of cyanopolyynes involves ion-molecule reactions yielding H$_3$C$_n$N$^+$ or H$_2$C$_n$N$^+$, followed by the dissociative recombination of these ions \\citep{Gla86, How90}. The low abundance of ions in IRC+10216 makes this path relatively slow \\citep{Che93}. Moreover, in the case of HCCN and HC$_4$N, the reaction of HCN with CH$_3^+$, leading to H$_2$C$_2$N$^+$ + H$_2$ (or H$_3$C$_2$N$^+$ + H), does not proceed \\citep{Gue91}, making the ion-molecule path even slower and less likely than in the case of cyanopolyynes.  Finally, a last process, formation on dust grains followed by photodesorption when the grains reach the outer envelope layers, has been advocated in the case of IRC+10216 \\citep{Gue93}. Such a mechanism may well produce HCCN, HC$_4$N, and the cyanopolyynes observed in the outer shell, but would provide no simple explanation to why the abundance decrement within each family seems to be so constant.   The HC$_4$N lines detected during this study belong to the linear triplet isomer. We have unsuccessfully searched for the singlet bent HC$_4$N isomer and for the c$_3$-ring cyclic isomer. Those latter have a permanent dipole moment almost as large as the linear isomer: 2.96D ($\\mu_a$=2.9$\\;$D, $\\mu_b$=0.5$\\;$D) and 3.48$\\;$D ($\\mu_a$=3.15$\\;$D, $\\mu_b$=1.49$\\;$D), respectively, versus 4.3 D for linear HC$_4$N \\citep{McC99a, McC99b, Iku00}. The upper limits to the column density of these species are $\\simeq$ 3 10$^{12}$ cm$^{-2}$, assuming the same rotational temperature than for linear triplet HC$_4$N (25 K).   The most accurate calculations of the HC$_4$N structure to date (level CCSD(T)/cc-pCVTZ) predict that the singlet ring isomer of HC$_4$N lies 4.2 Kcal/mol below the linear triplet isomer \\citep{Iku00}. If this result is confirmed, the non-detection of the ringed isomer (and of the bent isomer) would further stress the discrimination against ringed structures already noted in the case of carbon chain radicals.       \\begin{table}[htb] {\\bf Table 1: Observed HC$_4$N line parameters}\\\\ \\begin{tabular}{lccr} \\hline Obs. Freq.  & Calc. Freq. & Transition      & $\\int T_A^*dv$    \\\\ MHz        & MHz      &(N,J$\\rightarrow$N',J') &mK.kms$^{-1}$ \\\\ \\hline 82855.6(6) &     82855.52 & 18 17 -17 16   &110(20)\\\\ 82884.1(10)$^b$& 82883.57 & 18 18 -17 17   & 94(14)\\\\ 82905.4(3)     & 82905.41 & 18 19 -17 18  & 86(10)\\\\ \\\\ 87463.5(10)$^b$ & 87463.15 & 19 18 -18 17     &127(35)\\\\ 87487.9(10) &  87487.87   &  19 19 -18 18     &135(35) \\\\ 87507.6(5) &  87507.31   &19 20 -18 19 \t&142(20) \\\\ \\\\ 92070.0(10)$^b$ &92070.22  &20 19 -19 18 \t&90(24) \\\\ 92092.4(10)  & 92092.13    &20 20 -19 19 \t&105(24) \\\\ 92109.2(6)   & 92109.48    &20 21 -19 20 \t&90(18) \\\\ \\\\ 96677.1(5)   & 96676.82    &21 20 -20 19 \t&75(15) \\\\ 96696.0(10)$^b$ & 96696.32  &21 21 -20 20  &100(23) \\\\ 96712.1(5)    &96711.88     &21 22 -20 21  &75(15) \\\\  \\hline \\end{tabular} \\caption{ The number in parenthesis are 1$\\sigma$ uncertainties on the last digit, derived from least square fits; they do not include the 5\\% calibration uncertainty. $^b$ denotes a partly blended line. The frequencies in col. 2 are calculated from the spectroscopic constants measured by \\citet{Tan99} in the laboratory.} \\end{table}"
    },
    "0410/astro-ph0410182_arXiv.txt": {
        "abstract": "The existence of magnetic fields associated with the intracluster medium in clusters of galaxies is now well established through different methods of analysis.  Magnetic fields are investigated in the radio band from studies of the rotation measure of polarized radio galaxies and the synchrotron emission of cluster-wide diffuse sources. Other techniques include X-ray studies of the inverse Compton emission and of cold fronts and magneto hydrodynamic simulations.  We review the main issues that have led to our knowledge on magnetic fields in clusters of galaxies.  Observations show that cluster fields are at the $\\mu$G level, with values up to tens of $\\mu$G at the center of cooling core clusters.  Estimates obtained from different observational approaches may differ by about an order of magnitude. However, the discrepancy may be alleviated by considering that the magnetic field is not constant through the cluster, and shows a complex structure.  In particular, the magnetic field intensity declines with the cluster radius with a rough dependence on the thermal gas density. Moreover, cluster magnetic fields are likely to fluctuate over a wide range of spatial scales with values from a few kpc up to hundreds kpc.  Important information on the cluster field are obtained by comparing the observational results with the prediction from numerical simulations.  The origin of cluster magnetic fields is still debated.  They might originate in the early Universe, either before or after the recombination, or they could have been deposited in the intracluster medium by normal galaxies, starburst galaxies, or AGN. In either case, magnetic fields undergo significant amplification during the cluster merger processes. ",
        "introduction": "Experience indicates that most of the matter in the Universe is composed of ionized or partially ionized gas permeated by magnetic fields. Celestial objects are magnetized and magnetic fields of significant strength are found everywhere in the interstellar space, and over small and very large scales, in the extragalactic universe. In general, small compact objects have the largest magnetic field strengths, while larger low-density objects have weaker magnetic fields.  The Earth has a bipolar field of about 0.5 G at its surface, originating from an idealized current due to the charged fluid motion going circularly in a ring inside the liquid molten metallic core.  On the Jupiter surface the magnetic field is about 4 G, owing to the fast Jupiter rotation.  In the interplanetary space of the solar system the magnetic fields are of the order of 50 $\\mu$G.  On the Sun, the magnetic field is of 10 G at the poles, while localized sunspots on the surface near the equatorial zone of the Sun, and more generally of a star, can have magnetic field strengths of 2000 G. In protostellar envelopes and protostars, fields are of $\\sim$ 1 mG.  A bipolar field is ``frozen'' into the gas of a star during the contraction from a normal star to a degenerate star. It will remain bipolar-shaped but its intensity will increase as r$^{-2}$, thus magnetic fields of pulsars and neutron stars are of the order of 10$^{12}$ G, those of white dwarfs are around 10$^6$ G.  A widespread field of $\\sim$ 5 $\\mu$G, characterized by a spiral shape, is present in the Galaxy. At the Galaxy nucleus, highly organized filaments with strength of $\\sim$ 1 mG are detected. Fields in other spiral galaxies are of $\\sim$ 10 $\\mu$G on average, with values up to $\\sim$ 50 $\\mu$G in starburst galaxies and $\\sim$ 30 $\\mu$G in massive spiral arms.  Fields of $\\sim$ $\\mu$G are found in the radio emitting lobes of radio galaxies. Fields of similar or weaker strength are detected in the intracluster medium of clusters of galaxies, and in more rarefied regions of the intergalactic space.  Upper limits of $\\simeq 10^{-8}- 10^{-9}$ G have been obtained for the cosmological fields at large redshift.  In this review large-scale magnetic fields in clusters of galaxies will be analyzed.  In the last years the presence of cluster magnetic fields has been unambiguously proven and the importance of their role has been recently recognized. The study of cluster magnetic fields is relevant to understand the physical conditions and energetics of the intracluster medium. Cluster magnetic fields provide an additional term of pressure and may play a role in the cluster dynamics.  They couple cosmic ray particles to the intracluster gas, and they are able to inhibit transport processes like heat conduction, spatial mixing of gas, and propagation of cosmic rays.  They are essential for the acceleration of cosmic rays and allow cosmic ray electron population to be observed by the synchrotron radiation.  Despite many observational efforts to measure their properties, our knowledge on cluster magnetic fields is still poor. Overviews on observational and theoretical arguments can be found in the literature\\cite{kro94,bec96,grarub01,cartay02,wid02,gio04}.  The focus of this review is primarily observational, however, we present the basic theory needed for the interpretation of the data. We analyze some of the main issues that have led to our knowledge on magnetic fields in clusters of galaxies and discuss some of their limitations.  An outline of the review is as follows: In Sec. 2 we summarize some general properties of clusters of galaxies. Sec. 3 is devoted to theoretical background related to the detection of cluster magnetic fields and to the estimate of their strengths.  We recall the basic theory concerning synchrotron radiation, inverse Compton radiation and Faraday rotation. These are the main observed features which provide information on the cluster magnetic fields.  The observational results of cluster magnetic fields through synchrotron radio and inverse Compton hard X-ray emissions are described in Secs. 4 and 5. In Sec. 6 we give the results obtained by analyzing rotation measures of radio galaxies located within or behind clusters of galaxies. In Sec. 7 we present cluster magnetic fields detected through the study of cold fronts.  In Sec. 8 we report the evidence for a radial decline of cluster magnetic fields. In Sec. 9 we discuss how magnetic field values obtained with different approaches can be reconciled.  In Sec. 10 we summarize the results of a numerical technique which can significantly improve our interpretation of the data and thus the knowledge of the strength and structure of magnetic fields.  In Sec. 11 we briefly review the current knowledge on the cluster magnetic field origin and amplification.  Throughout this paper we assume the $\\lambda$CDM cosmology with $H_0$ = 71 km s$^{-1}$Mpc$^{-1}$, $\\Omega_m$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7, unless stated otherwise. ",
        "conclusions": "Our knowledge of the magnetic field properties in galaxy clusters has significantly improved in recent years, owing to the improved capabilities of radio and X-ray telescopes. It is well established that $\\mu$G level magnetic fields are widespread in the ICM, regardless of the presence of large-scale diffuse radio emission. The magnetic field strengths show almost an order of magnitude scatter between clusters, or within a given cluster, and are extreme in cluster cooling cores. For such large fields the magnetic pressure is comparable to or larger than the gas pressure derived from X-ray data, suggesting that magnetic fields may play a significant role in the cluster dynamics.  The observations are often interpreted in terms of the simplest possible model, i.e. a constant field throughout the whole cluster.  However, a decline with radius is expected if the intensity of the magnetic field results from the compression of the thermal plasma during the cluster gravitational collapse.  Observational evidence of magnetic field profiles has been derived in some clusters. Moreover the magnetic field could show complex structure with a range of coherence scales.  The study of cluster magnetic fields has gained a big interest in recent years, leading to several new observations as well as simulations.  There are, however, still many questions to answer: are the fields filamentary, what are the coherence scales, to what extent do the thermal and non-thermal plasmas mix in cluster atmospheres, how do the fields extend, what is the radial trend of the field strength, how does the field strength depend on cluster parameters such as the gas temperature, metallicity, mass, substructure and density profile, how do the fields evolve with cosmic time, and finally how were the fields generated?  New generation instruments in the radio band, as the EVLA, LOFAR and SKA, are rather promising and will establish a clear connection between radio astronomical techniques and the improvement in the knowledge of the X-ray sky. There are various satellite missions, as {\\it ASTRO-E2}, {\\it XEUS} and {\\it Constellation X}, which will map the X-ray sky at low and high energies in the next years. These will provide a more precise knowledge of the X-ray surface brightness of clusters, i.e. of their thermal gas density, allowing a more accurate and correct interpretation of  the sensitive RM measurements. The detection of HXR non-thermal emission will provide independent measurements of the magnetic fields. The accurate experimental determination of large-scale magnetic fields in the intracluster medium will thus be possible. The detection of synchrotron radiation at the lowest possible levels will allow the measurement of magnetic fields in even more rarefied regions of the intergalactic space, and the investigation of the relation between the formation of magnetic fields and the formation of the large-scale structure in the universe."
    },
    "0410/astro-ph0410307_arXiv.txt": {
        "abstract": "The wide-band X-ray spectra of the high mass X-ray binary Cygnus X-3 exhibits a pivoting behavior in the `low' (as well as `hard') state, correlated to the radio emission. The time scale of the soft and hard X-rays' anti-correlation, which gave rise to the pivoting feature, was found to be less than a day from the monitoring observations by {\\it RXTE}--ASM and {\\it CGRO}--BATSE. In this {\\it Letter} we report the detection of a lag of $\\lesssim$ 1000s in the anti-correlation of the hard X-ray emission (20--50 keV) to that of the soft X-ray emission (2--7 keV), which may be attributed to the viscous time scale of flow of matter in the accretion disk. This suggests the geometrical picture of a truncated accretion disc with a Compton cloud inside the disc, the relative sizes of which determine the spectral shape. Any change in the disc structure will take place in a viscous time scale, with corresponding anti-correlated change in the Compton cloud. We also report the pivoting in the spectra in one span of a pointed observation when an episode of the rearranging of the accretion system is serendipitously observed. This is the first such observation of hard X-ray delay seen in the persistent Galactic microquasars, within the precincts of the hard state. ",
        "introduction": " ",
        "conclusions": "The strong binary modulation of Cygnus X-3 has prevented any serious study on the variability of this source at times-scales of few hundred seconds to about a day. \\citet{singh02aa} have shown that the binary period (and the template of its variation) has remained consistent for more than 25 years  with a linear decay with no second derivative of the period. \\citet{choudhury04aa} have used this information and shown that this ephemeris can be used to remove the large amplitude smooth binary variation without recourse to any derivation of phase during individual observations - the only free parameters being the amplitude of variation and the mean count rate. Using this method, we have corrected for the large amplitude binary variations and detected anti-correlated hard X-ray delays in five of the observations.  The only assumption incorporated in our method is that the occasional jitters seen in the binary modulation is not due to any binary instability \\citep{choudhury04aa}, rather we explore whether any residual energy dependent phase jitter and/or variable energy dependent template could have caused the anti-correlated delay presented here. In all observations we find a strong positive correlation with zero phase indicating the similarity of the template in the two energy bands. Further, more than 90\\% of the binary variability in Cygnus X-3 is energy independent \\citep[see][]{rajeev94apj} and the large amplitude variations in high phase are correlated in all energies. The anti-correlated behavior is seen in the slow varying components which is unrelated to the binary variation. This was evident when we examined the variation of the hardness ratio compared to the lightcurves.  The strongest evidence for the reality of the anti-correlation comes from the model independent spectral pivoting seen in one of the observations (see Fig 4). \\citet{choudhury03apj, choudhury02aa} have shown that this type of spectral pivoting at 10--20 keV energies occurs when one takes spectrum at different days at different count rates. The fact that such spectral pivoting is indeed seen in one single observation shows that the soft and hard X-ray anti-correlation occurs in a time scale less than the binary period and the cross-correlation analysis reveals the delay in the time scale close to a thousand seconds.  \\vskip 0.3cm { \\includegraphics[width=7cm,height=4cm,angle=0]{figure_4.ps} \\\\ {{\\bf Figure 4.} The wide-band X-ray spectra of the source on MJD 50321, for the soft and the hard regions of the lightcurve, resulting in comparative softer and harder spectral distribution, displaying the `in situ' pivoting in the region of 10--20 keV.\\label{fig4}} } \\vskip 0.2cm  The three-way association between soft X-rays, hard X-rays and radio emission in Cygnus X-3 maybe explained by the geometry of truncated, optically thick, accretion disc, with a Comptonizing cloud of energetic electrons inside which may have advective \\citep{narayan94apjl} or bulk motion \\citep{chakrabarti95apj}. The soft X-ray flux determines the accretion disk structure, while the Compton cloud within the truncated disc is responsible for the hard X-ray emission and it also drives the outflow \\citep[see][for contrasting paradigms]{meier01apjl, das99cqg}. The delay presented in this work signifies the disk readjustment time scale which should correspond to the location of the truncation radius. The viscous time $t_{vis}$ can be calculated  for the radiation pressure dominated optically thick accretion disc as $$t_{vis} = 30 \\alpha^{-1} M^{-1/2} R^{7/2} \\dot{M}^{-2} s$$ where $\\alpha$ is the viscosity parameter in units of 0.01, M is the mass of the compact object in solar mass units, R is the radial location in the accretion disc in units of 10$^7$ cm, $\\dot{M}$ is the mass accretion rate in units of 10$^{18}$ g/s \\citep[see][]{belloni97apj_b,fender04araa}. Taking $\\alpha$ = 1, M = 10 and $\\dot{M}$ = 3 (corresponding to one fifth of the Eddington accretion rate), we get R $\\sim$ 7 for a viscous time scale of 1000 s. Thus, if Cyg X-3 harbors a black hole of 10 M$_\\odot$, the observed delay signifies a CENBOL location of $\\sim$ 25 Schwartzchild radius.  It should be noted that the X-ray emission shows a strong correlation with the radio emission, but the radio emission does not show any binary modulation. It would be interesting to repeat the type of work presented here with a simultaneous radio monitoring observations so that a clear handle can be obtained for the location of the radio emission."
    },
    "0410/astro-ph0410131_arXiv.txt": {
        "abstract": "Sixteen RR Lyrae variables from the QUEST survey that lie in the leading arm of the tidal stream from the Sagittarius dSph galaxy have been observed spectroscopically to measure their radial velocities and metal abundances.  The systemic velocities of 14 stars, which were determined by fitting a standard velocity curve to the individual measurements, have a sharply peaked distribution with a mean of $33$ km/s and a standard deviation of only 25 km/s.  The [Fe/H] distribution of these stars has a mean of $-1.76$ and a standard deviation of 0.22.  These measurements are in good agreement with previous ones from smaller samples of stars.  The mean metallicity is consistent with the age-metallicity relation that is observed in the main body of the Sgr dSph galaxy.  The radial velocities and the distances from the Sun of these stars are compared with recent numerical simulations of the Sgr streams that assume different shapes for the dark matter halo.  Models that assume a oblate halo do not fit the data as well as ones that assume a spherical or a prolate distribution.  However, none of the fits are completely satisfactory.  Every model fails to reproduce the long extent of the stream in right ascension ($36\\degr$) that is seen in the region covered by the QUEST survey.  Further modeling is required to see if this and the other mismatches between theory and observation can be removed by judicial choices for the model parameters or instead rule out a class of models. ",
        "introduction": "From the time of its discovery \\citep{iba94} the Sagittarius (Sgr) dwarf spheroidal (dSph) galaxy has been considered the prototypical example of the tidal destruction and assimilation of a small galaxy into the halo of a larger one.  Long tidal streams of stars from the Sgr dSph, which were predicted by theoretical models of the destruction process \\citep[e.g.][]{joh99}, have now been observed in several directions with different stellar tracers by the recent large-scale surveys of the halo.  An all-sky view is provided by the M giants detected by the 2MASS survey, which reveals two prominent tidal streams, one leading and one trailing the main body of the dSph galaxy \\citep{maj03}. In addition to these metal-rich stars, the streams also contain numerous RR Lyrae stars \\citep{ive00,viv01,vz04}, which, as we show later, are metal poor in addition to being very old.  The detailed study of the Sgr tidal streams is important for several reasons.  First, it gives us a close-up look into the process of tidal destruction and galaxy merging, which may be responsible for the formation of the halos of disk galaxies, as suggested by the hierarchical models of galaxy formation and by numerous studies of the Milky Way's halo over the past 25 years \\citep[see review by][]{fre02}. Second, comparisons between the stellar populations at different points in the streams and the main body of Sgr may reveal interesting facets of galaxy evolution in the presence of strong tidal forces, in particular if dynamical models can provide the chronology of the formation of the streams.  Finally, tidal streams are powerful probes of the shape of the potential of the Milky Way \\citep{joh98}, which at the large galactocentric distances of the Sgr streams may be dominated by the dark matter halo.  There has been some recent debate, however, over the usefulness of the Sgr streams for this purpose. N-body simulations of modern hierarchical galaxy formation picture predict that dark matter halos should be flattened by appreciable amounts \\citep[e.g.][]{bul02}. The fact that the Sgr tidal streams can be traced with both Carbon stars and M giants along a narrow great circle on the sky has been used as an argument in favor of an almost spherical halo \\citep{iba01,maj03}; otherwise the stars in the stream would disperse and the resulting stream would be much wider and perhaps not even recognizable as a feature above the background of halo stars \\citep{may02}.  While this alone may not rule out the hierarchical models, the apparent deviation of the most familiar dark matter halo from the predicted norm is somewhat discomforting.  More recently \\citet{hel04} has argued that only old streams will be dispersed by the precession generated by a flattened potential and that the Sgr streams are too young dynamically, with most of the stars stripped out in the past few perigalacticon passages \\citep{hel01,hel04,mar04}, for this to have occurred.  The theoretical models of the streams under different assumptions of halo shape, \\citep[e.g.][]{hel04}, are in fact different on small scales.  The precisions that are necessary to test observationally these models are achievable with RR Lyrae variables, but with few other halo tracers.  The spatial distribution of the Sgr streams is well documented by the investigations of \\citet{yan00,ive00,iba01,viv01,mar01,new02,maj03,mar04} and \\citet{vz04}. The tidal streams should be also coherent structures in velocity space as well. The most complete work on the subject to the date is that of \\citet{maj04}, who obtained radial velocities for $\\sim 300$ M giant candidates, most of them belonging to the Sgr trailing (Southern Galactic Hemisphere) stream.  They find that the feature is indeed a coherent structure in velocity space.  In this paper, we present a spectroscopic study of RR Lyrae stars from the QUEST survey \\citep{viv04} belonging to the Sgr leading arm (Northern Galactic Hemisphere). These observations provide a unique set of detailed observations in this part of the stream: location on the sky, distance from the Sun, spatial dimensions, radial velocities and metallicities. Many of these parameters are better determined in this region than in any other part of the stream. This is because RR Lyrae stars are excellent standard candles, which enables us to determine their distances with precisions of $\\sim 6\\%$.  The precisions obtained with other stellar tracers are substantially worse, for example they are $20-25\\%$ for K and M giants \\citep{doh01,maj03}. It is important to emphasize that our sample is made of {\\it bona-fide} RR Lyrae stars (not just candidates), which were discovered using well-sampled light curves.  The purpose of our observations was to constrain models of the disruption of the Sgr galaxy, and we show that none of recent models of Sgr streams are consistent with these observations.  This paper is organized as follows: in \\S2 we describe the target selection, spectroscopic observations and data reduction. Section \\S3 describes the method for obtaining radial velocities of these pulsating stars, and \\S4 explains the determination of metal abundances.  In \\S5, we compare our results with different theoretical models of the Sgr streams.  We briefly summarize our results in \\S6. ",
        "conclusions": "Our spectroscopic observations have yielded radial velocity and metallicity measurements for a subsample of the RR Lyrae variables that were discovered by the QUEST survey in the leading arm of the Sgr tidal stream.  These data and the positional information provided by the whole sample of QUEST RR Lyrae variables in the stream provide powerful diagnostics to test numerical simulations of the Sgr tidal stream.  While recent models of the destruction of the Sgr galaxy reproduce the general properties of the Sgr tidal streams, they fail to reproduce the details revealed by the RR Lyrae variables.  This is particularly true of models that assume a oblate flattening of the dark matter halo.  Models that assume spherical and prolate dark matter halos provide better fits to the data, but ones that are still marred by some inconsistencies.  The most striking of these is the failure of any model to reproduce is the observed span of the stream in right ascension ($36\\degr$ or $\\sim 30$ kpc).  Our metallicity measurements show that the RR Lyrae stars in this part of the stream belong to a metal-poor population with $\\langle {\\rm[Fe/H]} \\rangle = -1.77$.  This mean value is consistent with the age-metallicty relation observed by \\citet{lay00} in the central parts of the Sgr galaxy."
    },
    "0410/astro-ph0410241_arXiv.txt": {
        "abstract": "This paper presents an introduction to the astrophysics of cosmic rays and diffuse \\grays\\ and discusses some of the puzzles that have emerged recently due to more precise data and improved propagation models: the excesses in Galactic diffuse \\gray\\ emission, secondary antiprotons and positrons, and the flatter than expected gradient of cosmic rays in the Galaxy. These also involve the dark matter, a challenge to modern physics, through its indirect searches in cosmic rays. Though the final solutions are yet to be found, I discuss some ideas and results obtained mostly with the numerical propagation model GALPROP. A fleet of spacecraft and balloon experiments targeting these specific issues is set to lift off in a few years, imparting a feeling of optimism that a new era of exciting discoveries is just around the corner. A complete and comprehensive discussion of all the recent results is not attempted here due to the space limitations. ",
        "introduction": "Research in astrophysics of cosmic rays (CR) and \\grays\\ provides a fertile ground for studies and discoveries in many areas of particle physics and cosmology. Examples are the search for dark matter (DM), antimatter, new particles, and exotic physics; studies of the nucleosynthesis, acceleration of nuclei and their transport through CR spectra and composition analysis; the effects of heliospheric modulation; the origin of Galactic and extragalactic diffuse \\gray\\ emission; formation of the large scale structure of the universe as traced by \\grays.  In its turn, the astrophysics of CR and \\grays\\ depends very much on the quality of the data and their proper interpretation.  The accuracy of data from current CR experiments  on interplanetary spacecraft such as Ulysses, Advanced Composition Explorer (ACE), and the two Voyagers, specialized balloon-borne experiments such as Super-TIGER, BESS-Polar, CREAM, and the new space-based missions such as Pamela and AMS far exceeds the accuracy of the current propagation models, of which the ``leaky-box\" model has remained one of the main research tools for the last 50 years. The new major \\gray\\ observatory GLAST will improve the sensitivity for the diffuse high-energy \\grays\\ produced in CR interactions in the interstellar medium (ISM) by a factor of 30. These near-future missions are specifically designed to search for the signatures of DM, search for antimatter, study the diffuse Galactic and extragalactic \\grays, and provide outstanding quality data on \\gray\\ sources and CR species in a wide energy range. This presents a great opportunity for new discoveries that requires accurate, testable and readily accessible modeling to exploit.  On the other hand, the whole of our knowledge is based on measurements done only at one point on the outskirts of the Galaxy, the solar system, and the assumption that particle spectra and composition are (almost) the same at every point of the Galaxy. The latter may not necessarily be correct. \\grays\\ and radio-waves (synchrotron) are able to deliver the information directly from distant Galactic regions thus complementing that obtained from CR measurements. Some part of the diffuse \\grays\\ is produced in energetic nucleon interactions with gas via $\\pi^0$ production, another is produced by electrons via inverse Compton scattering and bremsstrahlung. These processes are dominant in different parts of the spectra of \\grays, therefore, if deciphered the \\gray\\ spectrum can provide information about the large-scale spectra of nucleonic and leptonic components of CR. Combining all the data collected by different experiments into a realistic interpretive model of the Galaxy, we have a better chance to understand the mechanisms of particle acceleration and the role of energetic particles in the dynamics and evolution of the Galaxy and make essential progress in related areas. ",
        "conclusions": "The choice of topics discussed in this paper is personal and by no means complete. More complete list would include the origin of 511 keV line from the inner Galaxy, \\gray\\ bursts, ultra-high energy CR, as well as a more comprehensive discussion of the DM, SUSY, and dark energy. Other contributions to the Conference will fill these gaps.  This work was supported in part by a NASA Astrophysics Theory Program grant."
    },
    "0410/astro-ph0410288_arXiv.txt": {
        "abstract": "Gaseous halos of nine nearby galaxies (with redshifts cz $<$ 6000 \\kms ) were probed at large galactocentric radii using background quasars observed with HST GHRS and STIS. The projected quasar-galaxy separations range from 55 to 387~$h^{-1}_{75}$ kpc.  Ly$\\alpha$ absorption lines were successfully detected in the spectra of five quasars, at impact parameters of up to $\\sim$ 170 $h^{-1}_{75}$~kpc from the center of the nearby galaxy, and in each case at wavelengths consistent with the galaxy's redshift.  Our observations include the lowest redshift Ly$\\alpha$ lines detected to date. HI velocity fields were obtained at the VLA for three of the galaxies in our sample (and in one case was available from the literature), to derive their rotation curves. When comparing the inner rotation curves of the galaxies with the velocity at large radius provided by the Ly$\\alpha$ line it is apparent that it is very difficult to explain the observed Ly$\\alpha$ velocity as due to gas in an extended rotating disk. In most cases one would need to invoke large warps in the outer gas disks and also thick gas disks in order to reconcile the observed velocities with the predicted ones. Indeed, in one case the Ly$\\alpha$ line velocity indicates in fact counter-rotation with respect to the inner disk rotation. In light of these results we conclude  that in a typical galaxy there is no longer detectable atomic gas corotating in an extended  disk at radii $> 35$ $\\alpha ^{-1}$, where $\\alpha^{-1}$ is the stellar disk exponential scale-length. The cosmic web is the most likely origin for the detected Ly$\\alpha$ lines. Our observations confirm  the Bowen \\etal (2002) correlation of equivalent widths  with the local volume density of galaxies around the sightline, and the observed equivalent widths of the lines are consistent with expectations of the cosmic web. ",
        "introduction": "Rotation curves are the best tool to study the dark matter halos of galaxies, and the most extended ones will better constrain the dark matter halo parameters. Rotation curves derived from neutral hydrogen (HI) observations enable us to probe the dark matter potential to distances of at most 20 to 25 $\\alpha^{-1}$ (where $\\alpha^{-1}$ is the optical disk exponential scale length), for gas-rich late-type spiral galaxies with particularly extended HI (eg. NGC2841 from Begeman 1987 or NGC2915 from Meurer \\etal ~1996). But it is not possible to obtain any kinematical data with 21cm emission further out than the radius at which the gas density falls to levels of a few times $10^{19}$ atoms cm$^{-2}$, even with very sensitive observations, as the HI column falls off sharply beyond this point (eg. the 'HI Edge' observed in NGC3198, van Gorkom 1991). This sharp truncation of the HI distribution is believed to be the result of ionization of the atomic gas by the extragalactic UV radiation field (as predicted by Silk \\& Sunyaev 1976). At lower column densities the neutral fraction is therefore predicted to drop off dramatically in comparison to the total (ionized) hydrogen column (see Maloney 1993).  Some studies have attempted to detect the faint ionized gas just beyond the HI limit from recombination emission in H$\\alpha $. Depending on the volume density of the plasma at the critical column density where the outer disk becomes optically thin to this radiation (about $N_{HI} \\sim 3\\times 10^{19}$~cm$^{-2}$ for NGC3198), the emission measure from recombination radiation expected to be emitted by this ionized gas is only 0.025 to 0.25~cm$^{-6}$pc (Maloney 1993).  The only reported detection of recombination radiation so far is that of NGC253 for which sensitive Fabry-Perot observations by Bland-Hawthorn \\etal ~(1997) managed to extend the HI rotation curve from 1.2 to 1.4 $R_{25}$, with measured surface brightness values in H$\\alpha$ from 80 to 40 mR (millirayleighs, where 1 rayleigh = 10$^6$/4$\\pi$ photons cm$^{-2}$s$^{-1}$sr$^{-1}$ or 2.78 cm$^{-6}$pc at H$\\alpha$). But at this level, as pointed out by the authors, the emission is probably not due to recombination after ionization by the metagalactic background, but rather from hot young stars near the center of the galaxy ionizing the warped outer HI disk.  If indeed spiral galaxies have extended (mostly ionized) gas disks it should then be possible to detect them in absorption in the spectrum of a bright background quasar, down to HI column densities of only $\\sim $10$^{13}$ atoms cm$^{-2}$. With such a suitably placed UV-bright background object one could obtain kinematical information on that galaxy's outer rotation curve out to 30 to 40 $\\alpha ^{-1}$ or more.  At such radii the data could potentially reveal the extent of the dark matter halo, and constrain the mass profile out to such large galactocentric radii.  This could then provide strong tests of cosmological galaxy formation scenarios.  For example, according to the rotation curves predicted by Cold Dark Matter-dominated N-body simulations (e.g.~Navarro, Frenk \\& White 1996) in low-mass dwarf galaxies we should detect a turn-down in the rotation curve if we could get data points just 15 or 20\\% further out in radius than currently mapped by HI observations.  This matching of a Ly$\\alpha$ absorption line velocity with the kinematics of the nearest galaxy has already been attempted, originally by Barcons \\etal (1995), who found that for two galaxies at z=0.075 and 0.09 the Ly$\\alpha$ velocities as measured from background QSOs spectra at impact parameters 64 $h^{-1}_{75}$kpc (= 16$\\alpha ^{-1}$) and 83 $h^{-1}_{75}$kpc (=35$\\alpha ^{-1}$), were consistent with gas in extended gaseous haloes co-rotating with the inner stellar disks (with $h^{-1}_{75} = H_0/75$ where $H_0$ is the Hubble constant and $q_0=0$). On the other hand, Hoffman \\etal (1998) found that the dwarf galaxy MCG+00-32-16, which is the closest galaxy to the low redshift Ly$\\alpha$ system in the sightline of 3C273, has an HI disk which would be counter-rotating compared to the Ly$\\alpha$ velocity. In this case however the impact parameter is 204 kpc which for the dwarf galaxy represents about 153 $\\alpha ^{-1}$, in other words an excessively large radius to expect the extended gas halo to reach. More recently Steidel \\etal (2002) have obtained long-slit spectra of 5 intermediate redshift galaxies ( 0.44 $\\leq$ z$\\leq$0.66) in the proximity of MgII absorbers sightlines (with projected separations from 19 to 95 $h^{-1}_{75}$kpc), and succeeded in reasonably matching the MgII absorbers' velocities with the galaxies' rotation curves, but only when setting various disk thicknesses and velocity scale heights (describing the velocity fall-off with z-height above the plane) case by case.  Our approach here is the reverse. Instead of identifying and analysing the galaxies found near the line-of-sight of a quasar exhibiting a known Ly$\\alpha$ line, we start by selecting nearby galaxies, normal and well-behaved kinematically, hence easy to subsequently model dynamically, and search for bright background quasars at some particular impact parameter away, ie: at the interesting radii beyond the galaxy's HI envelope and up to about 50 $\\alpha^{-1}$. We then observe these quasars with HST, first to attempt to detect a Ly$\\alpha$ line arising from the extended galaxy's disk, and then to use this Ly$\\alpha$ line velocity to probe the outer dynamics of the galaxy.   Our study does not aim at elucidating the overall nature of all Ly$\\alpha$ forest lines (other studies are better designed for this, by surveying all lines along several QSO sightlines and getting numerous spectra of galaxies in the surrounding fields eg. Penton \\etal 2002, Morris \\etal 2002).  Indeed, while it is now generally accepted that most metal-line absorption systems found in QSO's spectra are associated with gaseous galaxy halos, the situation for the Ly$\\alpha $ absorption systems has been long under debate.  Some studies claim that (at least the stronger) Ly$\\alpha$ lines are associated to extended halos of galaxies, while most believe that the majority of the lines occur in the intergalactic medium, tracing the 'cosmic web' (cf. van Gorkom et al.~1996), this intricate network of filaments and sheets of gas predicted by hydrodynamical simulations (eg. Dav{\\'e} \\etal 1999).  Some convincing arguments of the latter are the fact that several Ly$\\alpha$ lines have been detected in voids (McLin \\etal 2002), also confirmed by deep HI observations that would have detected dwarf galaxies and other low surface brigthness galaxies (Shull \\etal 1998). Also observations of double QSOs sightlines find Ly$\\alpha$ lines in common in both spectra on scales of typically half a Mpc (Dinshaw \\etal 1997, 1998). Here we are testing the simple conjecture that gaseous galactic disks extend further beyond that which can be probed in emission in HI, by searching for a corresponding Ly$\\alpha$ line.  More importantly, a successful detection, can be used as a dynamical probe of the outer halo of the galaxy.  In Section 2 we explain the selection of our targets, section 3 describes the observations (HST and radio), then section 4 discuss the results of matching the Ly$\\alpha$ velocities with the galaxies' rotation curves and the implications. ",
        "conclusions": "Sightlines of QSOs at impact parameters 55 to 387 $h^{-1}_{75}$~kpc from nine nearby galaxies (at systemic velocities from 566 to 5538 \\kms) were probed, using GHRS and STIS. In five cases a Ly$\\alpha$ absorption line was successfully detected at a velocity coincident with the galaxy's. Some of these Ly$\\alpha$ lines are the lowest-redshifts Ly$\\alpha$ lines ever detected, and in fact most of them are sitting on the wing of the Milky Way Ly$\\alpha$ damped line (hence introducing uncertainties in the derived equivalent widths).  The positive detections occur for QSO-galaxy pairs with separations from 55 to 169 $h^{-1}_{75}$~kpc, while non-detections are for pairs with separations from 162 to 387 $h^{-1}_{75}$~kpc. Although at first sight it appears that the Ly$\\alpha$ lines might be genuinely arising in the extended disks of the target galaxies, this is not what transpires from inspecting the detailed kinematics of these galaxies.  HI velocity fields were obtained at the VLA for our galaxies (and in one case was available from the literature), to derive their rotation curves. When comparing the inner rotation curves of the galaxies with the velocity at large radius provided by the Ly$\\alpha$ line it appears that it is very difficult to explain this Ly$\\alpha$ velocity as part of the extended gaseous rotating disk.  In most cases one would need to invoke large warps in the extreme outer gas disks to reconcile the observed velocities with the predicted ones. Worse, in some cases the Ly$\\alpha$ line velocity indicates in fact counter-rotation with respect to the inner disk rotation.   In light of these results it appears that there is no detectable gas, down to levels of about 10$^{13}$ cm$^{-2}$, corotating in an extended gaseous disk at radii $> 35$ $\\alpha ^{-1}$. The cosmic web is the most likely origin for the detected Ly$\\alpha$ lines.  The observed equivalent widths of the lines are consistent with this picture.  Indeed the equivalent widths are correlated with the local volume density of galaxies around the sightline. This makes sense if the Ly$\\alpha$ lines arise from the cosmic web, which is denser in regions of higher volume density of galaxies (since the galaxies have formed by condensating in the denser parts of the network). This correlation would be difficult to explain if the Ly$\\alpha$ lines arised from halos or extended disks of the most nearby galaxy. One could argue that the equivalent widths should be larger for sighlines close to larger galaxies, which tend to be found in regions of higher galaxy volume density, which could then explain the trend observed.  However if this were the case then one would observe presumably an even stronger correlation between the equivalent widths and the luminosity of the nearby galaxy, and although this was reported to be the case by Chen \\etal (1998) it is now clearly contested by the results of Bowen \\etal (2002), as well as in our sample for which no correlation is found between equivalent widths and luminosities.   Note that these results do not necessarily imply that the dark matter haloes of these galaxies do not exist at the radii probed by our QSO sightlines. It simply means that there is no longer some detectable gas in rotation in an extended gas disk at these radii, and that the Ly$\\alpha$ lines detected are due to foreground or background cosmic web gas surrounding the galaxy. In fact weak lensing studies seem to indicate that indeed the dark matter halo extend very far out, typically $>$ 300 $h^{-1}_{75}$kpc for a $L_*$ galaxy (eg. the weak-lensing studies such as Smith \\etal 2001). But there is no gas to trace its dynamics out there, at least down to our limiting column densities. HI as traced by background QSOs is therefore not a useful tracer of galactic potential in the far outer regions of a $\\sim $200 kpc halo. Only weak lensing, and perhaps H$\\alpha$ in recombination from the extended ionized disk (Bland-Hawthorn \\etal 1997) are hopeful tracers of dynamics at such radii. These QSOs turn out however to be extremely useful to probe the phases of the cosmic web gas. Many more sightlines will be required though before one can attempt to understand the cosmic gas distribution and density, as well as its evolution with redshift. The planned HST COS spectrograph would have allowed to probe the sightlines of fainter QSOs which would greatly increase the available sample, and hopefully the decision regarding HST can be reversed in the future. Eventually the next generation of sensitive giant radio telescopes as the CLAR (C\\^ot\\'e \\etal 2002) and eventually the SKA, will be able to survey efficiently large areas of the sky down to colum densities as deep as 10$^{16}$ cm$^{-2}$, and hence trace directly in emission the structure of such low column density gas in the cosmic web."
    },
    "0410/astro-ph0410077_arXiv.txt": {
        "abstract": "Cadmium Zinc Telluride (CZT) detectors are having a major impact on the field of hard X-ray astronomy. Without the need for cryogenic cooling they achieve good spatial and energy resolutions over the broad energy range from 10 keV to $\\sim$600 keV. In this paper, we briefly review the historical development of detectors used in X-ray astronomy. Subsequently, we present an evaluation of CZT detectors from the company Imarad. The standard 2$\\times$2$\\times$0.5~cm detectors, contacted with 8$\\times$8 In pixels and an In cathode, exhibit FWHM energy resolutions of 7 keV  at 59 keV, and 10 keV at 662 keV. A direct measurement of the 662~keV photopeak efficiency gives 67\\%. We have started a detailed study of the performance of Imarad detectors depending on surface preparation, contact materials, contact deposition, post-deposition detector annealing, and detector passivation techniques. We present first results from contacting detectors with Cr, Ag, Au, and Pt. ",
        "introduction": "X-ray astronomy has come a long way since the first detection of X-ray emission from an extra-solar object in 1962\\cite{Giac:62}. In the meantime, X-ray astronomy has emerged as one of the drivers of modern astronomy, as evident from the award of the Nobel prize in physics to X-ray pioneer Riccardo Giacconi in 2002, and the existence of two major X-ray observatories, {\\it Chandra} and {\\it XMM-Newton}, orbiting Earth since 1999.  The most important results of X-ray astronomy have so far been obtained from observations in the ``soft'' X-ray band with photon energies between 0.1 keV and several keV. To mention several important milestones, the {\\it Einstein} observatory\\cite{Giac:79} achieved in 1978 a first breakthrough in the 0.2 keV to 4 keV band by combining a grazing incidence mirror with position sensitive detectors. The observatory carried a ``High Resolution Imager'' (HRI), consisting of two cascaded Micro Channel Plates (MCP) read out by a crossed-grid charge detector. Another main instrument was the Imaging Proportional Counter (IPC), a 4 cm deep, 800 Torr Ar-Xe-CO$_2$ proportional counter which achieved a modest energy resolution of between 50\\% and 100\\%. The grazing incidence mirror and the HRI achieved an effective detection area of 10 cm$^2$. The mirror and the IPC had a detection area of 100 cm$^2$. The {\\it Roentgen Satellite (ROSAT)} was launched in 1990 and used similar detectors as the {\\it Einstein} observatory. Its HRI and IPC cameras collected respectively 8 and 2.4 times more photons per unit time than their {\\it Einstein} counterparts\\cite{True:83}. The {\\it Chandra} and {\\it XMM-Newton} satellites are equipped with powerful X-ray grazing incidence mirrors and state-of-the-art photon detectors. The {\\it Chandra} satellite carries the {\\it Advanced CCD imaging spectrometer (ACIS)}\\cite{Garm:03} made of Si-CCDs with energy resolutions of between 2\\% and 10\\% over the energy range from 0.1 keV to 10 keV. The {\\it XMM-Newton} satellite carries two conventional Metal Oxide Semiconductor (MOS) cameras, and a ``PN'' camera\\cite{Stru:01}. The fully depleted $pn$ CCDs (280 micron deep depletion region) can detect X-rays with energies up to 15 keV. The {\\it Chandra} X-ray mirror and ACIS camera achieve a 1 keV detection area of 340 cm$^2$ and an unique angular resolution of  0.5''. The {\\it XMM-Newton} mirrors and CCD cameras have a combined 1 keV detection area of 2,100 cm$^2$ and an angular resolution of 15''.  Compared to the soft X-ray band, the hard X-ray band ($\\sim$ 20-200 keV) has progressed more slowly. Hard X-rays are more difficult to focus, scintillation detectors had only modest spatial and energy resolutions, and Ge detectors required resource intensive cryogenic cooling. The {\\it High Energy Astrophysics Observatory-1 (HEAO-1)} was equipped with a 1.4$^\\circ$ $\\times$ 20$^\\circ$ fan collimator and 250~cm$^2$ of NaI scintillator, and scanned the hard X-ray sky in the years 1977-1979 \\cite{Matt:78}. More recently, the {\\it Rossi X-Ray Timing Explorer (RXTE)}\\cite{Brad:90} launched in 1995 and {\\it Satellite per Astronomia X, ``Beppo'' (BeppoSAX)}\\cite{Piro:95} launched in 1996 used essentially the same detector technology with 1$^\\circ$ pencil beam collimators. The wide bandgap semiconductors Cadmium Telluride (CdTe) and Cadmium Zinc Telluride (CZT) are now boosting the field of hard X-ray astronomy. The {\\it International Gamma-Ray Astrophysics Laboratory (INTEGRAL)}\\cite{Uber:03} launched in 2002 carries the {\\it Imager on Board the Integral Satellite (IBIS)} that uses CdTe detectors with a detector area of $\\sim$2,600 cm$^2$. The {\\it SWIFT} satellite\\cite{Gehr:00}, to be launched in October 2004, carries the {\\it Burst Alert Telescope (BAT)} with a CZT detection area of $\\sim$5,240 cm$^2$. The {\\it IBIS} and the {\\it BAT} use the approach of coded mask imaging. Embedded in a collimator assembly, the detectors see the sky through a patterned shadow mask. The detectors image the sum of all shadows cast by the X-ray sources in the field of view. A deconvolution algorithm is used to derive the X-ray surface brightness distribution from the detected image. In this approach, the effective detection area is roughly half of the active detector area (the other half of the detector area is covered by the coded mask), and is not given by the collection area of any focusing optics.  The congressionally approved Beyond Einstein Program\\cite{BeyondEinstein} of NASA's ``Structure and Evolution of the Universe'' theme plans for two major X-ray missions. In both missions, CZT plays an important role. The Constellation-X observatory \\cite{conX} will perform high throughput X-ray spectroscopy with high energy and angular resolution over the energy range from 0.25 keV to 10 keV employing X-ray gratings with Si-CCDs, as well as non-dispersive microcalorimeters. The design includes a hard X-ray telescope that uses CZT detectors to cover the energy range from 6~keV to 40 keV or higher. Constellation-X is currently receiving NASA funds for Pre-Phase A development, the ``advanced concepts'' phase. Given the current budget plan, the mission will be launched in 2016\\cite{Sala:04}. A second recommended mission is the ``black hole finder probe'', a large field of view, hard X-ray telescope. The {\\it Energetic X-ray Imaging Survey Telescope (EXIST)} design\\cite{Grin:03} proposes an assembly of coded mask imagers. Hard X-rays are detected with $\\sim$0.5~cm thick CZT detectors with a total detector area of $\\sim$80,000 cm$^2$. {\\it EXIST} has been selected for a 2-year NASA concept study, along with the {\\it Coded Aperture Survey Telescope for Energetic Radiation (CASTER)}\\cite{McCo}, a competitor design which employs the new scintillators LaBr$_3$ or LaCl$_3$ as primary detectors. There is no definitive time schedule for the black hole finder probe, as its funding is presently beyond the budgetary horizon of 2009\\cite{Sala:04}.  In this paper, we emphasize thick CZT detectors and their application in {\\it EXIST}-type hard X-ray telescopes. In Sect.\\ \\ref{requirements} we will discuss key requirements for CZT detectors in this application. In Sect.\\ \\ref{standard} we scrutinize the performance of In contacted standard detectors from Imarad. Furthermore, we present detailed detector simulations and compare simulated data with experimentally measured data. In Sect.~\\ref{contacts} we give first results of contacting the detectors with Cr, Ag, Au, and Pt. Finally, in Sect.~\\ref{summary}, we summarize our findings and describe work in progress. ",
        "conclusions": "\\label{summary} CZT has emerged as the detector material of choice for the detection of 10 keV to 600 keV X-rays without the need for cryogenic cooling. CZT is having a major impact on the field of hard X-ray astronomy. In this paper we have presented measurements of Imarad CZT standard detectors. We have shown that the cathode amplitude is better suited to correct the anode amplitudes for the DOI than the duration of the cathode pulse. After correction for the DOI, the detectors achieve a photopeak efficiency of 67\\% at 662 keV. The 59 keV and 662 keV energy resolutions are 7 keV and 10 keV, respectively. We have described first results from contacting Imarad detectors with high workfunction metals. Although the detectors exhibit low leakage currents, they perform in most cases worse than standard Imarad detectors. Together with our finding that the standard Imarad detectors are not limited by shot noise from the leakage current, our study indicates that the main emphasis of the detector optimization should be on finding contacts that do not generate electronic noise themselves, rather than focusing on contacts that minimize the leakage current.  In a collaboration of Washington University, Fisk University, and UCSD (J.~Matteson, R.~T.~Skelton), we have started a systematic study of CZT surface preparation, contact materials, contact deposition, annealing, and passivation techniques. The aim is to improve on the low-energy performance of the detectors and to achieve a high yield of good detectors. \\begin{figure}[bt] \\begin{center} \\includegraphics[width=14cm]{czt.eps} \\end{center} \\vspace*{-0.8cm} \\caption{ \\label{sem} The left panel shows a secondary electron microscope (SEM) image of an Imarad CZT detector taken at Washington University. Light regions show ``heavy'' elemental composition, while dark regions show ``lighter'' elemental composition. The right hand panels show two energy spectra of backscattered electrons taken from small (several microns diameter) regions with widely different chemical composition. Contamination of the CZT surface with Al and Si polish residuals can be recognized.} \\end{figure} To steer our developments into the right direction, we are guiding the optimization process with several diagnostic tools:\\\\[1.5ex] {\\it Secondary Electron Microscope (SEM):} SEM allows us to map the chemical composition of the CZT surfaces. The energy spectrum of backscattered electrons from the light regions makes it possible to determine the relative Ca, Zi and Te compositions to an accuracy of 1\\%-2\\%. As an example, Fig.\\ \\ref{sem} shows a SEM picture of an Imarad detector after polishing; dark (light) regions correspond to a light (heavy) composition. The analysis of $\\sim$1 micron diameter dark regions shows that the substrate is heavily contaminated with Al and Si residuals from the polishing process, and should be cleaned more thoroughly.\\\\[1.5ex] {\\it X-ray diffraction:} We have been examining the CZT with X-ray diffraction for crystal structure and orientation. In the future, defect-induced broadening of the $\\theta$-$2\\theta$ peaks will also be studied. An example of an X-ray $\\theta$-$2\\theta$ diffraction scan from an Imarad CZT crystal is shown in Fig.\\ \\ref{xd}. The line structure shows that the crystal is oriented with the (111) direction perpendicular to the anode and cathode side of the detector. The diffraction peaks do not show the K-$\\alpha$1 and K-$\\alpha$2 doublet from the Cu source proving that line broadening owing to crystal irregularities is significant.\\\\[1.5ex] {\\it Photoluminescence, photoconductivity, and Pockels effect:} Photoluminescence measurements use a laser beam tuned above the band gap energy that is directed onto the CZT surface and the emitted radiation is recorded as a function of wavelength below the band gap energy. Radiative recombination from the band edge states as well as defect and impurity levels can be resolved. The technique is very sensitive to defect structure and impurities from the intrinsic material as well as from any subsequent surface preparation\\cite{Leop:86}. Photoconductivity measurements can be used to measure the surface recombination rate\\cite{Cui:02}. These measurements can be combined with internal field mapping using the Pockels effect\\cite{DeAn:96} and the response to alpha particles. We aim to decouple the charge loss contributions of the bulk ($\\mu-\\tau$ product) and the surface (owing to fabrication related recombination effects)."
    },
    "0410/astro-ph0410594_arXiv.txt": {
        "abstract": "Type Ia supernovae (SNe\\,Ia) currently provide the most direct evidence for an accelerating Universe and for the existence of an unknown ``dark energy'' driving this expansion. The 5-year Supernova Legacy Survey\\footnote{see \\texttt{http://cfht.hawaii.edu/SNLS/}} (SNLS) will deliver $\\sim$1000 SN\\,Ia detections with well-sampled $g'r'i'z'$ light curves.  Using this definitive dataset, we will obtain a precise measurement of the cosmological parameters ($\\Omega_{m}$, $\\Omega_{\\lambda}$); our goal is to determine the cosmological equation of state parameter ``$w$'' to a statistical precision of $\\pm0.07$, testing theories for the origin of the universal acceleration. In this proceeding, we summarise the progress made during the first full year of the survey operation. ",
        "introduction": "The Hubble diagram for high-redshift Type Ia supernovae (SNe\\,Ia) provides the most direct current measurement of the expansion history of the universe -- and hence the most direct evidence for an accelerating expansion.  The ``first generation'' of SN\\,Ia cosmology work developed a systematic approach to this measurement (Riess et~al. 1998; Perlmutter et~al. 1999) that led to astonishing results ruling out a flat, matter-dominated Universe. This indicated the presence of a new, unaccounted-for ``dark energy'' driving the cosmic acceleration.  One of the most pressing questions in cosmology now is: ``What is the dark energy that causes this acceleration?''. There are fundamental differences between a Cosmological Constant and other proposed forms of dark energy (see Peebles \\& Ratra 2003 for a comprehensive review). The distinction can be addressed by measuring the dark energy's average equation-of-state, $<$$w$$>$$=$$<$$p/\\rho$$>$, where $w$$=$$-1$ corresponds to a Cosmological Constant -- current measurements of this parameter (e.g., Knop et~al.  2003; Tonry et~al. 2003; Riess et~al.  2004) are consistent with a very wide range of dark energy theories.  The importance of improving measurements to the point where $w$$=$$-1$ could be excluded has led to a second-generation of SN cosmology studies: large multi-year, multi-observatory programs benefiting from major commitments of dedicated time.  These ``rolling searches'' find and follow SNe over many consecutive months of repeated wide-field imaging, with redshifts and SN type classification from coordinated spectroscopy.  The Canada-France-Hawaii Telescope (CFHT) Legacy Survey is one such ambitious repeat-imaging wide-field survey conducted in 4 SDSS filters ($g'r'i'z'$), utilising an imager field three times larger than used in the next largest survey, with around twice as much time devoted to the survey. The five-year CFHT ``Supernova Legacy Survey'' (SNLS) will provide the biggest improvement in the determination of the dark energy parameters achievable over the next decade, using an order-of-magnitude larger statistical sample (i.e.  $\\sim$700) of SNe in the redshift range $z$$=$0.3-0.9 where $w$ is best measured.  With this sample, we aim to answer the key question: Is the dark energy something other than Einstein's $\\Lambda$? ",
        "conclusions": ""
    },
    "0410/gr-qc0410024_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}     In many theories of physics, realistic situations are often difficult to describe by an exact solution of a theory, because theories in physics and their exact solutions are often too idealized to properly represent natural phenomena. Given this situation, we have to consider perturbative approaches to investigate realistic situations. General relativity is one theory in which the construction of exact solutions is not so easy. Though there are many exact solutions to the Einstein equation\\cite{Exact-solutions}, these are often too idealized. For this reason, general relativistic perturbation theory is a useful technique to investigate natural phenomena\\cite{Cosmological-Perturbations,Gerlach_Sengupta}.     In addition to this technical problem, general relativity is based on the concept of general covariance. Intuitively, the principle of general covariance states that there is no preferred coordinate system in nature, though the notion of general covariance is mathematically included in the definition of a spacetime manifold in a trivial way. This is based on the philosophy that coordinate systems are originally chosen by us, and natural phenomena have nothing to do with our coordinate system. Due to this general covariance, the {\\it gauge degree of freedom}, which is an unphysical degree of freedom of perturbations, arises in general relativistic perturbations. To obtain physically meaningful results, we have to fix this gauge degrees of freedom or to extract the {\\it gauge invariant part of perturbations}. A similar situation has been found in recent investigations of the oscillatory behavior of a gravitating Nambu-Goto membrane\\cite{kouchan-papers,kouchan-flat}, which concern the dynamical degrees of freedom of extended gravitating objects.     On the other hand, higher-order multi-parameter perturbations, in which there are two or more small parameters, can be applied to many physical situations. One well-known application of two-parameter perturbation theory is perturbations of a spherical star\\cite{Kojima}, in which we choose the gravitational field of the spherical star as the background spacetime for the perturbations, one of the parameters for the perturbations corresponds to the rotation of the star, and the other is its pulsation amplitude. The effects due to the rotation-pulsation coupling are described at higher orders. A similar perturbation theory on the Minkowski background spacetime was developed by the present author to study the comparison of the oscillatory behavior of a gravitating string with that of a test string\\cite{kouchan-flat}. Even in the one-parameter case, it is interesting to consider higher-order perturbations. In particular, Gleiser et al.\\cite{Gleiser-Nicasio} reported that second-order perturbations yield accurate wave forms of gravitational waves. Hence, it is worthwhile to investigate higher-order multi-parameter perturbation theory from a general point of view.     Motivated by these physical applications, the general relativistic gauge invariant multi-parameter perturbation theory has been developed in a number of papers\\cite{kouchan-flat,kouchan-gauge-inv,Bruni-Gualtieri-Sopuerta}. In particular, the procedure to find gauge invariant variables for higher-order perturbations on a generic background spacetime was proposed by the present author\\cite{kouchan-gauge-inv} assuming that we have already known the procedure to find gauge invariant variables for a linear-order metric perturbations. The contents of this paper are based on this proposal. The main purpose of this paper is to present some formulae of the second-order perturbative curvatures within the two-parameter perturbation theory that are useful in some physical applications. When we derive these formulae, we follow the general framework of the second-order gauge invariant perturbation theory on an arbitrary background spacetime. This framework is originally proposed by Stewart et al.\\cite{J.M.Stewart-M.Walker11974} and developed by Bruni et al.\\cite{Bruni-Gualtieri-Sopuerta,M.Bruni-S.Soonego-CQG1997}, and the present author\\cite{kouchan-gauge-inv}. These perturbative curvatures have the same form as those given in the definitions of gauge invariant variables for arbitrary perturbative fields which are proposed in a previous paper\\cite{kouchan-gauge-inv}. As in that paper, we do not make any specific assumption regarding the background spacetime and the physical meaning of the two-parameter family. Because we make no assumption concerning the background spacetime, this framework has a wide area of applications.     The organization of this paper is as follows. In \\S\\ref{sec:gauge-freedom-of-perturbations}, we review the general framework of the second order gauge invariant perturbation theory. We mainly review the one-parameter perturbation theory. We emphasize that the review in \\S\\ref{sec:gauge-freedom-of-perturbations} of this paper is based on the idea of Stewart et al.\\cite{J.M.Stewart-M.Walker11974} and the development carried out by Bruni et al.\\cite{Bruni-Gualtieri-Sopuerta,M.Bruni-S.Soonego-CQG1997}. In \\S\\ref{sec:Formulae-of-perturbative-curvatures}, we present some formulae for the second-order perturbative curvatures within the two-parameter perturbation theory. We also give a derivation of these formulae. The final section, \\S\\ref{sec:summary}, is devoted to a summary and brief discussion of physical situations to which this framework of higher-order perturbation theory should be applied. We employ the notation of our previous paper\\cite{kouchan-gauge-inv} and use the abstract index notation\\cite{Wald-book}. ",
        "conclusions": "\\label{sec:summary}     In summary, we have derived some formulae for second-order gauge invariant perturbations, namely those of the Riemann, Ricci, scalar, Einstein, and Weyl curvature tensors. We also derived the formulae for the divergence of an arbitrary tensor field of second rank for perturbations at each order. These perturbative curvatures have the same forms as those in the definitions (\\ref{eq:matter-gauge-inv-def-1.0})--(\\ref{eq:matter-gauge-inv-def-1.1}) of the gauge invariant variables for arbitrary perturbative fields which are proposed in a previous paper\\cite{kouchan-gauge-inv}. These are useful for investigating physical problems using perturbation theory.     Through the derivation of these formulae, we have confirmed some facts. First, if linear-order gauge invariant perturbation theory is well established, its extension to higher orders and multi-parameter perturbations is straightforward. Second, the perturbative Weyl curvature at each order preserve the property of the tracelessness of the Weyl curvature. Third, the perturbative Einstein curvature of each order satisfy the perturbative Bianchi identities at each order. Fourth, the perturbative Einstein equations at each order necessarily take in gauge invariant forms. These properties we have confirmed are trivial from their derivations. In particular, the fourth result is trivial, because any equation can be written in a form in which the right-hand side is equal to ``zero'' in any gauge. This ``zero'' is gauge invariant. Hence, the left-hand side should be gauge invariant. However, we emphasize that these trivial results imply that the formulae derived here and the framework of higher-order perturbation theory applied here are mathematically consistent at this level.     Further, we note that in our framework, we specify nothing about the background spacetime nor about the physical meaning of the parameters for the perturbations. Our framework is based only on general covariance. For this reason, this framework is applicable to any theory in which general covariance is imposed, and thus it has very many applications. Actually, we are planning to apply it to some physical problems. The following are candidates of the physical situations to which the second order perturbation theory should be applied: the radiation reaction problem in gravitational wave emission\\cite{radiation-reaction}; stationary axisymmetric ideal MHD flow around a black hole or a star\\cite{MHD-ref}; the correspondence between observables in experiment and gauge invariant variables; dynamics of gravitating membranes (for example, topological defects\\cite{kouchan-papers}, brane world\\cite{brane-papers}, and so on); perturbations of a compact star with rotation and pulsation\\cite{Kojima}; Post-Minkowski expansion alternative to post-Newtonian expansion\\cite{Post-Minkowski-Detweiler}; higher-order cosmological perturbations and primordial non-Gaussianity\\cite{cosmological-second-non-gauss}.     In particular, the gauge invariant form of the second-order perturbation of the divergence of the energy momentum tensor should be useful in considering the gauge problem in the radiation reaction of gravitational wave emission. In astrophysical contexts, it is natural to consider the situation in which a solar mass object falls into a supermassive black hole of mass $\\sim 10^{6}M_{\\odot}$. This is one of the target phenomena of the observation of gravitational wave by LISA (laser interferometer space antenna for gravitational wave measurements)\\cite{LISA-home-page}. In such a situation, the perturbation parameter is the ratio of the mass of the compact object to that of the central supermassive black hole. We may regard the second order perturbations as describing the radiation reaction effect of gravitational wave emission. By applying the gauge invariant formulation discussed here, we can exclude gauge freedom completely and thus there is no gauge ambiguities in the results. It would be quite interesting to apply the formulation discussed here to this radiation reaction problem. We leave this application as a future work.     In addition to the radiation reaction problem of gravitational wave emission, there are many applications, some of which are listed above. Indeed, because there are few assumptions in our treatment, it is natural to expect that there are many applications that are not listed above. We also expect that the formulae derived here will be found to be very powerful tools in many applications."
    },
    "0410/astro-ph0410580_arXiv.txt": {
        "abstract": "The main astrophysical factories of fluorine ($^{19}$F) are thought to be Type~II supernovae, Wolf--Rayet stars, and the asymptotic giant branch (AGB) of intermediate mass stars. We present a model for the chemical evolution of fluorine in the Milky Way using a semi-analytic multi-zone chemical evolution model. For the first time, we demonstrate quantitatively the impact of fluorine nucleosynthesis in Wolf--Rayet and AGB stars. The inclusion of these latter two fluorine production sites provides a possible solution to the long-standing discrepancy between model predictions and the fluorine abundances observed in Milky Way giants. Finally, fluorine is discussed as a possible probe of the role of supernovae and intermediate mass stars in the chemical evolution history of the globular cluster $\\omega$~Centauri. ",
        "introduction": "The three primary astrophysical factories for fluorine ($^{19}$F) production have long been thought to be Type~II Supernovae (SNe~II), Wolf--Rayet (WR) stars, and asymptotic giant branch (AGB) stars (e.g. \\citealt{WW}; \\citealt{MA}; \\citealt{ForestiniETal}; Mowlavi, Jorissen \\& Arnould 1998, respectively). Previous attempts to model the Galactic production and evolution of $^{19}$F have been restricted to explore the role of SNe~II alone (e.g. Timmes, Woosley \\& Weaver 1995; Alib\\'es, Labay \\& Canal 2001).  The above problem has now been ameliorated by the release of the first detailed yield predictions for fluorine production from WR and AGB stars. We are now in a position to incorporate these yields into a Galactic chemical evolution framework, in order to assess the respective contributions of the three putative fluorine production sites. To do so, we will make use of {\\tt GEtool}, a semi-analytical multi-zone Galactic chemical evolution package which has been calibrated with extant observational data for the Milky Way (\\citealt{FG}; \\citealt{GibsonETal}).  Specifically, in what follows, we compare the model fluorine distribution in the Milky Way with the abundances observed by Jorissen, Smith \\& Lambert (1992) in near-solar metallicity giants. Further, our model predictions are contrasted with new fluorine determinations for giants in the Large Magellanic Cloud (LMC) and $\\omega$~Centauri \\citep{CunhaETal}. In addition, new results for more $\\omega$ Centauri giants from \\cite{SmithETal} are included. The latter two systems are likely to have had {\\it very} different star formation and chemical evolution histories from those of the Milky Way, but despite these obvious differences, a comparison against these new data can be valuable. In Section~2, we provide a cursory overview of the three traditional $^{19}$F nucleosynthesis sites; the chemical evolution code in which the nucleosynthesis products from these factories have been implemented is described in Section~3. Our results are then presented and summarised in Sections~4 and 5, respectively.  \\section[]{Nucleosynthesis of $^{19}$F} \\subsection{Type~II Supernovae}  The massive star progenitors to SNe~II produce fluorine primarily as the result of spallation of $^{20}$Ne by $\\mu$ and $\\tau$ neutrinos near the collapsed core (\\citealt{WH}; \\citealt{WoosleyETal}). A fraction of the $^{19}$F thus created is destroyed by the subsequent shock but most is returned to the ambient Interstellar Medium (ISM). The fluorine yields by neutrino spallation are very sensitive to the assumed spectra of $\\mu$ and $\\tau$ neutrinos \\citep{WHW}, which could be nonthermal and deficient on their high-energy tails, lowering the equivalent temperature of the neutrinos in the supernova model \\citep{MB}. An additional source of $^{19}$F derives from pre-explosive CNO burning in helium shell. However, fluorine production by neutrino spallation is largely dominant, as evident by comparing the models in \\cite{WW}, the only ones to-date including neutrino process, and recent models which do not include neutrino nucleosynthesis of fluorine \\citep{LC}. Most recently, \\cite{HegerETal} suggest that the relevant neutrino cross sections need to be revised downwards; if confirmed, the associated SNe~II $^{19}$F yield would decrease by $\\sim$~50\\%.  In light of the preliminary nature of the Heger et~al. claim, we retain the conservative choice offered by the \\cite{WW} compilation.  \\subsection{Asymptotic giant branch stars}  The nucleosynthesis pathways for fluorine production within AGB stars involve both helium burning and combined hydrogen-helium burning phases (e.g. \\citealt{ForestiniETal}; \\citealt{JorissenETal}; \\citealt{MowlaviETal}) and are companions for the nucleosynthesis by slow neutron accretion (s-process) \\citep{MowlaviETal}. Provided a suitable source of protons is available, fluorine can be synthesised via $^{14}$N($\\alpha$,$\\gamma$)$^{18}$F\\,($\\beta^{+}$)\\, $^{18}$O(p,$\\alpha$)\\,$^{15}$N($\\alpha$,$\\gamma$)\\,$^{19}$F. Primary sources of uncertainty in predicting fluorine nucleosynthesis in AGB stars relate to the adopted reaction rates, especially $^{14}$C$(\\alpha,\\gamma)^{18}$O and $^{19}$F$(\\alpha,p)^{22}$Ne, and the treatment of the nucleosynthesis occurring during the convective thermal pulses. Nucleosynthesis during the interpulse periods can also be important if protons from the envelope are partially mixed in the top layers of the He intershell (partial mixing zone), as \\cite{LugaroETal} have recently demonstrated. Nucleosynthesis in this zone may result in a significant increase in the predicted $^{19}$F yields. The magnitude of these systematic uncertainties for stellar models with mass $\\sim~3$~M$_{\\odot}$ and metallicities Z = $0.004 - 0.02$ are $\\sim~50\\%$, while for stellar models with mass M = $5$~M$_{\\odot}$ and metallicity Z = $0.02$ the uncertainty is a factor of $\\sim~5$, due to the uncertain $^{19}$F$(\\alpha,p)^{22}$Ne reaction rate. Characterising the mass- and metallicity-dependence of the partial mixing zone--$^{19}$F relationship needs to be completed before we can assess its behaviour self-consistently within our chemical evolution model of the Milky Way. For the present study, we have adopted the yields presented in Appendix, based upon the Karakas \\& Lattanzio (2003, and references therein) models, which themselves do not include $^{19}$F nucleosynthesis via partial mixing. This choice is a conservative one, and thus should be considered as a lower limit to the production of $^{19}$F from AGB stars.  For stars more massive than $\\approx 4$~M$_{\\odot}$, the convective envelope is so deep that it penetrates into the top of the hydrogen-burning shell so that nucleosynthesis actually occurs in the envelope of the star. Such ``hot-bottom-burning'' acts to destroy $^{19}$F, and should be treated self-consistently within the AGB models considered.  \\subsection{Wolf--Rayet stars}  Fluorine production in WR stars is tied to its nucleosynthesis during the helium-burning phase. At the end of this phase though, significant fluorine destruction occurs via $^{19}$F$(\\alpha,p)^{22}$Ne. Any earlier synthesised $^{19}$F must be removed from the stellar interior in order to avoid destruction. For massive stars to be significant contributors to net fluorine production, they must experience mass loss on a timescale that allows the removal of $^{19}$F before its destruction. This requirement is met by WR stars.  Recently, \\cite{MA} studied the role that such stars can play in the chemical evolution of fluorine by adopting updated reaction rates coupled with extreme mass-loss rates in not-rotating stellar models. They pointed out that WR mass-loss is strongly metallicity-dependent, and that the number of WR stars at low metallicities is very small. Their WR yields reflect such metallicity-dependence, with minimal fluorine returned to the ISM at low metallicities, but significant $^{19}$F returned at solar and super-solar metallicities. The WR yields are sensitive to the adopted reaction and mass-loss rates, while rotating models could favour an early entrance into the WR phase for a given mass, decrease the minimum initial mass for a star to go through a WR phase at a given metallicity, and open more nucleosynthetic channels because of the mixing induced by rotation. Therefore, after \\cite{MA}, we consider the aforementioned WR yields as lower limits.  \\section[]{The model}  In this study we employ {\\tt GEtool}, our semi-analytical multi-zone chemical evolution package to model a sample Milky Way-like disk galaxy (\\citealt{FG}; \\citealt{GibsonETal}). A dual--infall framework is constructed in which the first infall episode corresponds to the formation of the halo, and the second to the inside-out formation of the disk.  A Kroupa, Tout \\& Gilmore (1993) initial mass function (IMF) has been assumed, with lower and upper mass limits of 0.08~M$_{\\odot}$ and 120~M$_{\\odot}$, respectively. Stellar yields are one of the most important features in galactic chemical evolution models, yet questions remain concerning the precise composition of stellar ejecta, due to the uncertain role played by processes including mass loss, rotation, fall-back, and the location of the mass cut, which separates the remnant from the ejected material in SNe. The SNe~II yields are from \\cite{WW}; the yields for stars more massive than 60~M$_{\\odot}$ are assumed to be mass-independent. Such assumption is made to avoid extreme extrapolation from the most massive star in the Woosley \\& Weaver models (40~M$_{\\odot}$) to the upper end of the IMF (120~M$_{\\odot}$), and has negligible effect on the results, given the shape of the adopted IMF.  We have halved the iron yields shown in \\cite{WW}, as suggested by \\cite{TimmesETal}. The Type~Ia (SNe~Ia) yields of \\cite{IwamotoETal} were also employed. We adopted the metallicity-dependent yields of \\cite{RV} for single stars in the mass range $1 - 8$~M$_{\\odot}$. For the purposes of this work, which focuses on fluorine, the choice of the Renzini \\& Voli yield set does not affect the results. Metalllicity-dependent stellar lifetimes have been taken from \\cite{SchallerETal}.  We have constructed three Milky Way (MW) model variants that differ only in their respective treatments of $^{19}$F production:  1) MWa assumes that SNe~II are the only source of $^{19}$F; 2) MWb includes yields from both SNe~II {\\it and} WR stars; 3) MWc includes all three sources of fluorine - SNe~II, WR, and AGB stars.  We end by noting that within our adopted dual--infall framework for the chemical evolution of the Milky Way, our model is constrained by an array of observational boundary conditions, including the present-day star and gas distributions (both in density and metallicity), abundance ratio patterns, age-metallicity relation, and G-dwarf distribution \\citep{GibsonETal}.  While the modification of any individual ingredient within model framework will have an impact, to some degree, upon the predicted chemical evolution, this can only eventuate at the expense of one or more of the aforementioned boundary conditions that we require our model to adhere to. Within our framework, yield uncertainties will dominate the systematic uncertainties for the predicted evolution of $^{19}$F.   \\subsection{Fluorine yields} \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{rfgFig1} \\caption{Fluorine yields from a) SNe~II (Woosley \\& Weaver 1995), b) WR (Meynet \\& Arnould 2000), and c) AGB stars (Appendix).} \\label{fig1} \\end{center} \\end{figure}  We now summarise the $^{19}$F yields employed in our three ``Milky Way'' models.  1) SNe~II $^{19}$F yields are taken from \\cite{WW} and assumed to be mass-independent for stellar masses in excess of 60~M$_{\\odot}$.  2) WR $^{19}$F yields are taken from \\cite{MA} for stellar masses in the range $25 - 120$~M$_{\\odot}$: each star within this range is assumed to evolve through the WR stage. Such simplifying assumption could overestimate the WR contribution to fluorine, even though the adopted WR yields are themselves lower limits (Section~2.3). The WR fluorine contribution has been added to the corresponding SNe~II contribution (which comes from a different stage of the stellar evolution).  3) AGB $^{19}$F and oxygen yields in the $1 - 6.5$~M$_{\\odot}$ mass range have been derived from stellar models constructed with the Mount Stromlo Stellar Structure Code (\\citealt{FL}; \\citealt{KarakasETal}), and are presented in Appendix. The post-processing nucleosynthesis models with 74 species and time-dependent diffusive convective mixing are described in detail in \\cite{FrostETal} and \\cite{KL}.  To ensure internal consistency, we have also employed the AGB oxygen yields \\it in lieu \\rm of those of Renzini \\& Voli (1981), within this mass range.  The above fluorine yields are shown in Figure~1. In Figure~2, the yields are expressed as [F/O]\\footnote{Hereafter, [X/Y]$=log_{10}$(X/Y)$-log_{10}$(X/Y)$_{\\odot}$ and A(X)$=12+log_{10}$(n$_{\\rm X}$/n$_{\\rm H}$). An accurate determination of photospheric solar abundances requires detailed modeling of the solar granulation and accounting for departures from local thermodynamical equilibrium (e.g. Allende Prieto, Lambert \\& Asplund 2001). We adopt the solar fluorine abundance suggested by Cunha et~al. (2003), and the solar iron and oxygen abundances from Holweger (2001).} and $\\langle$[F/O]$\\rangle_{IMF}$, the latter corresponding to the mean [F/O] yields for SNe~II and AGB stars, weighted by the IMF over the SNe~II and AGB mass range, respectively. We have not shown a comparable entry for the WR stars as a self-consistent treatment of the oxygen production was not included in Meynet \\& Arnould (2000). Here, oxygen has been used as the normalisation to make easier the comparison with the observations, especially in $\\omega$ Centauri, though oxygen can be synthesised in various stellar sites, and its yields can be affected by different reaction rates and modeling of helium cores, semi-convection, convective boundary layers, and mass-loss (e.g. \\citealt{WHW}; \\citealt{DrayETal}).  \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{rfgFig2} \\caption{[F/O] and $\\langle$[F/O]$\\rangle_{IMF}$ for SNe~II and AGB yields (upper and lower panels, respectively). Here, A($^{19}$F)$_{\\odot} = 4.55$ (see discussion in Cunha et al. 2003) and A(O)$_{\\odot} = 8.736$ (e.g. Holweger 2001). The shaded regions in Figures~2a and 2b show the observed [F/O] in $\\omega$~Cen giants (Cunha et al. 2003). The $\\langle$[F/O]$\\rangle_{IMF}$ are weighted by the IMF over the SNe~II ($11 - 40$M$_{\\odot}$) and AGB ($1 - 6.5$~M$_{\\odot}$) mass range, respectively. In Figure~2d, both $\\langle$[F/O]$_{AGB}\\rangle_{IMF}$ and $\\langle$[F/O]$_{SNe~II}\\rangle_{IMF}$ are shown (closed boxes and open triangles, respectively).} \\label{fig2} \\end{center} \\end{figure} ",
        "conclusions": "We have studied the Galactic chemical evolution of fluorine, for the first time using new grids of stellar models which provide self-consistent predictions of fluorine nucleosynthesis for stars in both the WR and AGB phases of stellar evolution. We have shown that the WR contribution is significant at solar and super-solar metallicities because of the adopted metallicity-dependent mass-loss prescription employed in the stellar models. In contrast, the contribution of AGB stars to fluorine production peaks during the early epochs of the Galaxy's evolution (again due to the metallicity-dependent behaviour of the AGB models). In combination, the addition of the WR and AGB yields leads to a significant improvement in the galactic chemical evolution models when compared against observations.  The comparison between our MW models and the fluorine abundances in LMC and $\\omega$~Cen giants \\citep{CunhaETal} is not straightforward, as the latter two have star formation (and therefore chemical evolution) histories different from that of the MW. However, it is interesting to speculate on the possible origin of fluorine in $\\omega$~Cen, given the unique nature of this ``globular cluster'' (e.g. \\citealt{S}). Specifically, $\\omega$~Cen is the most massive Galactic cluster, and unlike most globulars, possesses a significant spread in metallicity ($\\sim$~1.5~dex) amongst its stellar population. It has been suggested that $\\omega$~Cen is actually the remnant core of a tidally-disrupted dwarf galaxy \\citep{BF}. Such a scenario could naturally drive radial gas inflows to the dwarf nucleus, potentially triggering starbursts.  Interestingly, SNe~II ejecta are characterised by low $\\langle$[F/O]$\\rangle_{IMF}$ (Figure~2c) and high $\\langle$[O/Fe]$\\rangle_{IMF}$ (Figure~4b), whereas AGB ejecta have higher $\\langle$[F/O]$\\rangle_{IMF}$ (Figure~2d). The observed $\\omega$~Cen giants have primarily low [F/O] (Figures~2a -- 2b, 4a -- 4b) and high [O/Fe] (Figure~4b) values, consistent with a picture in which their interiors have been polluted by the ejecta of an earlier generation of SNe~II, but not from a comparable generation of AGB. Given that the observed oxygen abundance in such $\\omega$~Cen giants is similar to that seen in comparable LMC and MW giants (Figure~4a), this would suggest that the chemical enrichment of $\\omega$~Cen proceeded on a short timescale (to avoid pollution from the lower mass progenitors to the AGB stars) and in an inhomogeneous manner (given the significant scatter in observed fluorine abundances), as discussed previously by \\cite{CunhaETal}. Should the (downward) revised neutrino cross sections alluded to in Section~2.1 be confirmed (Heger et~al. 2004), the concurrent factor of $\\sim$~2 reduction in SNe~II $^{19}$F production would improve the agreement of the model with the observed F/O ratio in $\\omega$~Cen giants.  This would consequently strengthen our conclusions which already support a picture whereby these giants have been polluted by earlier generations of SNe~II ejecta."
    },
    "0410/astro-ph0410063_arXiv.txt": {
        "abstract": "A reconstruction method for recovering the initial conditions of the Universe starting from the present galaxy distribution is presented which guarantees uniqueness of solutions. We show how our method can be used to obtain the peculiar velocities of a large number of galaxies, hence trace galaxies orbits back in time and obtain the entire past dynamical history of the Universe above scales where multi-streaming has not occurred. When tested against a $128^3$ $\\Lambda$CDM simulation in a box of $200$h$^{-1}$ Mpc length, we obtain $60\\%$ exact reconstruction on scales above $6$ h$^{-1}$Mpc. We apply our method to a real galaxy redshift catalogue, the updated NBG (Nearby Galaxies), containing 1483 galaxies, groups and clusters in a radius of $30$ Mpc/h, and reconstruct the peculiar velocity fields in the local neighbourhood. Our reconstructed distances are well-matched to the observed values outside the collapsed regions if $\\Omega_m(t)= 0.20\\, \\exp(-0.26 (t-13))$ where $t$ is the age of the Universe in Gyrs. ",
        "introduction": "Reconstruction of the initial condition of the Universe from the present distribution of the galaxies, brought to us by ever-more sophisticated redshift surveys, is an instance of the general class of {\\it inverse problems} in physics. In cosmology this problem is frequently tackled in an empirical way by a {\\it forward approach}. A {\\it statistical} comparison between the outcome of an N-body simulation and the observational data is made, assuming that a suitable {\\it bias} relation exists between the distribution of galaxies and that of dark matter. If the statistical test is satisfactory then the implication is that the initial condition assumed by the simulation is a viable one for our Universe, otherwise one changes the cosmological parameters until a statistical convergence between the observed and the simulated present Universe  is achieved.  Since Newtonian gravity is time-reversible, one could integrate the equations of motions back in time and solve the reconstruction problem trivially, if in addition to their positions, the present velocities of the galaxies were also known. However, the peculiar velocities of only a few thousands of galaxies are known out of thousands whose redshifts have been measured. Thus, a second boundary condition, in addition to the present redshifts of the galaxies, has to be provided : as we go back in time the peculiar velocities of the galaxies vanish. Contrary to the forward approach where one solves an {\\it initial-value problem}, in the reconstruction approach one is dealing with a {\\it two-point boundary value problem} (In this case, only the functional dependence of one of the boundary conditions is given, namely: ${\\rm time}\\rightarrow 0$ then ${\\rm peculiar\\,\\, velocities}\\rightarrow 0$). In the former, one has a {\\it unique} solution but in the latter this is not always the case.  The question remains whether unique reconstruction can be achieved. In this work, we report on a new method of reconstruction (Frisch et al. 2002, Mohayaee et al. 2003, Brenier et al. 2003) which guarantees uniqueness. ",
        "conclusions": ""
    },
    "0410/astro-ph0410549_arXiv.txt": {
        "abstract": "We perform a dynamical analysis of the recently updated set of  the radial velocity (RV) measurements of the HD~160691 ($\\mu$~Arae). The  purely kinematic, 2-Keplerian model of the measurements leads to the best-fit solution in which the eccentricity of the outer planet is about 0.7 and its semi-major axis is about 4~AU. The parameters of the inner planet are  well determined. The eccentricity is about~0.3 and the semi-major axis is about 1.65~AU. In such circumstances, the best 2-Keplerian model leads to a catastrophically unstable configuration, which is self disrupting in less than 20,000~yr. In order to derive dynamically stable configurations which are simultaneously consistent with the RV data, we use the so called GAMP (Genetic Algorithm with MEGNO Penalty) which incorporates the genetic fitting algorithms  and a verification of the stability by the fast indicator of the dynamics  (MEGNO). Using this method, we derive meaningful limits on the parameters of  the outer planet which also provide a stable behavior of the system. It appears,  that the best-fit solutions are located in a shallow valley of $\\Chi$, in the  $(a_{\\idm{c}},e_{\\idm{c}})$-plane,  extending over 2~AU (for the formal $1\\sigma$ confidence interval of the best fit). We find two equally good best-fit solutions leading to the qualitatively  different orbital configurations.  One of them corresponds to the center of the 5:1 mean motion resonance (MMR) and the second one describes a configuration  between the 6:1 and 17:2~MMRs.  Generally, the \\pstar system can be found in  the zone confined to other low-order MMRs of the type $p$:1 with $p>5$.   Our results support the classification of the \\pstara as a hierarchical   planetary system, dynamically similar to other known multi-planet systems  around HD~12661 and HD~169830. The results of the GAMP analysis on the extended  data set are fully consistent with our previous conclusions concerning a much  shorter observational window. It constitutes a valuable confirmation that the applied method is well suited for the analysis of the RV data series  which only partially cover the longest orbital period. ",
        "introduction": "Precise measurements of the radial velocity (RV) of two stars, HD~160691 (\\pstara) and HD~154857, from the Anglo-Australian survey for extrasolar planetary systems have been recently published by~\\cite{McCarthy2004}. Our particular attention is devoted to the \\pstar system.  \\cite{Butler2001} reported a planet with a 2~yr orbital period about the \\pstara.  This discovery was confirmed by \\cite{Jones2002b} and the observational window extended to 4~yr revealed an additional linear trend in the RV signal, interpreted as the signature  of a more distant planet. The data presented in the new paper by the Anglo-Australian  Planet Search team cover about 6~yr. This extended observational window allows  to confirm the existence of the second Jovian planet in the system.  Very recently,  \\cite{Santos2004} discovered a small, $14~$-Earth masses object in a circular orbit,  at the distance of only 0.09~AU from the \\pstara star and confirmed the findings of \\citep{McCarthy2004}. The Anglo-Australian team determined the  Keplerian elements of the Jovian companions: the orbital periods of 645~d, 3000~d  and the eccentricities of 0.2, 0.57, respectively \\citep{McCarthy2004}.  According to these authors, the current data set covers about 70\\% of the  longer period.  The detection of multiple, Jovian-like planets in stages from shorter to longer orbital periods are obviously expected from the Doppler surveys and \\citep{McCarthy2004} announce another possible multiple system around HD~154857. It is interesting to follow the subsequent estimations of the orbital parameters of such system, derived on the basis of an increasing number of observations. In the preprint by \\cite{Jones2002a}, the preliminary orbital solution for \\pstara suggested a proximity of the two putative planets to the 2:1~mean motion resonance (MMR).  Such configuration would be the third occurrence of this particular low-order resonance among only about 15 multi-planet systems known at that time. However, the formal 2-Kepler solution found  in the preprint appeared to be very unstable. In the published paper, \\cite{Jones2002b} excluded the possibility of the 2:1~MMR due to an increased orbital period of the outer planet (by about 300~d). Our independent analysis of the same RV data set \\citep{Gozdziewski2003e} has revealed that in the specific case of \\pstara, when one has only a part of the orbital RV signal of the additional planetary companion, the application of 2-Kepler  model of RV signal certainly leads to artifacts. The formally best-fit configurations are self-disrupting on a short time-scale! The newest set of measurements leads to the 2-Kepler solution with the outer planets' period increased to $\\simeq 3000$~d~\\citep{McCarthy2004}. This 2-Keplerian fit {\\rm still} describes a configuration which disintegrates after only $20,000$~yr.  It can be easily checked by direct integration of the 3-body problem. In fact even a simple analysis shows that this best fit is located  in the proximity of the planetary collision zone. The exact collision curve  can be determined through $a_{\\idm{b}} ( 1 + e_{\\idm{b}} ) = a_{\\idm{c}} ( 1 - e_{\\idm{c}} )$.  This means that the large eccentricity of the companion~c would be possible only if an orbital dynamical mechanism protecting the planets from an encounter existed. Another possibility is that the real orbital elements are in fact different from those determined by \\cite{McCarthy2004}.  Clearly, one needs a method of dynamical analysis of the orbital fits which will  reproduce an observed RV signal and simultaneously fulfill stability criteria. In our method, we not only account for the mutual interactions between the planets~\\citep{Laughlin2001} but also eliminate the orbital fits corresponding to strongly chaotic orbits, regardless of their goodness of fit expressed by $\\Chi$. We treat the dynamical behavior in terms of the chaotic and regular (or weakly-chaotic) states as an observable, at the same level of importance as the RV measurements are. Details of the algorithm which will be called from hereafter the Genetic Algorithm with MEGNO Penalty (GAMP) are presented in \\citep{Gozdziewski2003e}. This method seems to be particularly useful for systems undergoing strong mutual interactions such as the \\pstara system. According to the quoted preliminary estimates of the semi-major axes of the companions and their eccentricities, the likely configurations are related to the zone of strong low-order mean motion resonances like 5:2, 4:1, 5:1, 6:1, 7:1. The widths of these resonances rapidly increase with the growing eccentricity of the outer companion, finally  leading to their overlapping and a creation of a zone of global instability (see section~4 in this work). In this instance, the MMRs-originating chaotic diffusion leads to violent, macroscopic changes in the physical state of the system \\citep{Froeschle1999}. In a planetary system, it means close encounters between the planets and/or the parent star, collisions or ejection of a body from the system. Following this reasoning, to find configurations  which are stable over long periods of time, it is essential to eliminate the strongly chaotic solutions. These solutions originate in the unstable MMRs and lead to  a discontinuity of the searched parameter space, in the sense of the dynamical stability. The sophisticated structure of the phase space  can be understood in terms of the KAM theorem. Obviously, both the kinematic models of the Doppler signal as well as the gradient-like methods of the RV data fitting which require the $\\Chi$ to be a continuous function of its variables, are not well suited to this case. On the contrary, the discontinuous structure of the orbital parameter space justifies the application of non-gradient fitting and Monte-Carlo based techniques.  Taking into account these very basic considerations, we analyzed the RV data published by \\cite{Jones2002b}. We found that the quasi-linear trend present in the reflex motion of \\pstar permitted a continuum of equally good solutions in the $(a_{\\idm{c}},e_{\\idm{c}})$-plane. Moreover, using the GAMP algorithm  and assuming that $a_{\\idm{c}}<5$~AU,  it was possible to reduce the space of stable fits, permitted by the RV data, to a number of solutions confined to the zone of low-order mean motion resonances (MMRs) like 3:1, 7:2, 4:1, 9:2, 5:1, 11:2, 6:1. Moreover, the short observational window did not permit a distinction between them. At present, having the access to the updated observations, we have an occasion to verify these early predictions and to test  the reliability of the fits obtained with GAMP. Many stars accompanied by Jovian planets reveal linear trends in the Doppler signal indicating the presence of additional companions~\\cite[e.g.,][]{Butler2003}. Hence, we have a chance to test if our method is helpful in obtaining preliminary and realistic estimates of their orbital parameters. ",
        "conclusions": "The new RV data of the \\pstar system published by the Anglo-Australian Planet Search team \\cite{McCarthy2004} enable us to refine the orbital best-fit solutions. The preliminary observations published in~\\cite{Jones2002b} suggested the presence of the second companion inducing a linear trend in the RV signal. In  the light of the currently available data, it is very likely that the semi-major axis of the outer planet cannot be smaller than about $4$~AU. This makes the system relatively separated. Curiously, the kinematic 2-Keplerian model of the  RV signal constantly produces artifacts e.g. the apparent proximity to the 2:1~MMR  or large eccentricities of the outer companion. These configurations lead to  collisions between the planets in a very short time. In this work, we applied  the \"dynamical\" fitting by forcing a requirement that a real system should not be disrupted over a short time-span (thousands of years). Our analysis with the proposed GAMP algorithm reveals that the current data set still does not well constrain the semi-major axis of the outer companion. But the requirement of dynamical stability puts strong limits on the initial eccentricity of the outermost planet. The current set of observations is equally well modeled by a continuum of statistically equivalent best-fit solution which are distributed over $a_{\\idm{c}} \\in [4,6)$~AU and $e_{\\idm{c}}$ about $[0-0.4]$, along a flat valley of $\\Chi$  in the $(a_{\\idm{c}},e_{\\idm{c}})$-plane of osculating orbital parameters. Surprisingly, the initial longitudes of the coplanar configurations seem to be already well fixed.  The dynamical analysis of the best-fit solutions reveals that the currently determined parameters of the outer planet are confined to the zone of low order MMRs with the inner planet, like 5:1, 11:2, 6:1, 13:2, 7:1. The distribution of the best solutions, within the formal $1\\sigma$ confidence interval of the best fit, reveals a clear correlation between initial values of $a_{\\idm{c}}$ and $e_{\\idm{c}}$---the larger semi-major axis of the outer planet, the smaller its initial  eccentricity. According to the empirically observed law, the increase of the observational window results in the increase of the determined orbital periods. This would mean that in the current best fit $a_{\\idm{c}}\\simeq 4.79$~AU  is still smaller than the real value. It is difficult to predict  the real $a_{\\idm{c}}$ but very likely it is in the range $[5,6]$~AU.  During the final preparation of this paper, we have learned about a discovery of a Neptunian-mass body in the \\pstar system, close to the parent star \\citep{Santos2004}. The semi-major axis of this small planet is about $0.09$~AU.  Obviously, it is interesting to verify whether the presence of this body can affect the best-fit solutions and the dynamical behavior of the whole planetary system. The RV measurements published by the Anglo-Australian Planet Search team have the mean errors comparable to the RV semi-amplitude variations imposed by the new planet (4~m/s). For the moment, the newest observations, as reported  in the discovery paper, span a very short period of time (at present, about 80~d). Unfortunately, these measurements are not available.  They would be very helpful in improving our fits because they cover the end-part of the longer period and should constrain the parameters of the outer companion. We tried to refine our fits using the data from a scanned figure published by the Swiss team on their WWW~page\\footnote{\\url{http://obswww.unige.ch/~udry/planet/hd160691.html}}. The full data set, produced by a combination of  the Anglo-Australian measurements and these synthetic RV points, has been firstly modeled using the $N$-body signal. We used velocity offsets different for the two observatories. Still, we encountered similar problems as in the case of the Anglo-Australian data itself. The best $N$-body fits for the 3-planet  system converge to $a_{\\idm{c}}\\simeq 3.8$~AU and $e_{\\idm{c}}> 0.6$. According to the stability analysis, such fits locate the Jovian planets in a strongly chaotic zone. For the 3-planet best-fit solutions, we have $\\Chi \\simeq 1.56$ and the rms of $\\simeq 3.6$~m/s. Assuming that the innermost planet does not  disturb the motion of the Jovian companions too much and applying the GAMP for  these two planets only (otherwise, the CPU time requirement is very significant),  we derived solutions qualitatively similar to the GM1 and GM2 fits (although  the rms of these solutions exceeds $4.7$~m/s). They tend to have a large $e_{\\idm{c}} \\simeq 0.5$. Simultaneously, through the requirement of stability, the GAMP places the best fits about $a_{\\idm{c}} \\simeq 5.5~AU$. Thus, in spite of a better time coverage, the curious inconsistency between the pure  $N$-body and the GAMP-derived solutions is still evident. These results can  only be preliminary. A reliable modeling of the RV data is a difficult task  without having access to the actual measurements.  Although the elements of the innermost planet seem to be determined, due to the barely determined parameters of the outer planet, we think that  the dynamical studies concerning the whole system do not make at present a great sense.  As we have shown, the two Jovian planets can be coupled by a strong, low-order MMRs or can be well separated and relatively far from MMRs.  These situations create completely different dynamical environments for the whole planetary system.  In terms of the overall dynamical behavior, the \\pstara system appears to be similar to the HD~12661 and to the HD~169830. The Jovian planets in these systems are well separated. The variation of eccentricities spans a similar, relatively extended range from 0 to 0.5-0.6. Their semi-major axes ratio corresponds to a proximity of low-order MMRs, like 6:1 or 11:2 (for the HD~12661) or 7:1, 8:1 (similar to the HD~169830). Whether a real link exists between these systems is not clear at the moment. We still have to wait for new observations which eventually make it possible to derive significantly accurate determination of their parameters.  Using the GAMP algorithm we performed a very preliminary  analysis of the RV data of the HD~154857 system. Similarity to the \\pstar, these data reveal the presence of one Jovian planet  and a linear trend, indicating the presence of the second, more distant body. Our $N$-body fits  make it possible to approximate the elements of this planet: its mass is about 0.5~M$_{\\idm{J}}$, semi-major axis about $4$~AU and large eccentricity $\\simeq 0.78$. Of course, these elements are barely constrained. Nevertheless, the parameters of the inner planet are quite well fixed and we can already compute the stability map of this system.  It is shown in Fig.~\\ref{fig:fig13}. The relevant zone is filled up by strong MMRs, similarly to the \\pstar system.  By the same reasoning, as applied to the \\pstar system, the eccentricity of the outer planet should be strongly limited by the requirement of dynamical stability.  Our work advocates a dynamical analysis of the RV data. Even if the observational window  covers only partially the longest orbital period, we are able to derive a lot of characteristics of the  discussed system, avoiding falsifying the orbital elements by the purely kinematic approach. In fact,  the analysis presented in this paper shows that the application of the GAMP-like algorithm is essential for finding reliable solutions which well reproduce the RV data and simultaneously lead to orbitally stable configurations of the planetary companions."
    },
    "0410/astro-ph0410125_arXiv.txt": {
        "abstract": "In the neutron-rich internal shocks model for Gamma-ray Burts (GRBs), the Lorentz factors (LFs) of ions shells are variable, so are the LFs of accompanying neutron shells. For slow neutron shells with a typical LF $\\sim {\\rm tens}$, the typical $\\beta-$decay radius reads $R_{\\rm \\beta,s}\\sim {\\rm several}\\times10^{14}{\\rm cm}$, which is much larger than the typical internal shocks radius $\\sim 10^{13}{\\rm cm}$, so their impact on the internal shocks may be unimportant. However, as GRBs last long enough ($T_{90}>20(1+z){\\rm s}$), one earlier but slower ejected neutron shell will be swept successively by later ejected ion shells in the range  $\\sim 10^{13}{\\rm cm}-10^{15}{\\rm cm}$, where slow neutrons have decayed significantly. We show in this work that ion shells interacting with the $\\beta-$decay products of slow neutron shells can power a ultraviolet (UV) flash bright to 12th magnitude during the prompt $\\gamma-$ray emission phase or slightly delayed, which can be detected by the upcoming Satellite SWIFT in the near future. ",
        "introduction": "As realized by many authors, the fireball of Gamma-ray bursts (GRBs) may contain significant neutron component in essentially all progenitor scenarios (e.g., Derishev, Kocharovsky $\\&$ Kocharovsky 1999a, b; Beloborodov 2003a). The dynamics as well as the observable signatures (the neutrino emission from the proton-neutron decoupling and the changes in GRB lightcurves due to neutrons preceding the ion shell or catching up it) of a relativistic neutron-rich fireball have been investigated by Derishev, Kocharovsky $\\&$ Kocharovsky (1999a, b) firstly. Since then, many authors have had their attention on the neutron-fed GRBs. If the neutron abundance is comparable to that of proton, the inelastic collision between differentially streaming protons and neutrons in the fireball will provide us observable 5-10 GeV neutrinos (Bachall \\& M\\'{e}sz\\'{a}ros 2000, hereafter BM00; M\\'{e}sz\\'{a}ros \\& Rees 2000). The baryon-loading problem in GRB models can be ameliorated if significant fraction of baryons confined in the fireball are converted to neutrons (Fuller, Pruet $\\&$ Abazajian 2000). Further investigations on the observation implications have been presented in Pruet $\\&$ Dalal (2002, hereafter PD02) and Beloborodov (2003b, hereafter B03). In PD02, the short GRBs are assumed to be powered by external shocks and the neutron component lags behind the ion one. As the ion ejecta has been decelerated significantly by the external medium, the decayed neutron ejecta catches up the decelerated ion ejecta and powers keV emission. In B03, the neutron-ejecta keeps ahead the ion-ejecta. The $\\beta-$decay ($n\\rightarrow p+e^{-}+\\nu_{\\rm e}$) products share their momentum with the external medium immediately. As a result, the external medium have been accelerated to a ultra-relativistic velocity. The interaction between the ion ejecta and the accelerated medium is much different from the usual case (i.e., a fireball ejecta interacting with the static medium). So the presence of neutron ejecta qualitatively changes the early dynamical evolution of GRBs' remnant.  In PD02 and B03, only the ``fast'' neutron component has been taken into account. In fact, the long, complex GRBs are more likely to be powered by the interaction of shells with variant LFs, i.e., the internal shocks model (Paczy\\'{n}ski \\& Xu 1994; Rees \\& M\\'{e}sz\\'{a}ros 1994). The best fit to the multi-wavelength afterglows implies that a significant fraction of the initial kinetic energy has been converted into internal energy (Panaitescu $\\&$ Kumar 2002), which requires that the difference in velocities between the shells is significant (i.e., the corresponding LFs satisfy $\\eta_{\\rm f}\\gg \\eta_{\\rm s}$) and their masses are comparable ($m_{\\rm f}\\approx m_{\\rm s}$) (Piran 1999, hereafter P99) [Throughout this work, the subscript ${\\rm f,~s}$ represent the fast/slow shells respectively; ${\\rm n,~p}$ represent the neutron/proton component respectively]. Typically, $\\eta_{\\rm s}$ is in order of tens, and $\\eta_{\\rm f}$ is in order of hundreds. Thus, for the slow shell, the n, p component coast with $\\Gamma_{\\rm n,s}=\\Gamma_{\\rm p,s}\\simeq\\eta_{\\rm s}$ since $\\eta_{\\rm s}<\\eta_\\pi\\simeq 3.9\\times 10^2L_{52} ^{1/4}r_{0,7}^{-1/4}[(1+\\xi)/2]^{-1/4}$ (BM00), where $L$ being the total luminosity of the ejecta, $r_0$ being the radius of the central engine, $\\xi$ being the ratio of the neutrons to the protons contained in the shells. In this Letter, we adopt the convention $Q_{\\rm x}=Q/10^{\\rm x}$ for expressing the physical parameters, using cgs units. For the fast shell, the n, p components move with different LFs: $\\Gamma_{\\rm n,f}<\\Gamma_{\\rm p,f}$ since generally $\\eta_{\\rm f}>\\eta_{\\rm \\pi}$ (BM00). In the internal shocks phase, the fast ion shell catches the slower but earlier one at a radius $R_{\\rm int}\\sim 10^{13}{\\rm cm}$ and forms a new one moving with a LF $\\Gamma_{\\rm m}\\sim {\\rm several~hundred}$ (hereafter the formed ``new'' shell is named as the ``i-shell''). At $R_{\\rm int}$, the $\\beta-$decay of the neutron component is far from important unless the typical variability timescale is longer than 0.1 s (See Rossi, Beloborodov \\& Rees 2004 for more information).  As the duration of GRB is long enough, the much later ejected i-shell catches up the earlier slow neutron shell (hereafter, the ``n-shell'') at $R_{\\rm cat}\\approx 2\\Gamma_{\\rm n,s}^2c \\delta T/(1+z)$, where $\\delta T$ being the ejection time-lag of the earlier slow n-shell and the later i-shell; On the other hand, the $\\beta-$decay radius of slow neutrons reads $R_{\\rm \\beta,s}\\approx 2.6\\times 10^{13}\\Gamma_{\\rm n,s}~{\\rm cm}$, at which slow neutrons have decayed significantly. As long as $R_{\\rm cat}$ is comparable with $R_{\\rm \\beta,s}$, i.e., $\\delta T\\geq 14(1+z){\\Gamma^{-1}_{\\rm n,s,1.5}}$, the decayed products of earlier but slower ejected neutron  shells will be swept orderly by the later ejected ion shells in a range $\\simeq R_{\\rm int}-R_{\\rm \\beta,s}$. The possible emission signature powered by that interaction is of our interest in this Letter. Otherwise, for GRBs much shorter than $14(1+z){\\Gamma^{-1}_{\\rm n,s,1.5}}{\\rm s}$, the last i-shell crosses the first slow n-shell at a radius $\\ll R_{\\rm \\beta,s}$, where the $\\beta-$decay is unimportant and there is no observable signature. ",
        "conclusions": "In the standard baryonic internal shocks model for Gamma-ray bursts (Paczy\\'{n}ski \\& Xu 1994; Rees \\& M\\'{e}sz\\'{a}ros 1994), the prompt $\\gamma-$ray emission are powered by the interaction between ion shells with variable LFs. If these shells are neutron rich, the neutron component coupled with the ''fast ions'' is accelerated to a larger bulk LF ($\\sim {\\rm several~hundred}$), i.e., the fast neutron component, which is the focus of many publications (e.g., PD02, B03). However, the neutron component coupled with the ``slow ions'' can only be accelerated to a moderate bulk LF $\\sim {\\rm tens}$, for which the typical $\\beta-$decay radius $\\sim 10^{15}{\\rm cm}$ is significantly smaller than that of the fast neutrons. As long as the duration of GRB is long enough, the earlier n-shells will be swept by the much later i-shells successively till the radius $\\simeq R_{\\rm \\beta,s}$, where the slow neutrons have decayed significantly. Consequently the interaction may power observable emission.  In this Letter, with an extremely simplified model to describe the later but faster i-shells interacting with the much early but slow, decaying n-shells, we have shown that there comes a $\\sim$12th magnitude UV flash at the later time of long GRBs. That emission is bright enough to be detected by the upcoming Satellite SWIFT in the near future. The emission predicted here is independent on the poorly known medium distribution around the progenitor of GRBs, as long as it is not very dense; it is also independent of some other physical processes, such as the $e^{\\pm}$ pair emission (Fan \\& Wei 2004), the possible reverse shock emission (As usual, if we assume the number density of interstellar medium (ISM) is about $1~{\\rm cm}^{-3}$, $R_{\\rm \\beta,s}$ is much smaller than the deceleration radius $\\sim 10^{17}{\\rm cm}$, at which about half of the kinetic energy of the ejecta has been converted into the thermal energy of the shocked ISM, for a neutron-free ejecta) and so on.   The resulting net light-curve in this work increases rapidly firstly. Then there comes a flat lasting $\\sim T_{90}$, the duration of corresponding GRB. Finally, the light-curve drops sharply, where the observed emission are mainly contributed by the \"Equal-arriving surface\" effect (e.g., Kumar \\& Panaitescu 2000; Fan, Wei \\& Wang 2004), if other emission such as the reverse shock emission has not been taken into account. It should be noted that the model taken here is a great simplification of the real situation, in which both the ion/neutron shells are moving with variable LFs, it is natural to expect that the corresponding emission is variable with time, too.  However, to simulate such complicate processes is far beyond the scope of this Letter. Nonetheless, the actual flux may not be far dimmer than our prediction since averagely speaking, the LFs of the slow n-shells and the i-shells are significantly different, but their masses are comparable. So the interaction of them can power a bright UV flash in the range $\\sim 10^{13}-10^{15}{\\rm cm}$."
    },
    "0410/astro-ph0410313_arXiv.txt": {
        "abstract": "Pulsation modes have recently been observed in a handful of white dwarf (WD) primaries of cataclysmic variables, allowing an interesting new probe into the structure of accreting WD's. We briefly discuss the seismology of these objects, how stellar properties may be inferred from the observed mode frequencies, and mode driving mechanisms. For one pulsator, GW Lib, we have shown that a WD mass $M=1.05\\ M_\\odot$ and accreted envelope mass $M_{\\rm env}=0.4\\times 10^{-4}\\ M_\\odot$ give the best match to the observed pulsation periods. A first exploration of mode driving favors $T_{\\rm eff} = 14 000 K$ and a massive WD, but more work is necessary. ",
        "introduction": "Pulsation modes have been observed in isolated white dwarfs for several decades. Since the observed mode periods are sensitive to the WD structure, an industry has emerged to estimate parameters such as WD mass, light element envelope mass, and rotation rate from observed the mode periods (for a review, see e.g. Bradley 1998). GW Lib is the first discovered accreting, pulsating white dwarf \\citep{VanZetal00}. It is the shortest orbital period cataclysmic variable (CV) known at $P_{\\rm orb} = 77$ min \\citep{Thoretal02}, and is observed to have dwarf nova outbursts. In the last few years, a handful of additional accreting pulsators have been found in dwarf nova systems (see table \\ref{tab:periods}) and roughly a dozen more are expected, raising the possibility of using seismology to determine parameters for a number of accreting WD in CV's.  Mass determinations for WD in CV's are especially interesting due to possible implications for progenitor systems of Type Ia supernovae.  Thought to be accreting WDs near the Chandrasekhar mass \\citep{HillNiem00}, these progenitors are closely related to the CV population, but the precise nature of this relationship is unknown.  Important clues in this mystery lie in how the masses of CV primary WDs change over the accreting lifetime of the binary, but progress is hampered by the difficulty of measuring CV primary masses \\citep{Patt98}.  This mass evolution, as well as measuring the accreted mass directly with seismology, is also important for determining how much, if any, of the original WD material has been ejected into the ISM in classical nova outbursts, contributing to the ISM metallicity \\citep{Gehretal98}.  \\begin{table} \\label{tab:periods} \\caption{Known Accreting Pulsating White Dwarfs} \\smallskip \\begin{center} {\\small \\begin{tabular}{ccccc} \\tableline \\noalign{\\smallskip} System & Mode Periods & References & \\\\ & [sec] & \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} GW Lib & $236$, $377$, $646$ & Van Zyl et.al. (2000,2004) \\\\ SDSS 1610 & 345, 607 & Warner and Woudt (2004)\\\\ SDSS 0131 & 330, 600 & '' ''\\\\ SDSS 2205 & 330, 600  & '' ''\\\\ SDSS 1556 & 834, 1137, 1558 & Warner and Woudt poster, this meeting \\\\ SDSS 1556 & 834, 1137, 1558 & '' '' \\\\ HS 2331 & 60, 300 & Araujo-Betancour et.al (2004) \\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{center} \\end{table}  In this proceedings, we first review how seismology can be used to infer stellar parameters. Next we discuss possible mechanisms for how modes might be driven in accreting WD. We outline the differences between seismology for the accretors and isolated WD. ",
        "conclusions": ""
    },
    "0410/astro-ph0410639_arXiv.txt": {
        "abstract": "\\baselineskip=11.7pt We present a comprehensive model of the global properties of Alfv\\'{e}n waves in the solar atmosphere and the fast solar wind. Linear non-WKB wave transport equations are solved from the photosphere to a distance past the orbit of the Earth, and for wave periods ranging from 3 seconds to 3 days. We derive a radially varying power spectrum of kinetic and magnetic energy fluctuations for waves propagating in both directions along a superradially expanding magnetic flux tube. This work differs from previous models in three major ways. (1) In the chromosphere and low corona, the successive merging of flux tubes on granular and supergranular scales is described using a two-dimensional magnetostatic model of a network element. Below a critical flux-tube merging height the waves are modeled as thin-tube kink modes, and we assume that all of the kink-mode wave energy is transformed into volume-filling Alfv\\'{e}n waves above the merging height. (2) The frequency power spectrum of horizontal motions is specified only at the photosphere, based on prior analyses of G-band bright point kinematics. Everywhere else in the model the amplitudes of outward and inward propagating waves are computed with no free parameters. We find that the wave amplitudes in the corona agree well with off-limb nonthermal line-width constraints. (3) Nonlinear turbulent damping is applied to the results of the linear model using a phenomenological energy loss term. A single choice for the normalization of the turbulent outer-scale length produces both the right amount of damping at large distances (to agree with in situ measurements) and the right amount of heating in the extended corona (to agree with empirically constrained solar wind acceleration models). In the corona, the modeled heating rate differs by more than an order of magnitude from a rate based on isotropic Kolmogorov turbulence. ",
        "introduction": "Magnetic fields are known to play a significant role in determining the equilibrium state of the plasma in the solar atmosphere and solar wind (e.g., Parker 1975, 1991; Narain and Ulmschneider, 1990, 1996; Priest 1999). Much of the magnetic flux in the ``quiet'' photosphere seems to be concentrated into small (100--200 km) intergranular flux tubes. The physical processes that heat the chromosphere and corona have not yet been identified definitively, but there is little doubt that magnetohydrodynamic (MHD) effects are prevalent. Even many proposed nonmagnetic mechanisms depend on the underlying properties of the magnetically structured atmosphere. The outflowing solar wind is fed by open magnetic flux tubes, and many MHD processes have been proposed to deposit heat and momentum at locations ranging from the extended corona to interplanetary space.  The continually evolving convection below the photosphere gives rise to a wide spectrum of MHD fluctuations in the magnetic atmosphere and wind. The propagation of waves through the solar atmosphere has been studied for more than a half century (Alfv\\'{e}n 1947; Schwarzschild 1948; Biermann 1948). Although compressible (e.g., acoustic and magnetoacoustic) MHD waves are likely to be dynamically and energetically important in some regions of the atmosphere, it is the mainly incompressible Alfv\\'{e}n mode that has been believed for many years to be dominant in regions that are {\\em open to the heliosphere} (e.g., Osterbrock 1961; Kuperus, Ionson, \\& Spicer 1981). Indeed, the MHD fluctuations measured by spacecraft in the solar wind have a strongly Alfv\\'{e}nic character (Belcher \\& Davis 1971; Hollweg 1975; Tu \\& Marsch 1995; Goldstein, Roberts, \\& Matthaeus 1995).  Even though much has been learned about the generation, propagation, reflection, and damping of Alfv\\'{e}n waves in the solar atmosphere, most earlier studies have focused on only a finite range of heights and treated the interactions with regions above or below as boundary conditions. This necessarily involved the approximation that the relevant phenomena are mainly local, i.e., that they do not depend on the conditions very far away from the region being modeled. There are circumstances, however, for which this approximation breaks down. For example, the properties of reflecting Alfv\\'{e}n waves with long periods (i.e., of order 1 day) in the solar wind depend formally on the conditions ``at infinity,'' since they behave asymptotically as standing waves (e.g., Heinemann \\& Olbert 1980).  In this paper we construct a comprehensive model of the radially evolving properties of Alfv\\'{e}nic fluctuations in a representative open magnetic region of the solar atmosphere and fast solar wind. The model takes account of nonlocal interactions by tracing the waves from their origin as transverse flux-tube oscillations in the photosphere all the way out to the interplanetary medium (truncated for convenience at 4 AU). This is done with the smallest possible number of free parameters. There are two overall aims of this paper: \\begin{enumerate} \\item We wish to understand better the global ``energy budget'' of Alfv\\'{e}n waves, including relative amplitudes of inward and outward propagating waves, along the open flux tubes that feed the solar wind. \\item In order to determine how MHD turbulence contributes to the heating of the extended solar corona, we need to pin down precisely how Alfv\\'{e}n waves provide the natural preconditions for a turbulent cascade. \\end{enumerate} The second aim above was motivated by a recent study of the small-scale dissipation of MHD turbulence in the extended corona (Cranmer \\& van Ballegooijen 2003). This kinetic dissipation depends strongly on how the turbulence is excited at its largest scales, and in this paper we attempt to put the ``outer scale'' wave dynamics on firmer footing so that the resulting ``inner scale'' can be better understood. The work described by this paper builds on prior studies by Hollweg  (1973, 1978a, 1981, 1990), Heinemann \\& Olbert (1980), Spruit (1981, 1984), An et al.\\  (1990), Barkhudarov (1991), Velli (1993), Lou (1993, 1994), Lou \\& Rosner (1994), MacGregor \\& Charbonneau (1994), Kudoh \\& Shibata (1999), Matthaeus et al.\\  (1999), Hasan et al.\\  (2003), and many others.  One unique aspect of this paper is that the photospheric spectrum of transverse fluctuations---which constrains the Alfv\\'{e}n wave amplitudes at all larger radii---is specified directly from detailed observations of magnetic bright point (MBP) motions and is not (as in many other models) given as an arbitrary boundary amplitude. However, a complete physical description of the fluctuations in the photosphere (e.g., the relative inward/outward amplitudes and the kinetic/magnetic energy partition) is obtained only after the fully nonlocal wave reflection has been computed for all radii. Another way this work differs from many previous models is that the expansion and successive merging of flux tubes on granular and supergranular scales is described using a two-dimensional model of a magnetic network element in the inhomogeneous solar atmosphere.  Despite the attempted comprehensiveness of this model, we needed to make three specific approximations in order to render the calculations tractable. These approximations are summarized here, but they are also discussed further below and justified for certain regimes of applicability. First, we ignore all effects of compressible fluctuations (e.g., acoustic waves; fast and slow magnetosonic waves) despite their importance in understanding observed intensity oscillations and chromospheric heating. This is done in order to fairly assess the relative importance of the incompressible Alfv\\'{e}n mode before resorting to more involved models. Second, the Alfv\\'{e}n wave model is mainly linear, which limits its applicability in regions where the wave amplitudes become large in comparison to background field strengths and characteristic speeds.  (Some effects of nonlinearity are examined, though, in {\\S}~6.) Third, we do not model explicitly the back-reaction of the waves on the mean properties of the solar atmosphere and solar wind. We do, however, compute quantities such as the wave pressure acceleration and turbulent heating rate for use in future models of this back-reaction.  The remainder of this paper is organized as follows. In {\\S}~2 we give an overview of the physical processes to be incorporated in the wave models together with a ``cartoon'' description of the steady-state magnetic field topology. In {\\S}~3 we describe the adopted steady-state (i.e., zero-order) plasma conditions in detail. The specification of the photospheric frequency spectrum of transverse fluctuations is given in {\\S}~4, and the wave equations to be solved are given in {\\S}~5. Solutions of these equations, including some with various prescriptions for damping, are presented in {\\S}~6. We conclude with a summary of major results ({\\S}~7) and a discussion of remaining issues ({\\S}~8). Appendix A contains supplementary equations describing analytic solutions of the kink-mode wave equations for isolated flux tubes in an isothermal atmosphere. Appendix B compares various published formalisms for the non-WKB transport equations for Alfv\\'{e}n waves in an accelerating wind. Appendix C summarizes the properties of the fully developed anisotropic MHD turbulence spectrum discussed by, e.g., Cranmer \\& van Ballegooijen (2003). ",
        "conclusions": "The goal of this work was to produce a detailed and self-consistent description of the global energy budget of transverse incompressible waves in open magnetic regions of the solar atmosphere. Here we list the unique features of the model, key results, and some of the insights gained about the overall wave-plasma system. \\begin{enumerate} \\item Measurements of G-band bright points in the photosphere were used to set the power spectrum of transverse waves as the lower boundary condition of our model. The observationally inferred power spectrum was summed from two phases of MBP motion assumed to be statistically independent: isolated random walks and occasional rapid jumps due to MBP merging, fragmenting, or magnetic reconnection. \\item The steady-state plasma density, magnetic field, and flow velocity were constrained empirically from the photosphere to a distance of 4 AU. The successive merging of flux tubes on granular and supergranular scales in the atmosphere was described using a two-dimensional magnetostatic model of the magnetic network element. \\item Non-WKB wave transport equations, incorporating linear reflection terms, were solved for a grid of wave periods between 3 and $3 \\times 10^5$ s over the range of heights given above. Below a mid-chromospheric {\\em merging height} the waves were modeled as modified kink-mode flux tube waves, and above this height they were modeled as Alfv\\'{e}n waves propagating parallel to the background magnetic field. \\item The waves are reflected strongly at the transition region, with only about 5\\% of the wave energy transmitted and 95\\% reflected. Above the transition region, most periods are reflected only weakly by the large-scale radial gradient of the Alfv\\'{e}n speed, but periods exceeding 1 day are reflected strongly in interplanetary space. The period-averaged reflection coefficient from the extended corona ($r = 2 \\, R_{\\odot}$) to 1 AU ranges between $10^{-4}$ and $10^{-3}$. \\item At the photosphere, the wave periods containing the most power are between 1 and 30 minutes, but the kink-mode waves with periods greater than about 12 minutes are evanescent below the merging height. Only the shorter periods propagate to larger distances, and above the merging height the power spectrum is dominated by wave periods between about 1 and 6 minutes. In~situ observations of power spectra that are strongest at periods of several hours may be explained by many uncorrelated flux tubes rotating past spacecraft in the solar wind on these longer timescales. \\item The period-averaged transverse velocity amplitude of the waves agrees with observed off-limb nonthermal line widths from SUMER and UVCS when the MBP jump amplitude $\\sigma_j$ is taken to be $\\sim$3 km s$^{-1}$. For all reasonable values of $\\sigma_j$ between 0 and 6 km s$^{-1}$ the modeled wave amplitudes at heights greater than 0.3 AU are significantly larger than the in~situ measurements, implying that large-scale damping is needed. \\item We investigated the potential impact of the linear viscous damping of Alfv\\'{e}n waves. If the effective viscous stress timescale undergoes a transition from its collisionally dominated form to either of the possible collisionless forms (given by the $m=1$ or 2 cases of {\\S}~6.2.1), the resulting damping lengths are always much longer than the local height from the photosphere, and viscosity thus cannot damp the waves. \\item We also considered nonlinear turbulent damping using a phenomenological model of the energy loss terms in the wave transport equations. The one free parameter is the normalization of the outer-scale correlation length $L_{\\perp}$ of the turbulence, which scales with height as $B_{0}^{-1/2}$. Interestingly, a single choice for the constant $\\Lambda = 0.35$, which specifies the value of $L_{\\perp}$ at the merging height to be about 1100 km, produces both the right amount of damping above 0.3 AU to agree with the in~situ measurements and the right amount of heating in the extended corona to agree with empirically constrained solar wind acceleration models. \\item Because of the relatively weak degree of reflection in the corona, non-WKB effects are probably not important for computing the wave pressure acceleration in regions that contribute to the overall solar wind acceleration. However, non-WKB effects produce order-of-magnitude differences (from the WKB approximation) in the wave pressure acceleration at and below the transition region. The inclusion of damping also affects the total amount of wave pressure acceleration imparted to the solar wind. \\end{enumerate}"
    },
    "0410/hep-ph0410283_arXiv.txt": {
        "abstract": "We discuss the possibilities of high precision measurement of the solar neutrino mixing angle $\\theta_\\odot \\equiv \\theta_{12}$ in solar and reactor neutrino experiments. The improvements in the determination of $\\sin^2\\theta_{12}$, which can be achieved with the expected increase of statistics and reduction of systematic errors in the currently operating solar and KamLAND experiments, are summarised. The potential of LowNu $\\nu-e$ elastic scattering experiment, designed to measure the $pp$ solar neutrino flux, for high precision determination of $\\sin^2\\theta_{12}$, is investigated in detail. The accuracy in the measurement of $\\sin^2\\theta_{12}$, which can be achieved in a reactor experiment with a baseline $L \\sim (50-70)$ km, corresponding to a Survival Probability MINimum (SPMIN), is thoroughly studied. We include the effect of the uncertainty in the value of $\\sin^2\\theta_{13}$ in the analyses. A LowNu measurement of the $pp$ neutrino flux with a 1\\% error would allow to determine $\\sin^2\\theta_{12}$ with an error of 14\\% (17\\%) at 3$\\sigma$ from a two-generation (three-generation) analysis. The same parameter $\\sin^2\\theta_{12}$ can be measured with an uncertainty of  2\\% (6\\%) at 1$\\sigma$ (3$\\sigma$) in a reactor experiment with $L \\sim60 $ km, statistics of $\\sim$60 GWkTy and systematic error of 2\\%. For the same statistics, the increase of the systematic error from 2\\% to 5\\% leads to an increase in the uncertainty in  $\\sin^2\\theta_{12}$ from 6\\% to 9\\% at 3$\\sigma$. The inclusion of the $\\sin^2\\theta_{13}$ uncertainty in the analysis changes the error on $\\sin^2\\theta_{12}$ to 3\\%  (9\\%). The effect of $\\sin^2\\theta_{13}$ uncertainty on the $\\sin^2\\theta_{12}$ measurement in both types of experiments is considerably smaller than naively expected. ",
        "introduction": "There has been a remarkable progress in the studies of neutrino oscillations in the last several years. The experiments with solar, atmospheric and reactor neutrinos \\cite{cl,ga,sksolar,Ahmad:2002jz,Ahmed:2003kj, kearns04,kl162} have provided compelling evidences for the existence of neutrino oscillations driven by non--zero neutrino masses and neutrino mixing. Evidences for oscillations of neutrinos were obtained also in the first long baseline accelerator neutrino experiment K2K \\cite{K2K}.  The recent Super-Kamiokande  data on the $L/E$-dependence of multi-GeV $\\mu$-like  atmospheric neutrino events \\cite{kearns04}, $L$ and $E$ being the distance traveled by neutrinos and the neutrino energy, and the new more precise spectrum data of KamLAND and K2K experiments \\cite{kl766,K2Knu04}, are the latest significant contributions to this progress. For the first time the data exhibit directly the effects of the oscillatory dependence on $L/E$ and $E$ of the probabilities of $\\nu$-oscillations in vacuum \\cite{BP69}. We begin to ``see'' the oscillations of neutrinos. As a result of these magnificent developments, the oscillations of solar $\\nu_e$, atmospheric $\\nu_{\\mu}$ and $\\bar{\\nu}_{\\mu}$, accelerator $\\nu_{\\mu}$ (at $L\\sim$250 km) and reactor $\\bar{\\nu}_e$ (at $L\\sim$180 km), driven by nonzero $\\nu$-masses and $\\nu$-mixing, can be considered as practically established.  The SK atmospheric neutrino and K2K data are best described in terms of dominant 2-neutrino $\\nu_{\\mu} \\rightarrow \\nu_{\\tau}$ ($\\bar{\\nu}_{\\mu} \\rightarrow \\bar{\\nu}_{\\tau}$) vacuum oscillations. The best fit values and the 99.73\\% C.L. allowed ranges of the atmospheric neutrino oscillation parameters $|\\Delta m^2_{\\rm A}| = |\\Delta m^2_{31}|$ and $\\theta_{\\rm A} \\equiv \\theta_{23}$ read \\cite{kearns04}: $|\\Delta m^2_{31}| = 2.1\\times10^{-3}~{\\rm eV^2}$,~ $\\sin^22\\theta_{23} = 1.0$, $|\\Delta m^2_{31}| = (1.3 - 4.2)\\times 10^{-3}~{\\rm eV^2}$, $\\sin^22\\theta_{23} \\geq 0.85$. The sign of $\\Delta m^2_{31}$ and of $\\cos2\\theta_{23}$, if $\\sin^22\\theta_{23} \\neq 1.0$, cannot be determined using the existing data.  The combined 2-neutrino oscillation analysis of the solar neutrino and the new \\kl 766.3 Ty spectrum data shows \\cite{kl766,kl766us,kl766others} that the solar neutrino oscillation parameters lie in the low-LMA region : $\\Delta m^2_{\\odot} \\equiv \\Delta m^2_{21} = (7.9^{+0.6}_{-0.5})\\times 10^{-5}~{\\rm eV^2}$, $\\tan^2 \\theta_\\odot \\equiv \\tan^2 \\theta_{12} = (0.40^{+0.09}_{-0.07})$ \\cite{kl766}. The high-LMA solution is excluded at more than 3$\\sigma$. The value of $\\Delta m^2_{21}$ is determined with a remarkably high precision of 12\\% at 3$\\sigma$. Maximal solar neutrino mixing is ruled out at $\\sim 6\\sigma$.  The solar and atmospheric neutrino, and \\kl and K2K neutrino oscillation data require the existence of three-neutrino mixing in the weak charged lepton current. In this case the neutrino mixing is characterised by one additional mixing angle $\\theta_{13}$ -- the only small mixing angle in the PMNS matrix. Three-neutrino oscillation analyses of the solar, atmospheric and reactor neutrino data show that $\\sch < 0.05$ \\cite{kl766us,kl766others} \\footnote {{After the new background data published by the \\kl  collaboration \\cite{kl766} the bound changes to $\\sch < 0.055$ \\cite{kl766us} whereas inclusion of the recently published 391 days SNO salt phase data \\cite{sno05} gives $\\sch < 0.044$ \\cite{progress}}}.  Understanding the origin of the patterns of solar and atmospheric neutrino mixing and of $\\Delta m^2_{21}$ and $\\Delta m^2_{31}$, suggested by the data, is one of the central problems in neutrino physics today. A pre-requisite for any progress in our understanding of neutrino mixing is the knowledge of the precise values of solar and atmospheric neutrino oscillation parameters, $\\theta_{12}$, $\\Delta m^2_{21}$, and $\\theta_{23}$, $\\Delta m^2_{31}$ and of $\\theta_{13}$. In the present article we discuss the possibilities of high precision measurement of the solar neutrino mixing angle $\\theta_{12}$ in solar and reactor neutrino experiments.  The solar neutrino mixing parameter $\\sin^2\\theta_{12}$ is  determined by the current KamLAND and solar neutrino data with a relatively large uncertainty of 24\\% at 3$\\sigma$. In the future, more precise spectrum data from the KamLAND experiment can lead to even more accurate determination of the value of $\\Delta m^2_{21}$. However, these data will not provide a considerably more precise measurement of $\\sin^2\\theta_{12}$ owing to the fact that the baseline of the \\kl experiment effectively corresponds to a $\\bar{\\nu}_e$ Survival Probability MAXimum (SPMAX) \\cite{th12,shika2}. The analysis of the global solar neutrino data taking into account a possible reduction of the errors in the data from the phase-III of the SNO experiment shows that the uncertainty in the value of $\\sss$ would still remain well above 15\\% at $3\\sigma$ \\cite{skgd}.  We begin by summarising the results on $\\sin^2\\theta_{12}$, obtained using the current global solar and reactor neutrino oscillation data (section 2). We consider the improvements in the determination of $\\sin^2\\theta_{12}$, which can be achieved with the expected increase of statistics and reduction of systematic errors in the experiments which are currently operating. In particular, the effect of \\kl data, corresponding to a statistics of 3 kTy, as well as of the data of phase-III of SNO experiment, are analysed.  We turn next to future experiments. We discuss first (in section 3) the possibility of high precision determination of $\\sin^2\\theta_{12}$ in a LowNu solar neutrino experiment, designed to measure the $\\sim$MeV and sub-MeV components of the solar neutrino flux: $pp$, $pep$, $CNO$, $^7{Be}$. It is usually suggested that the LowNu experiments can provide one of the most precise measurements of the solar neutrino mixing angle \\cite{LENS,XMASS,lownu}. Detailed analysis was carried out in \\cite{roadmap} and it was concluded that a future $pp$ experiment should have accuracy better than 3\\% in order to improve on the knowledge of $\\tan^2\\theta_{12}$.  We consider a generic $\\nu-e$ scattering experiment measuring the $pp$ neutrino flux and perform a detailed quantitative analysis of the precision with which $\\sin^2\\theta_{12}$ can be determined in such an experiment. We examine the effect of including different representative values of the $pp$ neutrino induced event rate in the $\\chi^2$ analysis of the global solar neutrino data. Three values (0.68, 0.72, 0.77) of the (normalized) event rate from the currently allowed 3$\\sigma$ range are considered. The error in the measured rate is varied from 1\\% to 5\\%. We investigate how much the accuracy on $\\sin^2\\theta_{12}$ improves by adding the $pp$ flux data in the analysis. The dependence of the sensitivity to $\\sss$ on the central value of the measured $pp$ flux as well as on the measurement errors is studied. We compare the precision in $\\sss$ expected with assumed data on $pp$ neutrinos included in the analysis, with the sensitivity that can be achieved using prospective results from phase-III of SNO experiment and 3 kTy KamLAND data, and comment on the minimum error required in the LowNu $pp$ experiments to improve the precision of $\\sss$ measurement. The impact of the uncertainty due to $\\theta_{13}$ on the allowed ranges of $\\sin^2\\theta_{12}$ is studied as well.  In the section 4 we analyse in detail the possibility of a high precision determination of $\\sin^2\\theta_{12}$ in a reactor experiment with a baseline of $L \\sim (50 - 70)$ km, corresponding to a $\\bar{\\nu}_e$ Survival Probability MINimum (SPMIN). That such an experiment can provide the highest precision in the measurement of $\\sss$ was pointed out first in \\cite{th12}. A rather detailed study of the precision in $\\sss$, which might be achieved in an SPMIN experiment with the flux of $\\bar{\\nu}_e$ from the Kashiwazaki reactor complex in Japan and $L = 54$ km, was performed recently in \\cite{minat12} \\footnote{In ref. \\cite{minat12} this experiment is called SADO.}. We consider a generic SPMIN reactor  experiment with a KamLAND-type detector. We investigate the dependence of the precision on $\\sin^2\\theta_{12}$ which can be achieved in such an experiment on the baseline, statistics and systematic errors. More specifically, the spectrum data is simulated for four different true values of $\\ms$ and for each of these values the optimal baseline at which the most precise measurement of $\\sss$ could be performed is determined. We show, in particular, that an independent determination of $\\ms$ with sufficiently high accuracy would allow $\\sss$ to be measured with the highest precision over a relatively wide range of baselines. The effect of $\\sch$ uncertainty on the $\\sss$ determination is investigated in detail.  The results of the present study are summarised in section 5. ",
        "conclusions": "We have investigated the possibilities of high precision measurement of the solar neutrino mixing angle $\\theta_{12}$ in solar and reactor neutrino experiments. As a first step, we have analyzed the improvements in the determination of $\\sin^2\\theta_{12}$, which can be achieved with the expected increase of statistics and reduction of systematic errors in the currently operating solar and KamLAND experiments. With the phase-III prospective data from SNO experiment included in the current global solar neutrino and KamLAND data, the uncertainty in the value of $\\sin^2\\theta_{12}$ is expected to diminish from 24\\% to 21\\% at 3$\\sigma$. If instead of 766.3 Ty, one uses simulated 3 kTy KamLAND data in the same analysis, the 3$\\sigma$ error in $\\sin^2\\theta_{12}$ reduces to 18\\%.  We next considered the potential of a generic LowNu $\\nu-e$ elastic scattering experiment, designed to measure the $pp$ solar neutrino flux, for high precision determination of $\\sin^2\\theta_{12}$. We examined the effect of including values of the $pp$ neutrino induced electron scattering rates in the $\\chi^2$ analysis of the global solar neutrino data. Three representative values of the rates from the currently allowed 3$\\sigma$ range were considered: 0.68, 0.72, 0.77. The error in the measured rate was varied from 1\\% to 5\\%. By adding the $pp$ flux data in the analysis, the error in $\\sin^2\\theta_{12}$ determination reduces to 14\\% (19\\%) at 3$\\sigma$ for 1\\% (3\\%) uncertainty in the measured $pp$ rate. Performing a similar three-neutrino oscillation analysis we found that, as a consequence of the uncertainty on $\\sin^2\\theta_{13}$, the error on the value of $\\sin^2\\theta_{12}$ increases correspondingly to 17\\% (21\\%).  We also studied the possibility of a high precision determination of $\\sin^2\\theta_{12}$ in a reactor experiment with a baseline corresponding to a Survival Probability MINimum (SPMIN). We showed that in a $L \\sim 60 $ km experiment with statistics of $\\sim$60 GWkTy and systematic error of 2\\%, $\\sin^2\\theta_{12}$  could be measured with an uncertainty of 2\\% (6\\%) at 1$\\sigma$ (3$\\sigma$). The inclusion of the $\\sin^2\\theta_{13}$ uncertainty in the analysis changes this error to 3\\%  (9\\%). An independent determination of $\\ms$ with sufficiently high accuracy would allow, as we have shown, $\\sss$ to be measured with the highest precision over a relatively wide range of baselines. We investigated in detail the dependence of the precision on $\\sin^2\\theta_{12}$ which can be achieved in such an experiment on the baseline, statistics and systematic error. More specifically, with the increase of the statistics from 20 GWkTy to 60 GWkTy, the error diminishes from 3\\% (10\\%) to 2\\% (6\\%) at 1$\\sigma$ (3$\\sigma$). For statistics of (60 - 70) GWkTy, the increase of the systematic error from 2\\% to 5\\% leads to an increase in the uncertainty in  $\\sin^2\\theta_{12}$ from 6\\% to 9\\% at 3$\\sigma$.  We have found that the effect of $\\sin^2\\theta_{13}$ uncertainty on the $\\sin^2\\theta_{12}$ determination in LowNu $pp$ and SPMIN reactor experiments considered is considerably smaller than naively expected.  The results of our analyses for the currently running, the proposed LowNu and future reactor experiments show that the most precise determination of $\\sin^2\\theta_{12}$ can be achieved in a dedicated reactor experiment with a baseline tuned to SPMIN associated with $\\ms \\equiv \\Delta m^2_{\\odot}$.  \\vskip 0.4cm {\\bf Acknowledgments:} We would like to thank Y. Suzuki, A. Suzuki, M. Nakahata, F. Suekane, and C. Pe\\~na-Garay for useful discussions. This work was supported by the Italian INFN under the program ``Fisica Astroparticellare'' (S.C. and S.T.P.).  \\vspace{-0.5cm}"
    },
    "0410/astro-ph0410196_arXiv.txt": {
        "abstract": "We present the enigmatic galactic unclassified B[e] star Hen 2-90. Its optical spectrum is discussed and analysed in terms of the forbidden emission lines coming from a non-spherical wind. The evolutionary phase of Hen 2-90 is discussed. We raise the question whether an interacting binary nature seems to be a possible solution to account for all observed characteristics. ",
        "introduction": "Stars with the B[e] phenomenon have besides their spectral type several characteristics in common: strong Balmer emission lines, many permitted and forbidden emission lines of low ionized metals, e.g. [Fe{\\sc ii}] and [O{\\sc i}], and a strong near- or mid-infrared excess due to hot circumstellar dust.  The group of stars showing the B[e] phenomenon is heterogeneous and has been divided by \\citet{l1998} into subgroups according to their evolutionary phase. These subgroups contain supergiants, Herbig stars, symbiotic objects and compact planetary nebulae. The biggest group, however, are the unclassified B[e] stars whose evolutionary phase is not or not unambiguously known.  Here, we present the enigmatic object Hen 2-90, a galactic unclassified B[e] star. It has first been classified as a compact planetary nebula, e.g. by \\citet{h1967}; \\citet{s1992}; \\citet{c1993}; \\citet{l1998}. Later on, a bipolar jet-like structure was resolved with HST \\citep{sn2000} with several perfectly aligned knots extending up to $\\sim 10\\arcsec$~on both sides of the star. Such a structure reminds more of a young stellar object or a symbiotic object than of a compact planetary nebula. Further observations with HST \\citep{s2002} revealed even a bipolar high-ionized wind, a low-ionized wind at intermediate latitudes as well as a high-density circumstellar disk. ",
        "conclusions": "We presented the unclassified B[e] star Hen 2-90 and discussed the possible nature of this fascinating object. It shows a jet-like structure with several perfectly aligned knots on both sides of the star that seem to be ejected regularly (every 40\\,years). This jet structure is perpendicular to the disk-like structure seen with HST. This disk seems to be neutral in hydrogen since strong [O{\\sc i}] emission has been observed that can only come from this disk. During the analysis of the forbidden emission lines it turned out that C, O {\\bf and} N need to be depleted to reproduce the observed line luminosities. This fact is not consistent with any known stellar evolution scenario of single stars. In addition, a kind of accretion mechanism is certainly necessary to create the observed knots. From our comparison with other well known systems having jets and knots we could exclude the YSO and the symbiotic nature of Hen 2-90. We cannot exclude Hen 2-90 being an X--ray binary since only observations with Chandra in the soft X--ray bands have to date been performed. But it seems to be clear that some kind of interaction has been or is still going on in this system. We want to finish our contribution by raising questions open for further discussions and investigations: \\begin{enumerate} \\item No wiggling of the knots around the jet axis -- hint for a close binary? \\item The period of knot ejection of about 40\\,years -- if it is due to binary interaction must it then be an extremely excentric system? \\item The disk like structure seen on the HST image -- an outflowing disk or the remnant of an earlier common envelope phase? \\item The variety of line profiles -- hint for a complex velocity and density structure of the circumstellar matter due to ongoing mass transfer, accretion and ejection? \\end{enumerate} Only further observations (e.g. in the hard X--ray bands) will help to understand and disentangle the real nature of the enigmatic object Hen 2-90."
    },
    "0410/astro-ph0410475_arXiv.txt": {
        "abstract": "This paper presents the results of a detailed theoretical investigation of the impact of non-LTE effects and of granulation inhomogeneities on the derived iron and oxygen abundances in the metal-poor halo subgiant HD140283. Our analysis is based on both the `classical' one-dimensional (1D) stellar atmosphere models and on the new generation of three-dimensional (3D) hydrodynamical models. The non-LTE calculations presented here have been carried out without inelastic collisions with neutral hydrogen atoms. We find that if non-LTE effects are taken into account when synthetizing the Fe\\, {\\sc i} spectrum in both type of models, then the derived iron abundance turns out to be very similar in both cases. The emergent spectral line profiles in both models are very much weaker in non-LTE than in LTE because the UV overionization mechanism produces a very strong underpopulation of the Fe\\, {\\sc i} levels, in particular in the granular regions of the 3D model. As a result, the non-LTE effects on the derived iron abundance are very important, amounting to ${\\sim}0.9$ dex and to ${\\sim}0.6$ dex in the 3D and 1D cases, respectively. On the other hand, we find that non-LTE and 3D effects have to be taken into account for a reliable determination of the iron abundance from weak Fe\\, {\\sc ii} lines, because the significant overexcitation of their upper levels in the granular regions tend to produce emission features. As a result such Fe\\, {\\sc ii} lines are weaker than in LTE and the abundance correction amounts to ${\\sim}0.4$ dex for the 3D case.  We derive also the oxygen-to-iron abundance ratio in the metal-poor star HD140283  by using the O\\,{\\sc i} triplet at 7772--5 \\AA\\     and the forbidden [O\\,{\\sc i}] line at 6300 \\AA. Our 3D results for the oxygen abundance confirm the values reported in some recent investigations. While the oxygen abundance derived from the O\\,{\\sc i} IR triplet is not very sensitive to the presence of granulation inhomogeneities, such 3D effects amount to ${\\sim}-0.2$~dex for the [O\\,{\\sc i}] line. The non-LTE abundance correction for the O\\,{\\sc i} IR triplet turns out to be $-0.2$~dex, approximately.  Interestingly, when both non-LTE and 3D effects are taken into account there still remains significant discrepancies in the iron abundances derived from Fe {\\sc i} and Fe {\\sc ii} lines, as well as in the oxygen abundances inferred from the O\\,{\\sc i} and [O\\,{\\sc i}] lines. We conclude that the discrepancies could be due to uncertainties in the stellar parameters of this metal poor star. We argue that adopting $T_{\\rm eff}{\\approx}5600$ K (instead of $T_{\\rm eff}{\\approx}5700$ K) and [Fe/H]${\\approx}-2.0$ (instead of [Fe/H]${\\approx}-2.5$) substantially reduces the discrepancies in the abundances of iron and oxygen inferred from several spectral lines. Under such circumstances we find [O/Fe]${\\approx}0.5$ at [Fe/H]=$-2$. Obviously, our tentative conclusion that the metalicity of this type of metal-poor stars is significantly larger than previously thought may have far-reaching implications in stellar astrophysics. ",
        "introduction": "Reliable determinations of the oxygen and iron abundances in the atmospheres of metal-poor stars are of paramount importance for 21st Century astrophysics. This is because the abundances of oxygen and iron in metal-poor stars represent fundamental astrophysical quantities for determining the ages of the oldest stars, for constraining the models of the chemical evolution of the Galaxy, for estimating the yields of light elements (lithium, beryllium and boron) from spallation processes in the interstellar and circumstellar gas, for testing different scenarios of nucleosynthesis in supernovae, etc. Such studies largely rely upon the oxygen-to-iron abundance ratio (${\\rm [O/Fe]}={\\rm log}{\\rm (O/Fe)}_{\\star}-{\\rm log}{\\rm (O/Fe)}_{\\odot}$) versus the stellar metallicity [Fe/H].   \\subsection{The oxygen abundance problem in metal-poor stars}  The shape and spread of the [O/Fe] vs. [Fe/H] relation is still hotly debated in the literature (e.g., Israelian et al. 2001; Asplund \\& Garc\\'\\i a P\\'erez 2001; Mel\\'{e}ndez \\& Barbuy 2002; Nissen et al. 2002; Fulbright \\& Johnson 2003; Takeda 2003, and references therein).              The controversy is linked with discordant results obtained using different indicators of the oxygen abundance, such as: i) the  oxygen  [O\\,{\\sc i}] forbidden  $\\lambda$6300 \\AA\\ line; ii) the oxygen O\\,{\\sc i} infrared triplet at 7772--5 \\AA; iii) the OH lines in the near ultraviolet; iv) the OH lines in the infrared. More precisely, the oxygen abundance derived using all these lines agree more or less with each other for moderate metal-deficient disk stars having $-1 \\le {\\rm [Fe/H]} \\le 0$. However, there are clear discrepancies among the results obtained for metal-poor halo stars with lower metallicities. Some authors find that the [O/Fe] ratio shows a plateau  between 0.4 and 0.5 dex in the metallicity range -3${<}$[Fe/H]${<}-1$ (Barbuy 1988; Kraft et al. 1992; Fulbright \\& Kraft 1999; Asplund \\& Garc\\'{\\i}a P\\'{e}rez 2001; Caretta et al. 2000; Mel\\'{e}ndez \\& Barbuy 2002), while others find a linear increase with decreasing metallicity, reaching [O/Fe]$\\simeq +1$~dex at [Fe/H]$=-3$ (Abia \\& Rebolo 1989; Tomkin et al. 1992; Cavallo et al. 1997; Israelian et al. 1998, 2001; Boesgaard et al. 1999; Mishenina et al. 2000; and references therein). Recently, Nissen et al. (2002) have analyzed [O\\,{\\sc i}] and [Fe\\,{\\sc ii}] lines within the framework of 3D atmospheric models and LTE line formation concluding that there is a quasi-linear trend for [Fe/H]$<-2$, with [O/Fe]$\\simeq +0.5$~dex at [Fe/H]$=-2.5$. Furthermore, an extensive 1D re-analysis of published equivalent widths for the {O\\,{\\sc i}} IR and the forbidden [{O\\,{\\sc i}}] lines (see Takeda 2003) shows that the discrepancy between O\\,{\\sc i} and [O\\,{\\sc i}] lines in metal-poor halo stars tends to be smaller for stars with higher effective temperatures ($T_{\\rm eff}$) and gravities (${\\rm log} g$), i.e. preferentially for dwarfs and subgiants than for giants. Fulbright \\& Johnson (2003) support this conclusion. In addition, the oxygen abundances in dwarfs derived from the {O\\,{\\sc i}} IR lines tend to have values between those obtained with the UV OH and [O\\,{\\sc i}] lines.   \\subsection{The iron abundance problem in metal-poor stars}  A key problem directly related with the [O/Fe] issue in metal-poor stars is the reliability of the iron abundance determinations. In fact, the derivation of the iron abundance from Fe\\,{\\sc i} and/or Fe\\,{\\sc ii} lines is a very complex problem that requires detailed non-LTE investigations (hereafter NLTE) like the one carried by Shchukina \\& Trujillo Bueno (2001) and in this paper. In general, the Fe\\,{\\sc i} lines suffer from UV overionization in the atmospheres of late-type stars (Athay and Lites 1972; Rutten 1988; Shchukina \\& Trujillo Bueno 2001). This NLTE effect tends to be substantially more important in metal-poor stars than in solar-like stars as a result of the lower electron density and of the weaker UV blanketing (Th\\'{e}venin \\& Idiart 1999; Shchukina and Trujillo Bueno 2001; Israelian et al. 2001, 2004; Shchukina et al. 2003). We should mention also that Gratton et al.(1999) concluded that NLTE effects are rather small in metal-poor stars. However, their calculations presumably overestimated the efficiency of inelastic collisions with hydrogen atoms which thermalize the Fe\\,{\\sc i} levels and thus almost restore LTE conditions.  It is of interest to point out that while the Fe\\,{\\sc i} lines are (in principle) expected to be sensitive to the UV overionization mechanism, the LTE approximation is thought to be suitable for weak lines of Fe\\,{\\sc ii} (see, e.g., Asplund et al 1999; Nissen et al. 2002). However, as we shall see below this does not have to be necessarily the case because many subordinate Fe\\,{\\sc ii} lines whose upper levels are of odd parity (starting at $z\\rm{{^6}D^{o}}$) might be partially filled by emission due to an optical pumping mechanism similar to that investigated by Cram et al. (1980) for the solar case.   \\subsection{The solar oxygen abundance problem}  Obviously, the solar case constitutes a fundamental reference for a reliable quantification of chemical abundances in other stars. The solar iron abundance appears to be fairly well determined now (${\\rm log} \\epsilon_{\\odot}$(Fe)=7.50$\\pm 0.10$; see Shchukina \\& Trujillo Bueno 2001) --- that is, it is similar to the meteoritic value. In fact, this is the value adopted by Nissen et al. (2002) in their study of the O/Fe ratio in metal-poor stars.  Great progress has been also made towards a reliable determination of the solar oxygen abundance. The 1D modeling approach leads to discrepancies between the solar oxygen abundance determined from the permitted O\\,{\\sc i} lines and from the [O\\,{\\sc i}] and OH lines. Abundance determinations from OH lines using the 1D model of Holweger \\& M$\\ddot{\\rm u}$ller (1974) suggested first a high oxygen abundance value (${\\rm log} \\epsilon_{\\odot}$(O)=8.93$\\pm 0.03$; see Grevesse et al. 1984; Sauval et al. 1984; Anders \\& Grevesse 1989). Several years later, Grevesse \\& Sauval (1998) obtained an oxygen abundance that was lower by 0.1 dex. This lower value was found after using improved ${\\rm log} gf$ values and a slightly modified 1D model, which helped to remove the existence of a dependency of the derived oxygen abundance with the excitation potential of the lower level of the OH transitions. The forbidden [O\\,{\\sc i}] line gives also a high oxygen abundance value when the contribution of the blend due to a Ni\\,{\\sc i} line is ignored (Lambert, 1978; Sauval et al. 1984; Holweger 2001).  However, the permitted O\\,{\\sc i} lines at $\\lambda$7772--5 \\AA\\ yield a value for the oxygen abundance that is lower by 0.2 dex (Kiselman 1993; Holweger 2001). Some authors have argued that the existing NLTE calculations for these spectroscopic indicators of the oxygen abundance might not be valid due to the neglect of inelastic collisions with hydrogen. For this reason the solar oxygen abundance derived  from the [O\\,{\\sc i}] forbidden line was considered as the most reliable one, because of the suitability of the LTE approximation (hence, uncertainties of collisional rates with hydrogen atoms do not play any role). Another argument used to support a high value for the oxygen abundance is the small scatter in the oxygen abundances obtained from different OH lines. Kiselman \\& Nordlund (1995) tried to clarify the solar oxygen abundance issue by using a 3D solar photospheric model that is similar but less sophisticated than that of Asplund et al. (2000). They concluded that a downward revision of the solar oxygen abundance is required, with respect to the value derived from [O\\,{\\sc i}] and OH lines. More recently, Allende Prieto et al. (2001) have obtained a more precise estimate by taking into account the influence of the above-mentioned blend by the Ni\\,{\\sc i} line and by using the  3D hydrodynamical model of Asplund's et al. (2000). They found ${\\rm log}{\\epsilon}_{\\odot}$(O)=8.69$\\pm 0.05$, which shows that the [O\\,{\\sc i}]-based abundance derived via the 3D approach is significantly lower than that derived from classical 1D model atmospheres.  In a recent IAU Symposium on modeling of stellar atmospheres Shchukina et al. (2003) reported that the oxygen abundance obtained via NLTE synthesis of the solar O\\,{\\sc i} IR triplet in the 3D solar photosphere model of Asplund's et al. (2000) is ${\\rm log}\\epsilon_{\\odot}$(O)${\\approx}8.7$. This result is in good agreement with both Allende Prieto's et al. (2001) conclusion and with the recent 3D analysis of Asplund et al. (2004) who found excellent agreement between the averaged oxygen abundance values derived from [O\\,{\\sc i}], O\\,{\\sc i}, OH vibration-rotation and OH pure rotation lines\\footnote{It is interesting to mention that the oxygen abundances obtained by Asplund et al. (2004) from OH rotational lines using the 3D modeling approach show a large line-to-line scatter, and with sensitivity to the line strength and to the excitation potential of the transition's lower level. However, such presumably minor discrepancies are not the topic of this paper. Here we restrict ourselves to consider the NLTE determination of the solar oxygen abundance via O\\,{\\sc i} lines, for consistency with similar computations we have carried out for the metal-poor star HD140283.}.   \\subsection{Possible reasons for the abundance discrepancies}  It is currently believed that abundance discrepancies in metal-poor halo stars and the sun are dominated by deficiencies in the spectral line modeling. Errors in the measurements of equivalent widths are of minor importance with the exception of weak [O\\,{\\sc i}] and Fe\\,{\\sc ii} lines. All abundance indicators in metal-poor  stars are to a different degree sensitive to the adopted stellar parameters ($T_{\\rm eff}$, ${\\rm log} g$ and [Fe/H]). They have to  depend also on atmospheric properties such as the temperature gradient and the ratio of the line to continuous opacity. Large uncertainties could result from neglecting atmospheric inhomogeneities and velocity fields (i.e., granulation inhomogeneities produced by the stellar surface convection process). Part of the discrepancies are probably due to NLTE effects. There is now an abundance of literature concerning the iron and oxygen line formation problems in classical plane-parallel stellar atmospheric models (Th\\'{e}venin \\& Idiart 1999; Mishenina et al. 2000; Kiselman 2001; Takeda 2003; and references therein). The recent generation of 3D hydrodynamical models of stellar atmospheres have activated the present debates on the influence of granulation inhomogeneities and of NLTE effects on the determination of chemical abundances in metal-poor stars (see Asplund et al. 1999; Asplund and Garc\\'{\\i}a P\\'{e}rez et al. 2001; Shchukina \\& Trujillo Bueno 2001; Nissen et al. 2002; Asplund et al. 2004). In this respect, it is very important to emphasize that the 3D hydrodynamical simulations of Asplund and coworkers show that the atmospheres of metal-poor halo stars may have substantially lower kinetic temperatures than the radiative equilibrium temperatures of the `classical' theoretical 1D model atmospheres of Kurucz (1993) and the MARCS collaboration (Gustafsson et al. 1975).  \\subsection{Motivation}  The previous discussion makes evident that detailed, 3D+NLTE investigations of the iron and oxygen abundances are urgently needed. However, it is crucial that this type of analyses are carried out {\\em jointly} for iron and oxygen, and using the same set of atmospheric models, continuum opacity package and radiative transfer code. The present paper presents the results of such an investigation. We study the NLTE formation problem of the Fe\\,{\\sc i}, Fe\\,{\\sc ii} and O\\,{\\sc i} lines in a 3D hydrodynamical model of the atmosphere of the metal-poor star HD140283. This classical halo subgiant is widely used for studies of the very early Galaxy and the origin of the chemical elements. In order to quantify the impact of 3D simulations on the abundance determinations in HD140283 we carry out a detailed comparison with the 1D classical modeling. To this end, we use the 3D hydrodynamical atmospheric model of Asplund's et al. (1999) for the star HD140283, as well as a suitable grid of Kurucz's 1D models. Moreover, we evaluate  the extent to which 3D+NLTE effects influence the determination of the stellar parameters of this metal-poor subgiant.  Our paper considers also the case of the Sun as a reference for stellar abundance determinations with emphasis on the solar oxygen abundance derived from the O\\,{\\sc i} IR  triplet. As in our NLTE+3D analysis of the solar iron abundance (Shchukina \\& Trujillo Bueno 2001) we investigate here the solar oxygen abundance problem by using the same 3D hydrodynamical solar granulation model. ",
        "conclusions": "\\label{conclusion}  We have derived the solar oxygen abundance from the O\\,{\\sc i} IR triplet using Asplund's et al. (2000) 3D hydrodynamical model of the solar photosphere and a NLTE modeling approach. We find ${\\rm log} \\epsilon({\\rm O})=8.70\\pm 0.06$, which is in good agreement with the determinations of Allende Prieto et al. (2001), Shchukina et al. (2003) and Asplund et al. (2004). Actually, if for the O\\,{\\sc i} IR triplet we use the same oscillator strengths adopted by Asplund et al. (2004) we then obtain a value very similar to that reported by them (i.e., ${\\rm log}\\epsilon({\\rm O})=8.65$).    In agreement with previous investigations, we confirm that LTE synthesis of Fe\\,{\\sc i} lines in the 3D hydrodynamical model of this star gives flux profiles significantly different from those obtained in a 1D model for the same stellar parameters. However, when the LTE approximation is used for the synthesis of Fe\\,{\\sc ii} lines we find negligible differences between the 1D and 3D cases.  The differences between the iron abundances obtained from Fe\\,{\\sc i} lines when assuming LTE in the 1D and 3D models of the star HD140283 turn out to depend on the lower excitation potential of the spectral line under consideration. The NLTE abundance corrections for Fe\\,{\\sc i} lines are significant, especially for the low-excitation lines in the 3D model.  We have three particularly important results concerning the determination of the iron abundance in metal poor stars like HD140283:  (1) If NLTE effects are taken into account when synthetizing the Fe\\,{\\sc i} lines in the 1D and 3D models of this star, then the derived iron abundance turns out to be practically the same in both cases.  (2) Contrary to a generally accepted belief, the Fe\\,{\\sc ii} lines turn out to be significantly affected by NLTE effects. In particular, a full NLTE modeling should be carried out for the Fe\\,{\\sc ii} lines with upper levels of odd parity starting from $z\\rm{{^6}D^{o}}$, and especially the weaker lines.  (3) The iron abundance of the star HD140283 is estimated to be $\\sim 0.5$ dex higher than previously thought.   For the O\\,{\\sc i} IR triplet the NLTE correction of the oxygen abundance in the 1D and 3D models is practically the same and less than 0.2 dex. Such a NLTE correction is negligible for the [O\\,{\\sc i}] line. We have confirmed that the mean value of the oxygen abundance derived from the O\\,{\\sc i} IR lines is insensitive to 3D effects, while such effects reduce the abundance derived from the [O\\,{\\sc i}] line by $\\sim 0.2$ dex.   With the here presented NLTE calculations for Fe and O, both the 1D and 3D models lead to inconsistencies in the iron and oxygen abundance determinations. The discrepancies in the derived iron and oxygen abundances cannot be removed by taking into account NLTE and/or 3D effects. This result has led us to investigate whether these discrepancies can be resolved by a modification of the stellar parameters of the star HD140283. We have shown that previous studies may have underestimated the metallicity of this star and overestimated its effective temperature. We find [Fe/H]$\\simeq-2.0$ (instead of [Fe/H]$\\simeq-2.5$) and $T_{\\rm eff}=5600$ (instead of $T_{\\rm eff}=5700$). With these new stellar parameters the iron and oxygen abundances in the star HD140283 would be ${\\rm log}\\epsilon({\\rm Fe})=5.54\\pm 0.14$ and ${\\rm log}\\epsilon({\\rm O})=7.20\\pm 0.06$, respectively. Taking into account  the  {\\it low} value for the solar oxygen abundance (i.e., ${\\rm log}\\epsilon_{\\odot}$({\\rm O})=8.70$\\pm 0.06$) and our previously determined value of the solar iron abundance (i.e., ${\\rm log}\\epsilon_{\\odot}$({\\rm Fe})=$7.50\\pm 0.10$; see Shchukina \\& Trujillo Bueno 2001) we find that the oxygen-to-iron abundance ratio is ${\\rm [O/Fe]}\\approx 0.5$ at  [Fe/H]=$-2$.  Besides the reported NLTE effects for the Fe {\\sc ii} lines, our main conclusion here is that the metalicity of this type of metal-poor stars might well be significantly larger than previously thought. Obviously, the analysis of a single metal-poor star is not enough to fully resolve the puzzling behavior of the O/Fe ratio in metal-poor stars. However, we think that the present investigation at least indicates the road to be taken."
    },
    "0410/astro-ph0410369_arXiv.txt": {
        "abstract": "From the available STIS data of RR Tel,  that have provided a coverage with absolutely calibrated data in a wide wavelength range,  we have obtained a new determination of its reddening (E$_{B-V}$=0.00) from the comparison of the observed HeII Paschen lines decrement relative to HeII $\\lambda$ 4686 (for 24 HeII Paschen lines down to the region of the head of the series near lambda 2060 \\AA) with the theoretical one as given in Storey and Hummer (1995) for case B, T=10,000 K and log N$_{e}$=6. This new  E$_{B-V}$=0.0 value has been confirmed from a re-analysis of the IUE low resolution data. We recall that  the so far generally adopted value in the literature has been E$_{B-V}$=0.10 as obtained by Penston et al. (1983). ",
        "introduction": "We have retrieved from the HST archive  all STIS  spectra secured on Oct. 10, 2000 by Keenan et al., in particular the  echelle   and the grating spectra that cover the spectral regions of the HeII Paschen lines. These STIS data have been recently (Dec. 2003) re-calibrated  for the effects of increasing charge transfer inefficiency and for the effects of time-dependent optical sensitivity. ",
        "conclusions": ""
    },
    "0410/astro-ph0410643_arXiv.txt": {
        "abstract": "{ Weak gravitational lensing is considered to be one of the most powerful tools to study the mass and the mass distribution of galaxy clusters. However, the mass-sheet degeneracy transformation has limited its success. We present a novel method for a cluster mass reconstruction which combines weak and strong lensing information on common scales and can, as a consequence, break the mass-sheet degeneracy.  We extend the weak lensing formalism to the inner parts of the cluster and combine it with the constraints from multiple image systems. We demonstrate the feasibility of the method with simulations, finding an excellent agreement between the input and reconstructed mass also on scales within and beyond the Einstein radius. Using a single multiple image system and photometric redshift information of the background sources used for weak and strong lensing analysis, we find that we are effectively able to break the mass-sheet degeneracy, therefore removing one of the main limitations on cluster mass estimates. We conclude that with high resolution (e.g. HST) imaging data the method can more accurately reconstruct cluster masses and their profiles than currently existing lensing techniques. ",
        "introduction": "\\label{sec:introduction}  Clusters of galaxies have long been recognised as excellent laboratories for many cosmological applications. An especially important diagnostic is their number density as a function of mass and redshift. This can only be measured if reliable mass estimates of the observed clusters can be obtained. In addition, in the framework of the $\\Lambda$CDM cosmological model, the dark matter distribution in clusters likely follows the NFW profile \\citep{navarro97}.   Weak gravitational lensing is one of the most powerful tools currently available for studying the mass distribution of clusters of galaxies. The first weak lensing detection in clusters has been made by \\citet{tyson90}.  However, it was only after the pioneering work by \\citet{kaiser93} that the field began to flourish, and since then many cluster mass reconstructions have been carried out (see e.g. \\citealt{clowe01, clowe02, gavazzi04, lombardi05}). A disagreement occurring in some cases between the cluster mass estimated from the weak/strong lensing measurements with X-rays is still not well understood, although several scenarios have been proposed to resolve this issue (see e.g.  \\citealp{allen98,ettori03}).  In the absence of the redshift information, the main limitation for a precise weak lensing mass estimate is the mass-sheet degeneracy \\citep{seitz95}. If the redshifts of background sources and/or lens are not known, the transformation of the surface mass-density $\\kappa \\to \\kappa' = \\lambda \\kappa + (1-\\lambda)$, where $\\lambda \\ne 0$ is an arbitrary constant, leaves the expectation value of measured image ellipticities unchanged. In \\citet{bradac03b} we show that this degeneracy can be lifted using information on {\\it individual} source redshifts, however only if the weak lensing reconstruction is extended to the critical parts of the cluster.  Strong lensing is affected by this transformation as well. Namely, the mass-sheet degeneracy does not change the image positions (since the source position is not an observable) and flux ratios and therefore can not be broken if a single redshift multiple-image system is used. The mass-sheet degeneracy can in principle be broken using magnification effect \\citep[see][]{broadhurst95,broadhurst04b}.  In order to make full use of this method, the unlensed source counts at a given magnitude threshold must be known accurately. Given the photometric calibration uncertainties, which at the faint magnitudes one is usually dealing with easily amount to $0.1\\mbox{ mag}$, thus an uncertainty of $\\sim 10\\%$ in the unlensed source counts is typical. As shown by \\citet{schneider00}, this level of uncertainty removes a great deal of the power of the magnification method to break the mass-sheet degeneracy.   Several attempts have been made recently to measure the cluster mass profiles with weak lensing. However, as shown in \\citet{clowe01, clowe02}, it is extremely difficult to distinguish e.g. isothermal from NFW profiles at high significance using weak lensing data alone.  The authors also conclude that these difficulties mostly arise as a consequence of the mass-sheet degeneracy transformation. Therefore additional information needs to be included, such as combining the weak lensing data with strong lensing (see e.g. \\citealp{kneib03,smith04}). Another example was given by \\citet{sand03} using combined strong lensing and stellar kinematics data of the dominating central galaxy. This approach offers valuable extra constraints, however a detailed strong-lens modelling is required \\citep{bartelmann04,dalal03b}.   In this paper we use a combined strong and weak lensing mass reconstruction to determine the mass and the mass distribution of clusters. We reconstruct the gravitational potential $\\psi$, since it locally determines both the lensing distortion (for weak lensing) as well as the deflection (for strong lensing). The method extends the idea from \\citet{bartelmann96} and \\citet{seitz98}. Its novel feature is that we directly include strong lensing information. Further, weak lensing reconstruction is extended to the critical parts of the cluster and we include individual redshift information of background sources as well as of the source(s) being multiply imaged. This allows us to break the mass-sheet degeneracy and accurately measure the cluster mass and mass distribution. The method is tested using simulations, and in \\citet{bradac04b} (hereafter \\citetalias{bradac04b}) we apply it to the cluster {\\RXJ}.  In this paper we first briefly present the basics of gravitational lensing in Sect.~\\ref{sc:gravl}. In Sect.~\\ref{sc:cm} we give an outline of the reconstruction method (detailed calculations are given in the Appendix). We test the method using N-body simulations, and we present the results in Sect.~\\ref{sc:nbody}. The conclusions and summary are the subject of Sect.~\\ref{sc:concl}. ",
        "conclusions": "\\label{sc:concl}  In this paper we develop a new method based on \\citet{bartelmann96} to perform a combined weak and strong lensing cluster mass reconstruction. The particular strength of this method is that we extend the weak lensing analysis to the critical parts of the cluster. In turn, this enables us to directly include multiple imaging information to the reconstruction. Strong and weak lensing reconstructions are performed on the same scales, in contrast to similar methods proposed in the past where weak lensing information on much larger radii than the Einstein radius $\\theta_{\\rm E}$ was combined with strong lensing information (see e.g. \\citealp{kneib03}).   We test the performance of the method on simulated data and conclude that if a quadruply imaged system combined with weak lensing data and individual photometric redshifts is used, the method can very successfully reconstruct the cluster mass distribution. With a relatively low number density of background galaxies, $15 \\mbox{ arcmin}^{-2}$, we are effectively able to reproduce the main properties of the simulated cluster. In addition, with larger number densities $\\sim 100 \\mbox{ arcmin}^{-2}$ of background sources, accessible by HST, the substructures in the cluster can be resolved and the mass determination further improved.  We determine the enclosed mass within $340 {h^{-1}}\\mbox{ kpc}$ of the simulated cluster to be $(1.0 \\pm 0.1) \\times 10^{15} M_{\\odot}$, which is very close to the input value of $M_{\\rm s}(<340 {h^{-1}}\\mbox{ kpc}) = 0.99 \\times 10^{15} M_{\\odot}$. We have shown, that with the data quality we use we are effectively able to break the mass-sheet degeneracy and therefore obtain the mass and mass-distribution estimates without prior assumptions on the lensing potential.   In addition, the reconstruction algorithm can be improved in many ways. First, we use for the multiply imaged system only the information of the image positions.  The reconstruction method can, however, be modified to include the morphological information of each extended source. Instead of using a regular grid, one would have to use adaptive grids and decrease the cell sizes around each of these images. This will be a subject of a future work. Second, the photometric redshift determination does not only give the most likely redshift given the magnitudes in different filters, but also the probability distribution for the redshift. This information can be included in the reconstruction. In addition, source galaxies without redshift information can be included and different regularisation schemes can be considered.   Finally, the slight dependence on the initial conditions is getting weaker the higher the number density of background galaxies and/or multiple image systems are. In addition, it is of advantage to have a large spread in the redshift efficiency factors $Z$ of the background galaxies. For example, deep ACS images of clusters with a usable number density of $n\\sim 120 \\mbox{ arcmin}^{-2}$, or future observations with the James Webb Space Telescope will most likely make the dependence on the initial conditions negligible.   In \\citetalias{bradac04b} we will show the application of this method on the cluster {\\RXJ} and confirm that  a combination of strong and weak lensing offers a unique tool to pin down the masses of galaxy-clusters as well as their mass distributions."
    },
    "0410/astro-ph0410019_arXiv.txt": {
        "abstract": "{ The massive cluster of galaxies Abell 2219 (z = 0.228) with two spectacular gravitational lensing arcs was observed at 14.3\\,$\\mu$m (hereafter 15\\,$\\mu$m) with the Infrared Space Observatory and results were published by Barvainis et al. (\\cite{1999AJ....118..645B}). These observations have been reanalyzed using a method specifically designed for the detection of faint sources that had been applied to other clusters. Five new sources were detected and the resulting cumulative total of ten sources all have optical counterparts. The mid-infrared sources are identified with three cluster members, three foreground galaxies, an Extremely Red Object, a star and two galaxies of unknown redshift. The spectral energy distributions (SEDs) of the galaxies are fit with models from a selection, using the program GRASIL. Best-fits are obtained, in general, with models of galaxies with ongoing star formation. Infrared luminosities and star formation rates are obtained for six sources: the cluster members and the foreground galaxies. For the three cluster members the infrared luminosities derived from the model SEDs are between $\\sim5.7\\times10^{10}$\\,$\\mathrm{L}_{\\odot}$ and $1.4\\times 10^{11}$\\,$\\mathrm{L}_\\odot$, corresponding to infrared star formation rates between 10 and 24 $\\mathrm{M}_\\odot~\\mathrm{yr}^{-1}$. The two cluster galaxies that have optical classifications are in the Butcher-Oemler region of the color-magnitude diagramme. The three foreground galaxies have infrared luminosities between $1.5\\times 10^{10}$\\,$\\mathrm{L}_\\odot$ and $9.4\\times 10^{10}$\\,$\\mathrm{L}_\\odot$ yielding infrared star formation rates between 3 and 16 $\\mathrm{M}_\\odot~\\mathrm{yr}^{-1}$. Two of the foreground galaxies are located in two foreground galaxy enhancements (Boschin et al. \\cite{2004A&A...416..839B}). Including Abell 2219, six distant clusters of galaxies have been mapped with ISOCAM and luminous infrared galaxies (LIRGs) have been found in three of them. The presence of LIRGs in Abell 2219 strengthens the association between luminous infrared galaxies in clusters and recent or ongoing cluster merger activity. ",
        "introduction": "The study of clusters of galaxies is of fundamental importance in understanding the mechanisms that drive galaxy formation and evolution and in constraining cosmological theories for the formation of large-scale structures (e.g. Bahcall \\cite{1988ARA&A..26..631B}; Gladders et al. \\cite{2002AJ....123....1G}; Schuecker et al. \\cite{2003A&A...402...53S}). Moreover clusters offer the opportunity of studying environmental effects on the evolution of galaxies and, through the phenomenon of gravitational lensing, allow the observation of background galaxies that would normally be below detection thresholds (e.g. Mellier \\cite{1999ARA&A..37..127M}; Metcalfe et al. \\cite{2001IAUS..204..217M,2003A&A...407..791M}).  The comprehension of the properties of clusters of galaxies is growing rapidly thanks to a range of exceptionally good observations taken using both ground-based (e.g. Yamada et al. \\cite{2000PASJ...52...43Y}; B{\\\" o}hringer et al. \\cite{2001A&A...369..826B}; Gladders et al. \\cite{2002AJ....123....1G}; Andreon \\& Cuillandre~\\cite{2002ApJ...569..144A}) and space-based telescopes, particularly in the infrared, optical and X-ray bands (e.g. Duc et al. \\cite{2002A&A...382...60D}; Stanford et al. \\cite{2002ApJS..142..153S}; Burke et al. \\cite{2003MNRAS.341.1093B}).  Mid-infrared deep surveys with ESA's Infrared Space Observatory (ISO, Kessler et al. \\cite{1996A&A...315L..27K}) have revealed a population of starburst field galaxies that evolve significantly with redshift and have a median redshift of 0.8 (Aussel et al. \\cite{1999A&A...342..313A}; Oliver et al. \\cite{2000MNRAS.316..749O}; Serjeant et al. \\cite{2000MNRAS.316..768S}; Lari et al. \\cite{2001MNRAS.325.1173L}; Elbaz et al. \\cite{2002A&A...384..848E}; Gruppioni et al. \\cite{2002MNRAS.335..831G}; Metcalfe et al. \\cite{2003A&A...407..791M}). The mid-infrared emission from the ISOCAM sources accounts for most of the cosmic infrared background (Altieri et al. \\cite{1999A&A...343L..65A}; Franceschini et al. \\cite{2001A&A...378....1F}; Elbaz \\& Cesarsky \\cite{2003Sci...300..270E}; Metcalfe et al \\cite{2003A&A...407..791M}; Sato et al. \\cite{2003A&A...405..833S}). Mid-infrared data have been published for local clusters (e.g. Boselli et al. \\cite{1997A&A...327..522B,1998A&A...335...53B}; Contursi et al. \\cite{2001A&A...365...11C}) and distant clusters of galaxies (e.g. Pierre et al. \\cite{1996A&A...315L.297P}; L{\\' e}monon et al. \\cite{1998A&A...334L..21L}; Soucail et al. \\cite{1999A&A...343L..70S}; Fadda et al. \\cite{2000A&A...361..827F}; Metcalfe et al. \\cite{2003A&A...407..791M}; Coia et al. \\cite{coia}). The results of the observations revealed a wide dispersion in the number of luminous infrared galaxies (LIRGs, Genzel \\& Cesarsky \\cite{2000ARA&A..38..761G}) in clusters. New observations with the Spitzer Space Telescope (Rieke et al. \\cite{2001AAS...199.4401R}; Dole et al. \\cite{2003ApJ...585..617D}; Werner et al. \\cite{spitzer}) will further contribute to the understanding of the infrared population of galaxies in clusters.  The cluster of galaxies Abell 2219 ($z = 0.228$) was observed using the 15\\,$\\mu$m (LW3 : 12 to 18 $\\mu$m, $\\lambda_\\mathrm{eff}= 14.3\\,\\mu m$) filter of the ISOCAM instrument on board ISO (Cesarsky et al. \\cite{1996A&A...315L..32C}). Results were published by Barvainis et al. (\\cite{1999AJ....118..645B}, hereafter BAH99) and consist of five sources with flux densities between 530 $\\mu$Jy and 1100 $\\mu$Jy.  One of the sources was identified with the Extremely Red Object (ERO) \\object{J164023+4644} at redshift $z\\sim 1.048$ and the remaining four sources were identified with galaxies of unknown redshift.  In this paper we present the results of our re-analysis of the ISOCAM observations of Abell 2219 using a method specifically designed for the detection of faint sources (Metcalfe et al. \\cite{2003A&A...407..791M}). Section \\ref{sec:cluster} contains a short description of the cluster. Data reduction, source extraction and photometric calibration are described in section \\ref{sec:obs}.  The results are given in section \\ref{sec:results}. Spectral energy distributions, infrared star formation rates and luminosities are in section \\ref{sec:seds}. The discussion is given in section \\ref{sec:disc} and conclusions are summarized in section \\ref{sec:concl}.  We adopt H$_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_\\lambda=0.7$ and $\\Omega_\\mathrm{m} =0.3$.  With this cosmology, the luminosity distance to the cluster is $\\mathrm{D} = 1138\\,\\mathrm{Mpc}$ and 1\\arcsec\\ on the sky corresponds to 3.7 kpc.  The age of the Universe at the cluster redshift of 0.228 is 10.7 Gyr. ",
        "conclusions": "}      %  This paper presents new results from mid-infrared observations of the massive cluster Abell 2219.  The number of mid-infrared sources is increased from 5 to 10 with respect to the earlier work of Barvainis et al (\\cite{1999AJ....118..645B}). All the sources have optical counterparts. Three sources are identified with foreground galaxies, three with cluster galaxies, one with ERO J164023+4644 at $z \\simeq 1.05$ and one with a star.  The remaining two sources have undetermined redshifts and are likely to be background sources that are gravitationally lensed by the cluster.  SEDs were obtained for the three cluster members and the three foreground galaxies. Model SEDs were used to estimate parameters such as the infrared luminosity and SFR for the observed infrared sources. Models providing the best-fits to the observed SEDs are those of actively star-forming galaxies, either massive spirals, or starburst galaxies observed $\\sim 1$ Gyr after the burst event. Two of the cluster galaxies are classified as LIRGs and have a mean SFR of $\\sim 17\\,\\mathrm{M}_\\odot\\,\\mathrm{yr}^{-1}$, while the third has infrared luminosity and SFR comparable to those star forming galaxies found in other clusters. The foreground galaxies have infrared SFRs between 3 and 16 $\\mathrm{M}_\\odot~\\mathrm{yr}^{-1}$ and two of them are located in foreground galaxy enhancements.  The mid-infrared properties of the galaxies in Abell 2219 are compared with five distant clusters that were mapped with ISOCAM.  The number of LIRGs detected in Abell 2219 agrees with the expected number when compared with Cl 0024+1654 that is rich in luminous infrared galaxies."
    },
    "0410/astro-ph0410533_arXiv.txt": {
        "abstract": "About 30.5 days of nearly uninterrupted broadband photometry of the solar-type star $\\kappa^1$~Ceti, obtained with the MOST (Microvariability \\& Oscillations of STars) satellite, shows evidence for two large starspots with different rotation periods of 8.9 and approximately 9.3 days ($\\Delta\\Omega / \\Omega \\simeq 4\\%$). Ground based measurements in 2002 and 2003 of Ca II H \\& K emission reveal variations in chromospheric activity with a period of about 9.3 days. The data were obtained during the MOST commissioning phase. When the data are combined with historical observations, they indicate that the 9.3-day spot has been stable in its period for over 30 years. The photometry, with a sampling rate of approximately once per minute, was also used to search for acoustic (p-mode) oscillations in the star. We detect no clear evidence for p-modes in the $\\kappa^1$~Ceti photometry, with a noise level around 7 -- 9 $\\mu$mag at frequencies in the range 0.5 -- 4 mHz (3--$\\sigma$ detection limit of 21 -- 27 $\\mu$mag). There were no flares or planetary transits during the 30.5 days of MOST monitoring with light amplitudes greater than 2 mmag (durations greater than 200 minutes) and 3 mmag ($2 - 200$ min durations). While this rules out any close-in planets of $\\geq$0.5 Jupiter diameters with an orbital inclination close to $90^{\\circ}$, the scatter in differential radial velocities permit a close giant planet in a more highly inclined orbit. ",
        "introduction": "\\label{intro}  The production of solar flares, the evolution and migration of sunspots, and the very origin of the Sun's magnetic field are all believed to be associated with differential rotation beneath and at the solar surface. To test models of stellar dynamos, measurements of differential rotation in other stars are necessary for correlation with other parameters like magnetic variability and chromospheric activity.  Surface differential rotation (SDR) can be easily observed in the Sun -- with detailed observations going back to \\citet{Carrington1860} -- and subsurface rotation has been inferred from helioseismic data of the Sun's 5-minute p-mode oscillations (e.g., \\citet{GONG1996}). However, it is a challenge to observe SDR directly at the surfaces of other stars. With the notable exception of EK~Dra, where two spot modulations differing in period by 5.5\\% are simultaneously visible \\citep{MG2003}, evidence for SDR in most solar-type stars with prominent starspots is seen only by comparing data over many years. Light variations ascribed to rotational modulation sometimes change in period from epoch to epoch, and the most natural explanation is that the dominant spot systems appear at different latitudes and move with different rotation periods (e.g., \\citet{HB1990}; \\citet{HEHH1995}). Doppler Imaging of some solar-type stars from epoch to epoch has also revealed evidence for SDR (see reviews by \\citet{CollierCameron2002} and \\citet{strass2004}) and in at least one case, the rapidly rotating K dwarf AB~Doradus, SDR was observed in the Doppler maps in only a few consecutive rotations (\\citet{CCDS2002}). So far, there have not yet been measurements of solar-type acoustic oscillations in any star other than the Sun with sufficient frequency resolution to explore its interior rotation.  The MOST (Microvariability \\& Oscillations of STars) satellite \\citep{MOST} was pointed at $\\kappa^1$~Ceti as part of the commissioning phase of the mission, to obtain useful science during engineering tests and debugging of the MOST systems. Our space-based photometry has sufficient time coverage and precision to reveal SDR which is easily evident even in only three equatorial rotations of the star. The data also set the first meaningful limits on the amplitudes of p-mode oscillations in this star.  $\\kappa^1$~Ceti was chosen because of (1) its brightness and location in the sky at the time of these tests, (2) previous indications of spot migration and hyperactivity, and (3) its G5V spectral type, which made it a possible candidate for solar-type oscillations; we describe the star more fully in Section~\\ref{kappa}. The MOST observations are discussed in Section~\\ref{obs}. The MOST-orbit binned data produce a light curve which reveals the flux rotational modulation of a young active Sun-like star in unprecedented completeness and precision (Section~\\ref{lc}). We are able to relate the photometric modulations of spot visibility to periodic rotational variations of chromospheric activity seen in high-resolution spectroscopic observations of the Ca~II K emission in 2002 and 2003 (Section~\\ref{CaII}), the latter obtained just before the MOST observations; the same data provided a new accurate determination of $V \\sin i$ for the star. The full temporal resolution MOST photometry is used to set meaningful limits on the oscillation amplitudes in this star in Section~\\ref{p-mode}.  \\setcounter{footnote}{4}      % ",
        "conclusions": "\\label{concl}  We have presented data on the photometric and spectroscopic variability of the moderately young, solar-type star $\\kappa^1$~Ceti. We describe our continuous photometric observations over one month in November 2003 using the MOST satellite and our high-resolution spectroscopic CFHT observations over four observing runs in 2002 and 2003. The MOST data (averaged over the 101.4-min orbital period of the satellite) have an estimated point-to-point accuracy of $\\sim 0.2$ millimagnitude or $\\sim 200$ parts per million, and an unprecedented temporal coverage of 30.5 days with a completeness of 96\\%.  The MOST orbit-averaged photometry of $\\kappa^1$~Ceti shows two distinctly different periodicities -- a well-defined 8.9-day variation and a smaller, less regular variation whose period appears to be around 9.3 days or slightly longer.  The best interpretation of these variations is that two spots (or compact spot groups) are at different latitudes on the stellar surface, moving at angular rates which are different by $\\Delta\\Omega / \\Omega \\simeq 4$~\\%. This is one of the most direct evidences of surface differential rotation in a cool star other than the Sun and the second such after EK~Dra \\citep{MG2003}. All other measurements of SDR in other stars have been more indirect, either through season-to-season variations in spot period variability or through subtle broadening effects in spectral line profiles, so far observable only in the most rapidly rotating stars.  The details of the smaller slower moving spot on $\\kappa^1$~Ceti depend on how one describes and removes the larger spot, 8.9 day, variability and on the treatment of the slowly varying residuals. The multi-parametric modeling of the spots suggests that the smaller spots is located at a higher latitude, which would suggest a solar-like SDR pattern. However, independent spectroscopy of the star, some of which was obtained only a few variation cycles before the MOST photometry, reveals that maxima of the Ca~II K-line emission reversals are synchronized with a period of about 9.3 days, lending support to that period identification for the second spot. While showing significant evolution in size, shape and/or flux contrast even during our month of observations, the 9.3-day spot may in fact be the more stable one and perhaps visible over decades, based on earlier observations of $\\kappa^1$~Ceti in the literature. The shorter, 8.9 day, periodicity suggests that $\\kappa^1$~Ceti rotates faster than previously supposed and thus may be even younger, based on the correlation between equatorial rotation period and age in solar-type stars.  The unbinned MOST photometry, with an integration time of 40 seconds and a sampling of once per minute, sets limits on (1) the presence of p-mode oscillations, (2) the eruptions of flares over one month of monitoring of this historically active star, and (3) the occurrence of planetary transits in the system.  The data effectively rule out p-mode oscillations larger than 25 parts per million (ppm) at the 3--$\\sigma$ limit.  There were no flares with durations from 2 to 200 min whose white-light amplitudes could be larger than 2 -- 3 mmag. There can be no planets with orbital periods shorter than 30 days larger than about 0.5 Jupiter diameter in an orbit plane which would lead to transits as seen from Earth. Within the value of $\\sigma_{RV}\\sim $23 m~s$^{-1}$ we determine, there could be a close Jupiter-mass companion in an inclined orbit.  The information gleaned from this early set of MOST data, obtained before the satellite performance was optimized for normal operations on a target which was not part of the original planned scientific program, indicates the tremendous power of precise, nearly continuous photometry spanning many weeks, which is possible from a dedicated instrument in the right orbit in space.  There are many more applications like the ones described in this paper, and instruments like MOST can be an effective tool in studying the surface properties of rapidly rotating solar-type stars."
    },
    "0410/astro-ph0410705_arXiv.txt": {
        "abstract": "We present a photometric selection method for ULIRGs in the redshift range of $z=1$ -- 2. We utilize the most prominent spectral feature of ULIRGs, i.e. the silicate absorption feature at 9.7 $\\mu$m and an optimized filter system at mid-infrared wavelengths. These `{\\it silicate-break}' galaxies could be selected by using colour anomalies owing to the silicate absorption feature around 9.7(1+$z$) $\\mu$m.  Such filter systems are available on the {\\it Spitzer Space Telescope} but we suggest that the most promising selection criteria would be given with mid-infrared bands of {\\it ASTRO-F} satellite due to a more comprehensive set of filter bands than that of {\\it Spitzer}. We study the selection method of silicate-break galaxies with the SED model of both starbursts and quiescent star-forming galaxies, and then verify the method by using the observed spectra of nearby galaxies. We would expect that about 1000 candidates of silicate-break galaxies could be detected per square degree in current and future mid-infrared surveys. The silicate-break selection criteria will break the degeneracy between various galaxy evolution models for extragalactic source counts and place strong limits on the star formation activity at $z=1$ -- 2. Applying our silicate-break technique to initial Spitzer results we have tentatively identified the first candidate silicate-break galaxy at  $z=1.6$. ",
        "introduction": "\\label{sec:introduction}  Extreme examples of high-redshift counterparts of nearby Ultra-Luminous Infra-Red Galaxies (ULIRGs, ~\\cite{sand96}) have been found by observations with SCUBA (the Submillimetre Common-User Bolometer Array) at 850$\\mu$m (e.g. Smail et al.~\\shortcite{smail97}, Hughes et al.~\\shortcite{hugh98}). The star formation rates (SFRs) of submillimetre (submm) galaxies are estimated to be more than $10^3 M_\\odot$ yr$^{-1}$ (e.g. Smail et al. ~\\shortcite{smail02}, Chapman et al. ~\\shortcite{chapman03} ), i.e. significantly higher than that of nearby ULIRGs. The observed flux at submm wavelengths is insensitive to the classical flux-redshift relation at $z\\ga 1$, owing to the so-called negative K-correction. Nevertheless, the redshift distribution of submm galaxies seems to be rather narrow with a median redshift of $\\sim$2.4 (Chapman et al. ~\\shortcite{chapman03} ). A tentative detection of strong clustering within the submm population suggests that submm galaxies may be the progenitors of today's giant elliptical galaxies (Blain et al. ~\\shortcite{blain04}). A further clue to the evolutionary link between submm galaxies and ellipticals is presented by the Spectral Energy Distribution (SED) fitting analysis by Takagi et al. (2004).  Spheroids contain about 50 -- 70 \\% of the stellar mass in the local universe (e.g Schechter \\& Dressler ~\\shortcite{schecter87}, Fukugita, Hogan, \\& Peebles ~\\shortcite{fukugita98}). It is not yet clear whether the formation of such a large mass fraction can be explained by submm galaxies alone, including the recently found optically-faint radio galaxies at similar redshifts (Chapman et al. 2004), which could have the similar SFR with that of submm galaxies. The SFR density due to submm galaxies and optically-faint radio galaxies has a strong peak at $z=2$ -- 3 (Chapman et al. 2004), which corresponds to the time-scale of only $\\sim$1 Gyr.  Recent studies on near-IR selected galaxies in the HDF-N ~\\cite{dickinson03} and the HDF-S  ~\\cite{fontana03} show that a significant fraction of stars at the present epoch are formed at $z= 1$ -- 2, i.e. at later epochs than when most of submm galaxies are found. Therefore, ULIRGs at $z=1$ -- 2 could be a key galaxy population in understanding the formation process of spheroidal galaxies as well as submm galaxies at $z \\ga 2$. Furthermore it is probable that these infrared galaxies are the main contributors to the cosmic infrared background (c.f. Chary \\& Elbaz ~\\shortcite{chary01}). Such galaxies will be prime targets of new infrared satellites, such as {\\it Spitzer} ~\\cite{werner04} and {\\it ASTRO-F}~\\cite{cpp04a},~\\cite{murakami98}.  Various methods for the photometric pre-selection of interesting high-$z$ objects, such as Lyman-break galaxies ~\\cite{steidel03}, massive galaxies at high redshifts (van Dokkum et al. 2004; Cimatti et al. 2004), Lyman-$\\alpha$ emitters at $z \\ga 4$ (Cowie \\& Hu 1998; Taniguchi et al. 2003; Shimasaku et al. 2003), have revolutionized the field of the galaxy evolution at high redshift. These techniques are widely used to select objects for spectroscopy, and for statistical studies to derive luminosity functions and spatial correlation functions, which require a large sample at a similar redshift.  Until now, these photometric pre-selection techniques have been developed for galaxy populations observed at optical and near-infrared (NIR) wavelengths. Similar techniques for selecting luminous infrared galaxies at $z=1$ -- 2 are strongly demanded. Recently, the spectroscopic bump at 1.6 $\\mu$m due to the H$^-$ opacity minimum has been applied to galaxies detected at 24$\\mu$m with {\\it Spitzer} (Le Floc'h et al. 2004; Egami et al. 2004). This technique requires coordinated surveys at NIR wavelengths targeting the rest-frame 1.6$\\mu$m and MIR wavelengths to detect significant dust emission, and also unambiguous cross-identification between NIR and MIR sources, which is not easy for blended pairs of NIR sources (e.g. Egami et al. 2004). Therefore, any techniques using only the dust emission would be more efficient.   In the spectra of ULIRGs, the most prominent features are found at MIR wavelengths, i.e. silicate absorption at 9.7 $\\mu$m and the PAH features. Here, we investigate the possibility of selecting luminous infrared galaxies at $z\\ga 1$ with {\\it Spitzer} and {\\it ASTRO-F} by focusing on the MIR features. We hereafter refer to distant infrared galaxies selected by this method as {\\bf `silicate-break}' galaxies for simplicity. Note that Charmandaris et al. (2004) briefly discuss the possibility of estimating redshifts from the MIR features by using the Infrared Spectrograph (IRS, Houck et al. ~\\shortcite{houck04} ) peak-up imagers at 16 and 22 $\\mu$m onboard {\\it Spitzer}, although it is difficult to use these imagers for large area blank field surveys, because of the small filed-of-view.  Silicate-break galaxies could be selected by using the colour anomaly owing to the silicate absorption feature around a $9.7 (1+z)~ \\mu$m band. For example, a galaxy with the SED of Arp 220 at $z= 1.5$ would be too faint to be detected with {\\it Spitzer} in the MIPS 24 $\\mu$m band, while it would be detectable in the other infrared bands, as shown in Figure \\ref{example}. Thus, sources which are detected in all the {\\it Spitzer} bands but 24 $\\mu$m are candidates for heavily obscured galaxies at $z=1.5$. With the sensitivity of {\\it Spitzer} and {\\it ASTRO-F}, most of the silicate-break galaxies would be classified as ULIRGs.  The structure of this paper is as follows. In section ~\\ref{sec:templates} we describe the SED model used to derive the selection criteria for silicate-break galaxies. In section ~\\ref{sec:selection} we discuss the selection criteria for the {\\it Spitzer} and {\\it ASTRO-F} infrared satellite missions and predicted effectiveness of the filter combinations in detecting potential silicate-break candidates. The potential number of silicate-break galaxies at $z =1$ -- 2 are estimated using two infrared evolutionary models in section~\\ref{sec:numbers}. We give discussion and conclusions in section~\\ref{sec:conclusions}. Throughout this work we assume a flat cosmology of $\\Omega_m =0.3$, $\\Omega_\\Lambda =0.7$, and $H_0=75$ km sec$^{-1}$ Mpc$^{-1}$.    \\begin{figure}% \\resizebox{\\hsize}{!}{\\includegraphics{figure/silicate_break_lockmanhole.eps}} \\caption{An example of a potential silicate-break galaxy. The SED of Arp220 (Elbaz et al. 2002) at $z=1.5$ is compared with the 5 $\\sigma$ detection limits at {\\it Spitzer}/IRAC bands and MIPS 24 $\\mu$m achieved in the Lockman Hole region (Huang et al. 2004; Floc'h et al. 2004). At 70 and 160 $\\mu$m, we adopt the confusion limit based on the MIPS observations (Dole et al. 2004). With these detection limits, galaxies with the SED of Arp220 are detectable at all the {\\it Spitzer} bands, except for 24 $\\mu$m. Note that the luminosity is assumed to be twice of Arp220. } \\label{example} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  We have investigated the photometric pre-selection method of high-$z$ ULIRGs at $z= 1$ -- 2, by focusing on the silicate absorption feature at $9.7 \\mu$m in the rest-frame. This technique requires observations only at MIR wavelengths, and is therefore very efficient and self-consistent. Applying our silicate-break technique to the sample of Charmandaris et al. (2004) we have tentatively identified the first candidate silicate-break galaxy at  $z=1.6$ .  With {\\it Spitzer}, we could select ULIRGs at $z\\sim 1.5$ as silicate-break galaxies with $\\log (f_{24}/f_{70}) < -1.7$. This colour cut is applicable to the selection of optically thick ULIRGs; i.e. the selection is biased towards high optical depth, although the sample is free from the contamination from quiescent star-forming galaxies. A more complete selection of silicate-break galaxies is possible with the IRS peak-up imagers with additional photometry at 70 $\\mu$m to remove the contamination from quiescent star-forming galaxies at $z<1$.  With {\\it ASTRO-F}, a similar selection criterea of silicate-break galaxies to that of the IRS imagers is possible with L24 and L20W bands. The advantage of {\\it ASTRO-F} is the larger field-of-view (100 arcmin$^2$) for these bands, compared to that of the IRS peak-up imagers, allowing for larger samples of silicate-break galaxies.  The proposed selection criteria is useful to construct samples of heavily obscured galaxies at $z= 1$ -- 2. The majority of such galaxies are likely to be powered by starbursts rather than AGN. Therefore, such samples could be used to place strong limits on the star formation activity at $z=1$ -- 2. Note that spectroscopic redshifts might be difficult to measure for such heavily obscured galaxies, which are expected to be faint at optical wavelengths. Furthermore, the expected redshift range of silicate-break galaxies overlaps with the so-called `{\\it redshift desert}', in which strong emission lines are not accessible from the ground. Thus, the proposed silicate-break selection method could play important role to investigate the nature of infrared galaxies at $z>1$ bridging the gap between the IRAS-ISO and SCUBA populations.  In the {\\it ASTRO-F} NEP surveys, we expect to detect $\\sim400$ -- 1500 candidates of silicate-break galaxies depending on the evolutionary model. This large sample will allow us to derive important characteristics of ULIRGs at $z=1$ -- 2, such as their luminosity and spatial correlation function. Such data sets can also discriminate or tightly constrain multi-component number count models.  The available filters on {\\it Spitzer} and {\\it ASTRO-F} are found to be too broad to use individual PAH emission features for photometric pre-selections. Similarly, we find that the {\\it ISO} 15 $\\mu$m band is too broad to produce the colour anomaly for galaxies at $z\\sim 0.5$, owing to the silicate absorption. Thus, using the silicate-break method, {\\it Spitzer} and {\\it ASTRO-F} can provide the first opportunity to pre-select high-$z$ infrared galaxies with MIR bands."
    },
    "0410/hep-th0410119.txt": {
        "abstract": "Assuming that Quantum Einstein Gravity (QEG) is the correct theory of gravity on all length scales we use analytical results from nonperturbative renormalization group (RG) equations as well as experimental input in order to characterize the special RG trajectory of QEG which is realized in Nature and to determine its parameters. On this trajectory, we identify a regime of scales where gravitational physics is well described by classical General Relativity. Strong renormalization effects occur at both larger and smaller momentum scales. The latter lead to a growth of Newton's constant at large distances. We argue that this effect becomes visible at the scale of galaxies and could provide a solution to the astrophysical missing mass problem which does not require any dark matter. We show that an extremely weak power law running of Newton's constant leads to flat galaxy rotation curves similar to those observed in Nature. Furthermore, a possible resolution of the cosmological constant problem is proposed by noting that all RG trajectories admitting a long classical regime automatically give rise to a small cosmological constant. ",
        "introduction": "\\label{s1} % During the past few years, in the light of a series of investigations \\cite{mr,percadou,oliver1,frank1,oliver2,souma,percacciperini,frank2,litimgrav}, it appeared increasingly likely that Quantum Einstein Gravity (QEG), the quantum field theory of gravity whose underlying degrees of freedom are those of the spacetime metric, can be defined nonperturbatively as a fundamental, ``asymptotically safe'' \\cite{wein} theory. By definition, its bare action is given by a non--Gaussian renormalization group (RG) fixed point. In the framework of the effective average action \\cite{avact,avactrev,ym} a suitable fixed point is known to exist in the Einstein--Hilbert truncation of theory space \\cite{mr,souma,oliver1} and a higher--derivative generalization \\cite{oliver2} thereof. Detailed analyses of the reliability of this approximation \\cite{oliver1,frank1,oliver2} and a conceptually independent investigation \\cite{max} suggest that the fixed point should indeed exist in the exact theory, implying its nonperturbative renormalizability.  The general picture regarding the RG behavior of QEG as it has emerged so far points towards a certain analogy between QEG and non--Abelian Yang--Mills theories, Quantum Chromo--Dynamics (QCD) say. For example, like the Yang--Mills coupling constant, the running Newton constant $G=G(k)$ is an asymptotically free coupling, it vanishes in the ultraviolet (UV), i.\\,e.\\ when the typical momentum scale $k$ becomes large. In QCD the realm of asymptotic freedom, probed in deep inelastic scattering processes, for instance, is realized for momenta $k$ larger than the mass scale $\\Lambda_{\\text{QCD}}$ which is induced dynamically by dimensional transmutation. In QEG the analogous role is played by the Planck mass $m_{\\text{Pl}}$. It delimits the asymptotic scaling region towards the infrared (IR). For $k \\gg m_{\\text{Pl}}$ the RG flow is well described by its linearization about the non--Gaussian fixed point \\cite{frank1}.  Both in QCD and QEG simple local truncations of the running Wilsonian action (effective average action) are sufficient above $\\Lambda_{\\text{QCD}}$ and $m_{\\text{Pl}}$, respectively. However, as the scale $k$ approaches $\\Lambda_{\\text{QCD}}$ or $m_{\\text{Pl}}$ from above, many complicated, typically nonlocal terms are generated in the effective action \\cite{gluco,frank2,oliver0}. In fact, in the IR, strong renormalization effects are to be expected because gauge (diffeomorphism) invariance leads to a massless excitation, the gluon (graviton), implying potential IR divergences which the RG flow must cure in a dynamical way. Because of the enormous algebraic complexity of the corresponding flow equations it is extremely difficult to explore the RG flow of QCD or QEG in the IR, far below the UV scaling regime, by purely analytical methods. In QCD lattice techniques can be used to study the IR sector, but despite recent progress on dynamical triangulations \\cite{amb,ajl} there exists no comparable tool for gravity yet.  In QCD we have another source of information about its small momentum or large distance regime. If we take it for granted that QCD is the correct theory we can exploit the available experimental data on the strong interaction, interpret them within this theory, and thus obtain information about the quantum dynamics of QCD, in particular its nonperturbative IR sector, from the purely phenomenological side. An example to which we shall come back in a moment are the non--relativistic quark--antiquark potentials extracted from quarkonium data (and confirmed on the lattice). They suggest that nonperturbative IR effects modify the classical Coulomb term by adding a confinement potential to it which increases (linearly) with distance: \\begin{align} V (r) = - \\frac{a}{r} + \\kappa \\,r. \\label{1} \\end{align} Here $a$ and the string tension $\\kappa$ are constants \\cite{joos}.  In this paper we are going to apply a similar ``phenomenological'' strategy to gravity. Under the assumption that QEG is the correct theory of gravity on all distance scales, we try to describe and characterize the distinguished RG trajectory which is realized in Nature as completely as possible. We use both observational input and the available analytical RG studies.  We shall start from the flow equations in the Einstein--Hilbert truncation, determine which type of its RG trajectories the one realized in Nature belongs to, identify a regime on it where standard General Relativity is valid, and finally argue that this regime does not extend to arbitrarily large distances. In fact, for $k$ smaller than the momenta typical of the regime of standard gravity, the truncation predicts a strong increase of $G(k)$ with decreasing $k$, which, at a certain critical value of $k$, becomes infinite even. The diverging behavior is clearly an artifact of an insufficient truncation, but we shall see that the growth of Newton's constant with the distance can be understood on general grounds as due to a potential IR singularity. It is the main hypothesis of the present paper that a ``tamed'' form of this nonperturbative IR growth of $G(k)$ is a genuine feature of exact QEG.\\footnote{Using different methods or models, IR quantum gravity effects have also been studied in refs.\\ \\cite{tsamis,mottola,bertodm,bertoproc}.}  \\vspace*{2\\baselinestretch pt} The problem of the missing mass or ``dark matter'' is one of the most puzzling mysteries of modern astrophysics and cosmology \\cite{padman}.  It has been known for a long time that the luminous matter contained in a galaxy, for instance, does not provide enough mass to explain the gravitational pull the galaxy exerts on ``test masses'' in its vicinity. Typically their rotation curves $v(r)$, the orbital velocity as a function of the distance, are almost flat at large distances rather than fall off according to Kepler's law \\cite{combbook}.  Similar, but even stronger mass discrepancies are observed on all larger distance scales and in particular in cosmology. The recent high--redshift supernova and CMBR data show very impressively that the known forms of baryonic matter account only for a small percentage of the matter in the Universe. A possible way out is the assumption that the missing mass is due to some sort of ``dark matter'' which would manifest itself only by its gravitational effects. However, as to yet it has not been possible to convincingly identify any dark matter candidate, and so it might be worthwhile to think about alternatives.  It is a very intriguing idea that the apparent mass discrepancy is not due to an unknown form of matter we have not discovered yet but rather indicates that we are using the wrong theory of gravity, Newton's law in the non--relativistic and General Relativity in the relativistic case. In fact, Milgrom \\cite{mond} has developed a phenomenologically very successful non--relativistic theory, called MOdified Newtonian Dynamics or ``MOND'', which explains many properties of galaxies, in particular their rotation curves, in a unified way without invoking any dark matter. In its version where gravity (rather than inertia) is modified, a point mass $M$ produces the potential \\begin{align} \\phi (r) = - \\frac{\\overline{G} M}{r} + \\sqrt{a_{0} \\, \\overline{G} M \\,} \\, \\ln (r) \\label{2} \\end{align} where $\\overline{G}$ and $a_{0}$ are constants. The second term on the RHS of \\eqref{2} is responsible for the flat, non--Keplerian rotation curves at large distances. So far no wholly satisfactory relativistic extension of MOND is known.  Also the relativistic theory proposed by Mannheim \\cite{mannheim} where the Lagrangian is the square of the Weyl tensor tries to explain the rotation curves as due to a non--Newtonian force. The corresponding potential is of the form \\begin{align} \\phi (r) = - \\frac{\\overline{G} M}{r} + \\widetilde \\kappa \\, r. \\label{3} \\end{align} The resulting rotation curves do not become flat but still seem to be in accord with the observations.  For a detailed discussion of other attempts at modifying gravity at astrophysical distances and a comprehensive list of references we refer to \\cite{aguirre}. A possible relation to quintessence has been speculated about in \\cite{cwgal}.  \\vspace*{2\\baselinestretch pt} In the present paper we are going to explore the idea that the IR quantum effects of QEG, in particular the growth of $G$ at large distances, induces a modified Newtonian potential similar to \\eqref{2} or \\eqref{3}. If so, one can perhaps solve the missing mass problem in a very elegant and ``minimal'' manner by simply quantizing the fields which are known to exist anyhow, without having to introduce ``dark matter'' on an ad hoc basis.\\footnote{See refs.\\ \\cite{bertodm,bertoproc} for a related analysis within a perturbatively renormalizable higher derivative gravity.}  It is particularly intriguing that the potentials we would like to derive within QEG are strikingly similar to the nonperturbative quark--antiquark potentials generated by (quenched) QCD. In particular eqs.\\ \\eqref{1} and \\eqref{3} are mathematically identical, describing ``linear confinement'', while the MOND potential increases slightly more slowly at large distances. In view of the many similarities between QCD and QEG it is hard to believe that this should be a mere coincidence.  The purpose of the present paper is to learn as much as possible about the gravitational RG trajectory Nature has chosen and to investigate the possibility that the IR renormalization effects of QEG are ``at work'' at galactic and cosmological scales.  The remaining sections of the paper are organized as follows. In Section \\ref{s2} we discuss the Einstein--Hilbert truncation of theory space with an emphasis on the strong IR renormalization effects it gives rise to. In Sections \\ref{s3} and \\ref{s4}, we analyze the RG trajectory realized in Nature, first using mostly analytical results (Section \\ref{s3}) and then also phenomenological input (Section \\ref{s4}). In Section \\ref{s5} we employ a plausible model of the trajectory in the deep IR to demonstrate that an extremely tiny variation of Newton's constant would explain the observed flat rotation curves. The results are summarized in Section \\ref{s6}. % % % % %===================================================================== % % % % ",
        "conclusions": "\\label{s6} % In this paper we assumed that Quantum Einstein Gravity correctly describes gravity on all length scales and tried to identify the underlying RG trajectory realized in Nature. Along this trajectory, we found a regime where the running of the gravitational parameters is unmeasurably small, the domain of classical General Relativity. The renormalization effects become strong both at momentum scales $k$ much larger (UV regime) and much smaller (IR regime) than those of the GR regime. In the UV regime, for $k \\to \\infty$, the trajectory approaches a non--Gaussian fixed point. By definition, QEG is the ``asymptotically safe'' theory whose bare action equals the fixed point action, guaranteeing its nonperturbative renormalizability. We analyzed a potential IR divergence (instability) in the flow equation which occurs in the Einstein--Hilbert truncation already, and we argued that it triggers significant renormalization effects in the IR, i.\\,e.\\ at distances larger than those where classical GR was successfully tested. In particular, Newton's constant increases for decreasing $k$ so that $G$ should appear to be larger at galactic or cosmological scales than in a terrestrial or solar system ``laboratory''. Looking for possible manifestations of this effect we analyzed the possibility that the almost flat plateaus observed in the rotation curves of spiral galaxies, usually attributed to the presence of dark matter, are actually due to the RG running of Newton's constant. While the galactic regime is inaccessible with the Einstein--Hilbert truncation, a phenomenological analysis in the framework of the improved action approach revealed that a power law--type scale dependence $G (k) \\propto k^{-q}$, corresponding to an (approximately) constant anomalous dimension $\\eta_{\\text{N}}$, leads to rotation curves which are in qualitative agreement with those found in Nature. In particular they become perfectly flat at sufficiently large distances.  As a by--product we arrived at a new understanding of the cosmological constant problem: By analyzing the RG trajectories among which Nature could have chosen we saw that there exist no trajectories which, at the same time, would predict a long classical regime and a large cosmological constant. If there is a Universe inhabited by human beings who manage to verify that gravity behaves according to classical General Relativity over many different length scales, they unavoidably will observe a cosmological constant which is extremely small compared to the value of $1/G \\equiv m_{\\text{Pl}}^{2}$ they measure. Conversely, if $\\Lambda$ is not small, it is also not constant, so GR is not valid, and there is no point wondering about the value of a non--constant ``constant''.  This argument does not seem to require a detailed understanding of the IR effects since the data indicate that in our world, between the end of the GR regime and the Hubble scale, $G$ and $\\Lambda$ changed by comparatively small factors only.  As for the IR effects, it will be exciting to see whether the properties of the RG trajectory of QEG which we predicted can be confirmed by more advanced ab initio calculations. Clearly more general truncations for the effective average action, involving nonlocal invariants for instance, should be explored for this purpose. In view of recent progress \\cite{ajl} made on causal (Lorentzian) dynamical triangulations \\cite{amb} one also may hope that at some point it will be possible to make contact with numerical simulations. This could lead to an independent confirmation of the picture we have drawn in the present paper.  % % % % \\vspace*{3.3mm} \\noindent Acknowledgement: We would like to thank A.~Bonanno, O.~Lauscher, R.~Loll, F.~Saueressig, and J.~Smit for helpful discussions. % % % % %===================================================================== % % % %"
    },
    "0410/astro-ph0410274_arXiv.txt": {
        "abstract": "Direct numerical simulations of turbulent Hall dynamos are presented. The evolution of an initially weak and small scale magnetic field in a system maintained in a stationary turbulent regime by a stirring force at a macroscopic scale is studied to explore the conditions for exponential growth of the magnetic energy. Scaling of the dynamo efficiency with the Reynolds numbers is studied, and the resulting total energy spectra are found to be compatible with a Kolmogorov type law. A faster growth of large scale magnetic fields is observed at intermediate intensities of the Hall effect. ",
        "introduction": "In recent years, the relevance of two-fluid effects has been pointed out in several astrophysical \\citep{Balbus,Sano,MGM1,MGM2} as well as laboratory plasmas \\citep{Mirnov}. The standard magnetohydrodynamic (MHD) framework for the study of astrophysical plasmas  may not be adequate in the presence of strong magnetic fields and/or low ionization; the electric conductivity then is not isotropic and nonlinear effects arise in Ohm's law.  In low temperature accretion disks around young stellar objects or in dwarf nova systems in quiescence, for example, the plasma is only partially ionized with a small abundance of charged particles. As a result, two new effects appear in Ohms's law: ambipolar diffusion and Hall currents. The predominance of either of these effects is determined by the ionization fraction and the plasma density \\citep{Sano,Braginskii}. The impact of ambipolar diffusion on dynamo action has been already studied by \\citet{Zweibel} (see also \\citet{BraSub}). The Hall effect affects the dynamics of protostellar disks \\citep{Balbus}. In neutron stars magnetic fields are so strong that the Hall term can be even more important than the induction term for the magnetic field evolution \\citep{Muslimov}. The Hall effect is also known to be relevant in others astrophysical scenarios \\citep[and references therein]{Sano,MGM3}.  The Hall effect, when strong, is expected to seriously affect the MHD results on the  generation of magnetic fields in astrophysical and laboratory plasmas by inductive motions in a conducting fluid (dynamo effect). More specifically, it is expected to modify the growth and evolution of magnetic energy, since the addition  of the Hall term to the MHD equations leads to the freezing of the magnetic field to the electron flow (in the non-dissipative limit) rather than to the bulk velocity field. The first studies on the impact of Hall currents on dynamo action \\citep{Helmis1,Helmis2} were carried out using mean field theory and the first-order smoothing approximation \\citep{Krause}. Helmis obtained decreasing dynamo action as the strength of the Hall terms increased. Recently, the impact of the Hall effect has been studied using the kinematic approximation or focusing on particular geometries, both in numerical simulations \\citep{Galanti} and in experimental and theoretical studies \\citep{Heintzmann,Hantao,Rheinhardt,Mirnov}. A general closure scheme was also proposed to compute the contribution of the Hall term to the dynamo action $\\alpha $-effect \\citep{MGM1} using mean field theory and the reduced smoothing approximation \\citep{Blackman}; it was found that  the Hall effect could either suppress or enhance dynamo action. Expressions for the turbulent diffusivity using this closure were also derived for particular cases in \\citet{MGM2}.  In the present work, we report results from direct numerical simulations of the dynamo action in MHD and Hall-MHD at moderate Reynolds numbers  with strong kinetic helical forcing. The study of the scaling of the dynamo efficiency with increasing Reynolds numbers is the main aim of this paper. In a previous paper \\citep{MGM3}, we showed that three distinct dynamo regimes can be clearly identified: (1)  dynamo activity is enhanced (Hall-enhanced regime), (2) it is inhibited (Hall-suppressed regime), (3) it asymptotically approaches the MHD value (MHD regime). These regimes arise as a result of the relative ordering between the relevant lengthscales of the problem. Namely, the energy-containing scale of the flow, the Hall length, and the correlation length of the magnetic seed. Simulations in \\citet{MGM3} were performed with $64^3$ spatial grid points and fixed Reynolds numbers.  This study left unanswered the  question whether the enhancement of dynamo action by Hall effect  would increase or decrease with increasing Reynolds numbers and scale separation.  Theoretical estimates using mean field theory and a particular choice for the small scale fields suggest that the efficiency of Hall-MHD dynamos compared with the MHD counterpart  will  increase as the scale separation is increased \\citep{MGM1}. In this work we present direct simulations with higher spatial resolutions and Reynold numbers that confirm this result. Also, the magnetic, kinetic, and total energy spectra developed in Hall-MHD turbulence are calculated and studied. These spectra correspond to the dynamo regime, where no imposed currents are present, i.e, the  magnetic fields are purely self-generated. This regime is rather general and probably common to most astrophysical flows \\citep{Haugen04}.  The paper is organized as follows. In Section 2 we present the general equations describing the evolution of the fields. In Section 3, the code used to numerically integrate the Hall-MHD system is described, and all the simulations made with different resolutions and parameters are listed. Section 4 presents the results obtained in the MHD limit of our equations; the results obtained are similar to those of other authors \\citep{Meneguzzi,Brandenburg}, and provide a reference  set to compare the Hall-MHD solutions with. Section 5 is devoted to the results of Hall-MHD simulations. In Section 6 we discuss a subset of MHD and Hall-MHD simulations made to study the evolution of large scale magnetic fields. Finally, in Section 7, we give a brief summary of the current effort. ",
        "conclusions": "In this paper we have  presented the  results of direct numerical simulations of turbulent dynamo action in Hall-MHD. We find that with increasing Reynolds number and scale separation, the Hall MHD  dynamo works more efficiently when the Hall length is close but larger than the dissipation scale (Hall-enhanced regime). For larger values of $\\epsilon$ the Hall MHD dynamo is less efficient. In addition, the value of $\\epsilon$ (which measures the strength of the Hall term) at which the dynamo is most efficient decreases at higher Reynolds numbers.  An acceleration of the process responsible for the growth of a large scale magnetic field is observed at moderate values of $\\epsilon$. Although these simulations are made at Reynolds numbers which are far away from realistic values for astrophysical plasmas, the results obtained are encouraging; the dynamos tend to work better at high Reynolds numbers  By calculating the magnitude and nature of the generated magnetic field as the amplitude of the Hall term is varied, we obtain new evidence showing that the Hall dynamo can be fundamentally different from its classical MHD counterpart."
    },
    "0410/astro-ph0410568_arXiv.txt": {
        "abstract": "We present observations of the quiet solar photosphere in the \\ion{Fe}{1} lines at 6302 \\AA \\, where at least four different spatial locations exhibit upwards-directed supersonic flows. These upflows can only be detected in the circular polarization profiles as a double-peaked structure in the blue lobe of both \\ion{Fe}{1} lines. We have detected cases of either magnetic polarity in the data. The polarization signals associated with the upflows are very weak, which is probably why they had not been seen before in this type of observations. We propose that the observed flows are the signature of aborted convective collapse, similar to the case reported by \\citeN{BRRHC+01}. Our data indicates that this phenomenon occurs frequently in the quiet Sun, which means that many magnetic elements (although the fraction is still unknown) are destroyed even before they are formed completely. The spectral signatures of supersonic upflows reported here are probably present in most spectro-polarimetric observations of sufficient signal-to-noise and spatial resolution. ",
        "introduction": "\\label{sec:intro}  The term {\\it quiet Sun} is commonly used to refer to the solar atmosphere away from sunspots, faculae and other large-scale concentrations of magnetic flux. The adjective ``quiet'', however, seems now less of a synonym of ``unintersting'' than it did only a few years ago. This is due to both recent numerical simulations (e.g., \\citeNP{EC01}; \\citeNP{CSN+04}) and new spectro-polarimetric observations of high sensitivity and resolution (e.g., \\citeNP{SAL00}; \\citeNP{SNSA02}; \\citeNP{KCS+03}; \\citeNP{DCKSA03}). From the point of view of the magnetic energy balance it is worth mentioning that, according to these works, there is probably more (unsigned) magnetic flux in the quiet Sun than in all the active regions during solar maximum. The small-scale structure of the magnetic elements is also a hot topic of debate (e.g., \\citeNP{SNSA03}).  The present work adds to this more exciting view of the quiet photosphere by reporting on the detection of supersonic upflows, which seem to be a relatively frequent event. We argue that these flows are likely the last stages of the process observed by \\citeN{BRRHC+01} (hereafter BRCKR) in the infrared. According to their scenario, the upflows are the signature of an aborted convective collapse and subsequent disintegration of the magnetic structure. Our observations support this picture to some extent and remove any caveats that might have existed on their observation (e.g., structures moving in or out of the slit). ",
        "conclusions": "\\label{sec:conc}  The observations presented in this paper reveal supersonic upflows in the quiet solar photosphere and indicate that this is a rather common occurrence. The upflows are probably the signature of an exploding magnetic element after an aborted convective collapse proposed by BRCKR. It is important to point out that we do not observe the flux destruction directly. However, the assumption that we are observing the same process seems appropriate based on the following facts: \\begin{itemize} \\item The only observations reported in the literature with upflows inside photospheric magnetic elements is the aborted convective collapse of BRCKR. \\item The magnitude of the upflows is the same, very close to the sound speed of the photosphere at the heights where the lines form. This is what one would expect from a shockwave propagating inside the magnetic element. \\item If one assumes that we are observing the same process and that their collapse time scales are typical, it is possible to estimate the probability that BRCKR observed an aborted convective collapse in their time-series. We find that this probability is of the order of 1, which would be an unlikely coincidence if our assumption were incorrect. \\end{itemize}  From a 2-component inversion of the double-peaked profiles we obtain a model atmosphere that is consistent with a shockwave propagating in one of the components. The inversion even retrieves a zero field below the front shock, which would also be in agreement with the destruction of the magnetic element reported by BRCKR.  Interestingly, we only observe instances in which the two atmospheric components have similar Stokes~$V$ amplitudes. A possible explanation to this could be the following. The upflowing component cannot have a larger amplitude because, if it is indeed an exploding magnetic element, the field must be weaker in this component. If, on the other hand, it had a significantly lower amplitude then its Stokes~$V$ profile would be buried in the blue lobe of the quiet component and we would not be able to see it. Therefore, it is likely that we are only observing a small fraction of all exploding elements.  Because of the selection effect mentioned above and the fact that we are close to the sensitivity limit of the data, the frequency of occurrence reported here should be taken as a (perhaps severley underestimated) lower limit.  Finally, we find that the spatial extent of the upflows are quite puzzling, especially region~A. The upflows span here a significantly large area of more than 2 by 2 arc-seconds. Considering that the magnetic filling factor is very small in this area ($\\sim$10\\%), this is most certainly not a single magnetic structure disintegrating. However, if this area is composed of a collection of small magnetic elements (as is generally thought that network elements are), then what physical mechanism would syncrhonize the exploding elements as far as 2 arc-seconds apart?  It is important to continue investigating these upflows in more detail, given the potential implications of this work on the generation and destruction of small-scale magnetic elements in the quiet Sun. As mentioned earlier, these fields may be carrying most of the solar magnetic flux. It is fundamental to determine whether convective collapse is an efficient mechanism for concentrating field and these aborted instances constitute only a small fraction of the cases, or if on the contrary most (or perhaps even all) of the collapsing elements are ultimately doomed to disintegrate sending shockwaves into the photosphere. The amount of energy deposited in the photosphere, and also in the upper atmosphere, by the shocks might be important if a large number of these events are taking place continuously in the quiet Sun. Clearly, further study is required with both new simulations and observations to shed some light into these issues."
    },
    "0410/astro-ph0410042_arXiv.txt": {
        "abstract": "We present initial results of a project to measure mean longitudinal magnetic fields in a group of sdB/OB/O stars. The project was inspired by the discovery of three super-metal-rich sdOB stars, each having metals (e.g. Ti, V) enhanced by factors of 10$^3$ to 10$^5$. Similar behaviour is observed in chemically peculiar A stars, where strong magnetic fields are responsible for the enrichment. With this in mind, we obtained circularly polarised spectra of two of the super-metal-rich sdOBs, two \"normal\" sdBs and two sdOs using FORS1 on the ESO/VLT. By examining circular polarisation in the hydrogen Balmer lines and in helium lines, we have detected magnetic fields with strengths of 1-2 kG in most of our targets. This suggests that such fields are relatively common in hot subdwarfs. ",
        "introduction": "\\par Recently, we have found several sdBs that show a very large enrichment of iron group elements (Edelmann et al. 2001). Three out of thirteen targets showed abundances of these elements of up to 10$^5$ times the solar value. Around 23\\% of the lines in the peculiar sdBs remain unidentified. These metal-enriched sdBs are reminiscent of the chemically peculiar A stars, where strong magnetic fields play a major role in the surface abundance patterns. Although no direct evidence for magnetic fields was found, we felt a proper examination of these peculiar objects for such a field was clearly required.  \\par Strittmatter \\& Norris (1971) studied the role of magnetic fields in chemically peculiar A stars, and several of their conclusions may be relevant to sdBs. The abundance pecularities in Ap stars are, in general, confined to the stellar surface. If the magnetic field is strong enough, Strittmatter \\& Norris suggest that magnetic braking will occur, causing angular momentum loss and resulting in a slowly rotating star. The slowing of the star's rotation can lead to peculiar surface abundances, however there is no well established or detailed theory. In the case of sdB stars, the vast majority are slowly rotating, with observational limits of $v\\sin i<5$\\,km\\,s$^{-1}$. The only sdB star for which rotation has been detected is PG1605+072 (Heber et al. 1999). The metal lines of all the stars studied by Edelmann et al. (2001) are very sharp, consistent with projected rotational velocities of $v\\sin i<5$\\,km\\,s$^{-1}$. Could the slow rotation of sdB stars, coupled with a magnetic field of 1-2 kG, be the cause of the peculiar abundance patterns of some sdB and sdOB stars? To test this idea, we used the ESO/VLT equipped with FORS1 to investigate circular polarisation in hot subdwarfs with peculiar surface abundances. We have also observed apparently \"normal\" sdB and sdO stars for comparison.  \\par Elkin (1996) conducted a search for magnetic fields in sdB and sdO stars using the 6m Special Astrophysical Observatory, however, his study used several different measurement techniques and was therefore quite inhomogeneous. Our study is the first homogeneous search for magnetic fields in sdB stars. Elkin (1996) reported the detection of fields in two targets, the sdO star BD\\,$+75^\\circ 325$ and the sdOB star Feige\\,66. The other results were presented in Elkin (1998) and were largely inconclusive, with errors for some targets approaching $\\pm$1 kG. We do not believe this is sufficient for sdB stars, since, based on Elkin's detection of a variable field in the sdOB star Feige 66 with extrema of $-$1.3 kG and +1.75 kG, we expect field strengths to be no larger than $\\sim$1-2 kG. By measuring the polarisation of the Balmer lines using the method of Bagnulo et al. (2002), we expected our typical error for each target would be 200-300 G. ",
        "conclusions": ""
    },
    "0410/astro-ph0410618_arXiv.txt": {
        "abstract": "We present results of joint modeling of the probability distribution function (PDF) and the one-dimensional power spectrum (PS) of the \\op forest flux decrement.  The sensitivity of these statistical measures to the shape and amplitude  of the linear matter power spectrum is investigated using N-body simulations  of two variants of the $\\Lambda$CDM cosmology. In the first model, the linear power spectrum has a scale-invariant spectral  index $n_s=1$, whereas in the second, it has a negative running index  (RSI), $\\dd n/\\dd\\ln k<0$. We generate mock catalogs of QSO spectra, and  compare their statistical properties to those of the observations at  $z=3$. We perform a joint fit of the  power spectrum and  the PDF. A scale-invariant model with $\\sigma_8=0.9$ matches well  the data if the mean IGM temperature is $T\\lsim 1.5\\times 10^4\\kel$.  For higher temperature, it tends to overestimate the flux power spectrum  over scales $k\\lsim 0.01\\skm$. The discrepancy is less severe when the PS alone is  fitted. However, models matching the PS alone do not yield a  good fit to the PDF.  A joint analysis of the flux PS and PDF tightens the constraints on the model parameters and reduces systematic biases.  The RSI model is consistent with  the observed PDF and PS only if the  temperature is $T\\bsim 2\\times 10^4\\kel$.  The best fit models reproduce the slope and normalisation of the column  density distribution, irrespective of the shape and amplitude of the linear power spectrum. They are also consistent with  the observed  line-width distribution given the large uncertainties. Our joint analysis suggests that $\\sigma_8$ is likely to be in the range 0.7-0.9 for a temperature $1\\lsim T\\lsim 2\\times 10^4\\kel$ and a reasonable reionization history. ",
        "introduction": "\\label{data}  In this paper, we will compare the simulated statistics to the flux power spectrum  and PDF measured by McDonald \\etal (2000) (hereafter M00), which were  obtained from a sample of eight high-resolution QSO spectra.  Regarding the flux power spectrum, we consider the data points in the  range $0.005<k<0.04\\skm$. The lower limit $k=0.005\\skm$ is set by the size of our simulations (cf. Section~\\S\\ref{nbody}). Note that it is larger than the characteristic   wavenumber $k\\sim 0.003\\skm$ below which continuum fitting errors are expected to become important (e.g. Hui \\etal 2001). The upper limit  $k=0.05\\skm$ is chosen because of concern about metal contamination on smaller scales (e.g. M00; Kim \\etal 2004).  The observed flux PDF is very sensitive to continuum fitting, especially in the high transmissivity tail. We therefore exclude the data points with $F\\geq 0.8$ from the analysis to avoid dealing with continuum fitting  error. In addition, we discard also the measurement at $F=0$ since it is  strongly affected by noise (e.g. Rauch \\etal 1997).  Finally, we emphasize that we work with the power spectrum of the transmitted flux $F$, $\\dfk$, which  differs from that of the flux density contrast,  $\\delta_{\\rm F}\\equiv\\left(F/\\la F\\ra -1\\right)$, by a multiplicative  factor $\\la F\\ra^2$. $\\dfk$ has the advantage that it does not depend explicitly on the mean flux level $\\la F\\ra$.  \\begin{table} \\caption{The main parameters of the $\\Lambda$CDM simulations considered in the present paper. The normalisation of the RSI simulations follows from the results of Spergel \\etal (2003). The spectral index $n_s$ and its derivative $\\dd\\ln n_s/\\dd\\ln k$ are given at $k=0.07\\hmmpc$. The mass of a dark matter particle is about $8\\times 10^7\\mdh$ in all the simulations.} \\vspace{1mm} \\begin{center}% \\begin{tabular}{ccccccc} \\hline & $\\om{m}$ & $\\om{\\Lambda}$ & h & $\\sigma_8$ & $n_s$ & $\\dd\\ln n_s/\\dd\\ln k $ \\\\ \\hline S1 and S2  & 0.30 & 0.70 & 0.70 & 0.90 & 1 & 0 \\\\ R1 and R2  & 0.31 & 0.69 & 0.71 & 0.84 & 0.93 & -0.03 \\\\ \\hline\\hline \\end{tabular} \\end{center} \\label{table1} \\end{table}  \\begin{figure*} \\resizebox{0.45\\textwidth}{!}{\\includegraphics{fig1a.eps}} \\resizebox{0.45\\textwidth}{!}{\\includegraphics{fig1b.eps}} \\caption{ {\\it Left panel}~: The particle distribution in a slice of thickness $\\Delta h=0.05\\hmpc$ through the various $\\Lambda$CDM simulations at $z=3$~: the scale-invariant realisations S1 and S2 (top panels), and the RSI realisations R1 and R2 (bottom panels). {\\it Right panel}~: The corresponding 3D power spectra $\\dmk$ computed at $z=3$ from the DM particle distribution. The dashed curves show $\\dmk$ for the S2 and R2 simulations. The dotted curves show the linear power spectrum used to seed the fluctuations.} \\label{fig1} \\end{figure*} ",
        "conclusions": "\\label{discussion}  \\begin{table} \\caption{Parameter values of the models which best fit the flux PS and PDF, assuming a temperature-dependent smoothing $\\kf\\thg^{1/2}=$const. The  values listed in the upper half of the table were obtained using only the diagonal  terms of the PDF error matrix whereas, in the lower half, they were obtained  using the full covariance matrix of the flux PDF.  The filtering scale $\\kf(1)$ is given in unit of $\\hmmpc$.  The adiabatic index has a fixed value $\\gamma=1.3$, while the mean  IGM temperature is $\\thgg=1$, 1.5 and 2. The last columns gives the  best fit values of $\\la F\\ra$, $\\sigma_8$ and the corresponding  $\\chi^2$ (for 24 degrees of freedom). } \\vspace{1mm} \\begin{center}% \\begin{tabular}{cccccc} \\hline filtering & $\\gamma$ & $\\thgg$ & $\\la F\\ra$ & $\\sigma_8$ & $\\chi^2$ \\\\ \\hline $\\kf(1)=10$ & 1.3 & 1 & 0.68 & 0.67 & 24.9 \\\\ &     & 1.5 & 0.69 & 0.72 & 21.7 \\\\ &     & 2 & 0.69 & 0.76 & 25.3 \\\\ $\\kf(1)=8$  & 1.3 & 1 & 0.68 & 0.84 & 22.4 \\\\ &     & 1.5 & 0.68 & 0.94 & 25.4 \\\\ &     & 2 & 0.68 & 0.93 & 38.7 \\\\ \\hline $\\kf(1)=10$ & 1.3 & 1 & 0.70 & 0.70 & 26.9 \\\\ &     & 1.5 & 0.71 & 0.72 & 23.9 \\\\ &     & 2 & 0.71 & 0.72 & 27.7 \\\\ $\\kf(1)=8$  & 1.3 & 1 & 0.71 & 0.83 & 23.8 \\\\ &     & 1.5 & 0.71 & 0.82 & 27.3 \\\\ &     & 2 & 0.71 & 0.88 & 38.1 \\\\ \\hline\\hline \\end{tabular} \\end{center} \\label{table3} \\end{table}  We have used N-body simulations of  variants of the $\\Lambda$CDM models to generate  synthetic spectra of the \\op forest.  We have simulated the standard $\\Lambda$CDM cosmology with scale-invariant spectral index models, as well as a running spectral index model (Spergel \\etal 2003).  The one-dimensional flux power spectrum (PS) and flux probability  distribution function (PDF) have been computed from mock catalogs, and compared to the  observational data in order to constrain the cosmological models and the  physical parameters that dictate the properties of the \\op forest.  In addition to the one- and two-point correlations, we have also computed the neutral hydrogen column density distribution and the line width distribution.  These last two statistics have only been used as a consistency check for the models which give the best fit to  the observed PS and PDF.  \\begin{figure} \\resizebox{0.45\\textwidth}{!}{\\includegraphics{fig12.eps}} \\caption{A comparison between the $\\Delta \\chi^{2}$ versus $\\sigma_{8}$ computed from the PS alone (curves with  squares), and from the PS+PDF data (curves with triangles).  $\\Delta\\chi^2$ was obtained after varying $\\sigma_8$ and $\\la F\\ra$, and marginalizing over the value of the mean flux.  The mean IGM temperature is $\\thgg=1$ (solid curves), 1.5 (long-dashed) and 2 (short-dashed). The filtering is $\\kf=10\\hmmpc\\thgg^{-1/2}$ (cf.  Section~\\S\\ref{tdep}), and the adiabatic index has a fixed value $\\gamma=1.3$. The spline interpolation was performed using the diagonal elements of the PDF covariance matrix solely.} \\label{fig12} \\end{figure}  The RSI model matches the PS somewhat better than our scale-invariant model with rms fluctuation $\\sigma_8=0.9$, but overestimates the  PDF in the range $F\\bsim 0.6$ unless the temperature is very high, $\\thgg\\bsim 2$.  A scale-invariant model with $\\sigma_8=0.9$ matches better the flux PDF but tends to overestimate $\\dfk$ on scale $k\\lsim 0.01\\skm$ when the temperature is large, $\\thgg\\bsim 1.5$.  However, the agreement with the observed flux power spectrum can be improved by lowering the value of the normalisation amplitude, $\\sigma_{8}$.  Most importantly, models which match best the PS data alone  usually do not yield a good fit to  the PDF.  The discrepancy worsens with decreasing temperature. In the case of the $\\sigma_8=0.9$, scale-invariant cosmology, the disagreement with the  observed PDF is particularly severe for $\\thgg\\bsim 2$. We also computed the line statistics for some of our best fit  models. We found that they all reproduce successfully the slope  and normalisation of the column density distribution $f(\\nhi)$, as  well as the overall line-width distribution $f(b)$. The simulated line-width distribution is probably affected by resolution effects. Yet the  constraints on  the primordial power spectrum are insensitive to the details of  the line-width statistic.  Degeneracies among the parameters diminish the ability of the data to constrain the shape and normalisation of the primordial  power spectrum (e.g. Zaldarriaga, Hui \\& Tegmark 2001; Zaldarriaga, Scoccimarro \\& Hui 2003). In this respect, we have found that the pairs  $\\gamma$ - $\\kf$ and $\\thg$ - $\\sigma_8$ are degenerate when we marginalize  over the mean transmitted flux. To proceed further, we have assumed  that the filtering wavenumber $\\kf$ is related to the IGM temperature according to $\\kf\\propto\\thgg^{-1/2}$. The reasons which motivate this choice are discussed in Appendix~\\S\\ref{filtering}.  Under that assumption, a chi-squared minimization indicates that the normalisation amplitude is likely to be $0.7\\lsim\\sigma_8\\lsim 0.9$ for a reasonable reionization scenario and temperatures $1\\lsim\\thgg\\lsim 2$. For   $\\thgg\\bsim 1.5$, the data favour models with normalisation  $\\sigma_8\\lsim 0.8$ rather than models with $\\sigma_8\\bsim 0.8$, the later slightly overestimating $\\dfk$ in the range $k\\lsim 0.01\\skm$.  It should also be noted that taking into account the full error matrix  of the flux PDF systematically lowers the best fit values of $\\sigma_8$  by a few percents.  Our results  suggest that the constraints inferred from measurements of the flux power spectrum alone might be  biased. Most previous estimates  of  $\\sigma_{8}$ from \\op absorption measurement relied on the flux PS alone (e.g. Croft \\etal 1999,  2002; McDonald \\etal 2000, 2004b; Viel, Haehnelt \\& Springel 2004; see also Viel, Weller \\& Haehnelt 2004; Seljak \\etal 2004.  In  light of our results however, it is unclear whether their best fit  models also match the PDF of the flux (and other statistics of the  \\op forest as well). To emphasize the importance of combining several  statistics  we compare in Fig.~\\ref{fig12} the constraints on $\\sigma_{8}$  obtained from the PS data alone (curves with square) to those from from the whole PS+PDF data set (curves with triangles). We  assumed a  filtering wavenumber of $\\kf=10\\thgg^{-1/2}\\hmmpc$.  The curves $\\Delta\\chi^2$ were calculated for a mean IGM temperature $\\thgg=1$ (solid curves), 1.5 (long-dashed) and 2 (short-dashed), but for a fixed adiabatic index $\\gamma=1.3$. Note that the best fit values of $\\sigma_8$ in the case of the PS data only are consistent with those inferred in previous works.  Fig.~\\ref{fig12} clearly demonstrates that the inclusion of the PDF in the minimization has a large impact on the normalisation amplitude, decreasing the best fit value of $\\sigma_8$ by about 20\\%. Consequently, the analysis of Viel, Haehnelt \\& Springel (2004)  and McDonald \\etal (2004b) may  overestimate the  normalisation amplitude. Furthermore, including the PDF significantly improves the  confidence on $\\sigma_8$ as we can see from the shape of $\\Delta\\chi^2$. Finally, as we discussed in Section~\\S\\ref{errors}, it  also noticeably reduces the sensitivity of the results to the mean flux level $\\la F\\ra$.  A joint fit of the flux power spectrum and PDF can thus greatly improve over the standard analysis based on the flux power spectrum  alone.  Although including the flux PDF tends to lower the  best fit value of $\\sigma_{8}$, overall our result of $\\sigma_{8}=0.85-0.95$ is consistent with the studies of Viel, Haehnelt \\& Springel  (2004), and McDonald \\etal (2004b). Both of these studies were based on flux measurements at several redshifts (McDonald \\etal 2004b use eleven redshift bins spanning the range $2.2\\leq z\\leq 4.2$). They also relied on a more detailed  treatment of the gas physics by means of hydrodynamical simulations. Our assumption of uniform smoothing  does not account for the dependence of the smoothing on the local IGM density and temperature.  Yet the exact relation  between the gas and dark matter depends on  the rather poorly constrained thermal and reionization history of the universe.  Simulations incorporating a wider variety of  reionization histories are needed for a better understanding of the systematics in inferring the cosmological parameters from the  forest.   It is also unclear whether the currently available hydrodynamical simulations give  reasonable flux PDFs.  We have demonstrated that a joint analysis of the forest can potentially  yield more robust estimates of the cosmological parameters than  those inferred from the flux power spectrum alone. Here we have included mainly the PDF and the flux power spectrum. But other statistical measures can also be used.  The flux bispectrum (e.g. Mandelbaum \\etal  2003) offers a promising possibility as it should be less sensitive to systematics than the PDF."
    },
    "0410/astro-ph0410332_arXiv.txt": {
        "abstract": "We compute the effective values of apparent transverse velocity and flux boosting factors for the VLBI radio knots of blazar jets, by integrating over the angular distributions of these quantities across the widths of jets with finite opening angles but constant velocities. For high bulk Lorentz factors ($\\Gamma > 10$) variations across the jet can be quite large if the opening angle, $\\omega$, is even a few degrees on sub-parsec scales. The resulting apparent speeds are often much lower than those obtained from the usual analyses that ignore the finite jet opening angles.  We can thus reconcile the usually observed subluminal or mildly superluminal speeds with the very high  ($\\gtrsim 20$) $\\Gamma$ factors, required by the inverse Compton origin and rapid variability of TeV fluxes, as well as by intraday radio variability. Thus it is possible to associate the VLBI radio knots directly with shocks in the ultra-relativistic main jet flow, without invoking very rapid jet deceleration on parsec scales, or extremely unlikely viewing angles. ",
        "introduction": "In order to avoid excessive photon--photon collisions, the highly variable TeV emission from blazars has been interpreted in terms of inverse Compton radiation emerging from ultrarelativistic jets with bulk Lorentz factors, $\\Gamma \\sim 15-100$ (e.g., Krawczynski, Coppi \\& Aharonian 2002; Piner \\& Edwards 2004; Ghisellini, Tavecchio \\& Chiaberge 2004). Earlier, very high  values of $\\Gamma$ ($\\sim 100$) were inferred from the intraday radio variability (IDV) which is a common feature of blazars (Wagner \\& Witzel 1995; Kraus et al.\\ 2003).  It now appears that the bulk of this IDV may arise from interstellar scintillations (e.g., Kedziora-Chudczer et al.\\ 2001; % Bignall et al.\\ 2003). Nonetheless, the required microarcsecond angular sizes of the scintillating components would still need $\\Gamma > 30$, and probably substantially larger, in order to reconcile apparent brightness temperatures, $T_B \\sim 10^{13-14}$K (Blandford 2002; Rickett et al.\\ 2002) with the canonical limit of $T_B < 10^{12}$K needed to avoid the inverse Compton catastrophe (e.g., Kellermann \\& Pauliny-Toth 1969). Direct evidence for high values of $T_B$ comes from the VSOP survey (Horiuchi et al.\\ 2004). Standard models for gamma-ray bursts also invoke bulk ultrarelativistic jet flows with $\\Gamma \\sim 100-1000$ (e.g., % Sari, Piran \\& Halpern 1999; M{\\'e}sz{\\'a}ros 2002).  In contrast, the only direct probe of extragalactic jet motion, namely the radio knots detected by Very Long Baseline Interferometry (VLBI), reveal typical proper motions corresponding to apparent % speeds, $v_{app} \\equiv c\\beta_{app}$, of 5--10$c$ or less (Jorstad et al.\\ 2001; Cohen et al.\\ 2003; Giovannini 2003). % Recent measurements have shown that roughly one-third to one-half of the VLBI components measured in TeV blazars are found to be subluminal or even essentially stationary (Piner \\& Edwards 2004; Giroletti et al.\\ 2004). As emphasized by Piner \\& Edwards (2004), the apparent lack of relativistically moving shocks that are assumed to be responsible for the high energy flaring activity, and are convincingly seen in most EGRET blazars (Kellermann et al.\\ 2004), % is intriguing.  One possible explanation for the slow VLBI components of the TeV blazars would be a dramatic deceleration of the jet between sub-pc and pc-scales (e.g., Georganopoulos \\& Kazanas 2003); but then the problem becomes to understand how such a deceleration is avoided in the case of EGRET blazars for which distinctly higher apparent speeds are measured (Piner \\& Edwards 2004; Kellermann et al.\\ 2004). Essentially perfect alignment (to within $1^{\\circ}$) of the TeV blazar jets in our direction is another possible way out; however, given their substantial number, angles of a few degrees are much more likely (Piner \\& Edwards 2004). Other attempts to explain the relatively slow apparent motions of the knots involve postulating a spine and sheath geometry for the jets as was proposed for FR I sources (e.g., Sol,  Pelletier \\& Asseo 1989) % such that the radio knots are frequently associated with the slower moving outer layer (Komissarov 1990;  Laing et al.\\ 1999; Chiaberge et al.\\ 2000; Ghisellini et al.\\ 2004). Some support for this picture comes from the limb-brightening marginally observed in a few parsec-scale jets (e.g., Giovannini 2003; Giroletti et al.\\ 2004).  Here we shall argue that the observed mildly superluminal, or even subluminal, speeds of the parsec-scale VLBI knots do not necessarily demand that the knots be associated exclusively with an outer sheath; the knots may well be formed by shocks occurring in the energetic, ultra-relativistic flow of the spine. Nor is any overall slowing, or an extraordinarily close alignment, required to account for the observed substantial fraction of slow VLBI knots. Since the jets on the parsec scale are likely to be a few degrees wide, if they are pointed close to our direction and have a high bulk Lorentz factor, one expects very large variations of the apparent velocity and Doppler boosting across the face of the jet or blob. The calculations we present here show that when the distribution of apparent velocity across the jet cross-section is weighted by the corresponding Doppler boosting factor distribution, as is appropriate when the jet cannot be transversely resolved, there can be a marked overall reduction in the resulting apparent velocity, compared to the usually  assumed case where both the Doppler boosting factor and the apparent velocity are taken to be constant across the jet.  Therefore, the Lorentz factors, $\\Gamma \\gtrsim \\beta_{app}$, usually inferred from the VLBI motion data, often may be gross underestimates of the actual bulk (or pattern) Lorentz factors. ",
        "conclusions": "We have presented  computations indicating that a significant reduction in the apparent motion can be expected for VLBI radio knots associated with an ultrarelativistic jet by taking into account the jet's opening angle.  On the pc-scale which could be identified with the collimation regime, the jets are likely to be at least several degrees wide; they may be substantially wider, as indicated for the best resolved case of M87 (Junor, Biretta \\& Livio 1999).   For this situation, we find that the apparent transverse velocity peaks at lower values, and these peaks can occur at significantly greater  angles to the line of sight than they do for the usually assumed case of an infinitesimally small jet opening angle on parsec scales (Fig.\\ 1). These trends become stronger for high $\\Gamma$, for then $v_{app,w}$ is sharply peaked around some $\\theta > 1/\\Gamma$. A reversal of the sign of contributions to $v_{app,w}$ arising from some parts of the jet's cross-section occurs if $\\theta < \\omega/2$. The resulting cancellation, which is further enhanced because of the sharply declining Doppler boost with $\\theta$, can often lead to a fairly drastic reduction in the apparent velocity of the knots, as compared to the canonical peak value of $\\beta_{max} \\simeq \\Gamma$ (Fig.\\ 1).  In the core dominated samples that are characteristic of BL Lacs, the usual expectation, based on $\\omega = 0$ jets, is to find sources with $v_{app}$ widely distributed up to $\\Gamma c$, but with the distribution actually skewed toward the higher part of that range (as shown in Fig.\\ 1; see also Vermeulen \\& Cohen 1994). Thus, when the $v_{app}$'s measured for VLBI knots are frequently $< 10$c,  one of the following conclusions discussed in $\\S$1 is usually drawn:  bulk $\\Gamma$ values are modest; pattern velocities due to shock motions are  slower than the jet flow; the VLBI knots are associated with a slower sheath of the jet; the angle to the line of sight is extremely small. While all these are possible, here we have shown that none is necessary; rather, if the jet has a modest full-opening angle ($\\omega$) on the scales probed by VLBI, then there is a very large reduction in the probability of measuring apparent velocities approaching $\\Gamma c$.  For instance, even for the extreme case of $\\Gamma = 100$ and a modest jet opening angle, $\\omega = 5^{\\circ}$, over 73\\% of the radio components would be detected with $v_{app} < 10c$, while for $\\omega = 10^{\\circ}$, over 87\\% would fall into this category. Over 41\\% (for $\\omega = 5^{\\circ}$) and over 69\\% (for $\\omega = 10^{\\circ}$) would actually be seen as subluminal sources. Similarly, for  $\\Gamma = 50$ and $\\omega = 5^{\\circ}$, over 64\\% of the sources would be detected with $v_{app} < 10c$, the median value is $\\beta_{app} = 6$, and still some 15\\% would appear as subluminal sources.  Therefore, the predominance of marginally superluminal or even subluminal VLBI knots among TeV blazars does not imply that these radio knots cannot be physically associated with their ultrarelativistic jets. Instead, a combination of  high  $\\Gamma$ factors and modest jet opening angles  can just as well explain the preponderance of low $v_{app}$ values. At the same time, high $\\Gamma$ factor jets ($> 15$) are needed by the standard one-zone models to  efficiently produce the TeV photons by inverse Compton scattering (such relatively modest values emerge from models only when de-reddening of the TeV spectrum by the IR background is ignored) and higher values ($\\Gamma > 40$) are usually required when the TeV spectrum is appropriately de-reddened (e.g., Krawczynski et al.\\ 2002; Ghisellini et al.\\ 2004). Only multi-zone (rapidly decelerating or spine-sheath) models can give satisfactory fits to TeV blazar spectra with $\\Gamma \\sim 15$ even with de-reddening (Georganopoulos \\& Kazanas 2003; Ghisellini et al.\\ 2004).  Clearly, the applicability of our interpretation is not restricted to TeV blazars.  Relatively slow proper motions ($v < 3c$) are now known to occur frequently among the radio selected samples of quasars (for instance, the 2-cm VLBA survey, Kellermann et al.\\ 2004). In our picture, such slow motions would not be at variance with the values of $\\Gamma > 30$ which seem to be needed to explain the scintillating radio components of IDV blazars (e.g., Blandford 2002). The values of $\\theta$ for all blazars should be small, but are more likely to be a few degrees rather than less than 1 degree, as shown by the Monte Carlo simulations of Lister \\& Marscher (1997) and by the detection of quiescent x-ray emission (Giebels et al.\\ 2002), so explaining the  slow apparent motions through extremely precise jet alignment is a less attractive alternative.  We recall that essentially all models for gamma-ray bursts also invoke very high $\\Gamma$ factors for their jets (e.g., M{\\'e}sz{\\'a}ros 2002). If the processes that accelerate jets on these very different scales are basically similar, then this observation from GRBs may provide further support for the high $\\Gamma$ scenario in quasars and blazars (e.g., Kundt \\& Gopal-Krishna 2004).  Although our picture explains the preponderance of slow velocities, there is always a tail to the computed $v_{app}$ distribution which extends up to a substantial fraction of $\\Gamma c$. While the highest well measured values of $v_{app}$ do reach up to 35--40$c$ (Cohen et al.\\ 2003; Kellermann et al.\\ 2004), there is a strong possibility of bias against picking out VLBI components showing anomalously high apparent velocities.  The difficulty in tracking such fast moving knots in the usually sparsely temporally sampled databases is compounded by their expected rapid fading.  Upper limits on $\\omega$ for powerful blazars can be set, since the total jet power, $L_j$, is related to the {\\it inferred} bolometric luminosity $L_{bol}$ via $L_j = (L_{bol}/\\epsilon)(\\Omega/4 \\pi) = L_{bol}\\omega^2/\\epsilon$, with $\\epsilon$ the efficiency of converting jet power into radiation.  Since $\\epsilon < 0.1$ is expected  and $L_{j}$ should not exceed the Eddington limit, we have $\\omega < 0.37[(\\epsilon/0.1)(M_{BH}/10^9M_{\\odot})/(L_{bol}/10^{48}{\\rm erg~ s}^{-1})]^{1/2};$ here $\\omega$ is in radians and we have scaled $L_{bol}$ and the supermassive black hole mass, $M_{BH}$, by typical values. Unless $\\epsilon < 0.01$, $\\omega > 5^{\\circ}$ should be allowed for even very powerful blazars.   It is interesting to examine our picture in the context of the ``spine + sheath'' model of the jet, e.g., as proposed by Chiaberge et al.\\ (2000) to bring the  data into accord with the orientation based unification model (e.g., Gopal-Krishna 1995; Urry \\& Padovani 1995). By considering the nuclear emission in radio, optical and x-ray bands from BL Lacs and their presumed misaligned counterparts, the FR I radio galaxies, they inferred that a mildly relativistic ($\\Gamma = 1 - 2$) sheath component is the prime contributor to the observed nuclear emission from FR I RGs. At the same time they showed that the spectral energy distribution of the beamed counterparts, i.e., BL Lacs, all the way from radio to $\\gamma$-rays can be modeled in terms of a relativistic spine component moving with a bulk-Lorenz factor $\\Gamma \\simeq 15-20$. Still higher $\\Gamma$ values have been estimated in several other studies, particularly for TeV blazars, as mentioned above. In this work we have only discussed the spine component and tried to address the question: can such large $\\Gamma$ values be reconciled with the statistics of superluminal motion of the VLBI knots of blazars/BL Lacs, which suggest a typical $\\Gamma \\sim 3$ (e.g., Piner \\& Edwards 2004; Kellermann et al.\\ 2004)? We have argued that the discrepancy can be resolved by considering a modest opening angle ($\\sim 5^{\\circ}$) for the jet/spine on parsec scales. The blazar observations can then be understood without invoking a rapid jet deceleration on such small scales, or relegating the VLBI knots exclusively to a slower sheath (and thus totally decoupling them from the shocks occurring within the spine), or postulating extremely unlikely viewing angles for the jet.  The main question addressed here can be rephrased as follows: if indeed the fast spine component alone is relevant for the observed blazar emission, then how large a typical Lorentz factor can be reconciled with their VLBI data? We have argued that even when the VLBI knots are associated with shocks in the spine itself, their usually observed modest speeds ($v_{app} \\sim 3-5c$) would not be inconsistent with $\\Gamma > 30 - 50$ of the spine, provided one takes into account a $5 - 10^{\\circ}$ opening angle of the spine on the pc-scale. Thus, the VLBI results can be reconciled with the TeV and IDV observations indicating such large $\\Gamma$ factors (\\S 1).  One prediction of this picture is that when adequate resolution and sensitivity is achieved so as to be able to transversely resolve the fastest moving VLBI scale knots, then different portions of those knots would sometimes evince  different apparent velocities because variations across the jet opening angle could then be detected.  If sufficient dynamic range also becomes available, such components might be seen to fragment, or, perhaps more likely, appear to be smeared out quickly."
    },
    "0410/astro-ph0410104_arXiv.txt": {
        "abstract": "{The lobes of radio-loud AGNs are huge reservoirs of energy so even if the central engine turns off, the radio source is still observable for a substantial amount of time. It is quite obvious, however, that the ignition of a new period of activity can occur before the lobes have faded completely. Restarted activity has been observed in several giant radio galaxies but until recently, it has hardly ever been observed in more compact sources. Here we present an early result of a pilot VLBA programme targeted at compact core-dominated triples. It is found that the milliarcsecond structure of the central components in that type of objects normally take the form of core-jets, which in many cases do {\\em not} point toward the relic lobes. Therefore, it looks as if the phenomenon of the restarted activity can occur regardless of the source size and it is often accompanied by a repositioning of the central engine. } ",
        "introduction": "The large linear sizes of the majority of radio sources, which, for some radio galaxies termed Giant Radio Galaxies (GRGs), can be more than 1\\,Mpc, are a consequence of the constant growth of the radio structures during an appreciable period of time. According to \\citet{al87} and \\citet{liu92} the spectral ages of extended radio sources are $\\sim$$10^7$ to $10^8$~years. However, even in the case of a GRG the spectral age is two orders of magnitude less than the age of the galaxy itself which, in general, can be of the same order as the age of the Universe; the Milky Way Galaxy is a good example of this \\citep{pasq04}. It seems that the phenomena collectively called `activity' are episodic --- galaxy evolution can be interrupted at some stage and then it can continue after a certain period of time.  If the energy supply from the central engine to the hotspots and the lobes cuts off, the radio source is still observable for a substantial period of time: $10^8$~yr \\citep{kg94}. Therefore, it is quite likely that a new epoch of activity can occur before the lobes have faded completely. The observable effect of this would be the presence of a new, bright component(s) located in or straddling the centre of a larger, double-lobed relic structure. The signature of the renewed activity is most obvious if it takes the form of a smaller, double-lobed radio source giving rise to the so-called double--double structure \\citep{schoen00a,schoen00b}. J0116$-$473 \\citep{sarip02} and PKS\\,B1545$-$321 \\citep{sarip03} are also clear examples of double--double radio galaxies (DDRG).  Alternatively, `restarted activity' may result in `X-shaped' radio sources \\citep[and references therein]{helgephd} such as 3C\\,223.1 or 3C\\,403 \\citep{dt02}. It can be argued that these two sources resemble DDRGs such as J0116$-$473 or PKS\\,B1545$-$321 except for a misalignment between the inner (active) and the outer (inactive) parts. The mechanism that triggered the development of the new structure in these two kinds of objects could be the same but that DDRGs are just `special cases' in which the misalignment is (close to) zero.  In principle, different mechanisms could be responsible for activity renewal in particular types of objects, especially if they depend on the presence/absence of the misalignment between the old and the new structures and the magnitude of the ratio of those sizes. However, mergers seem to be the most `natural' explanation of the activity. A review of other possible mechanisms has been given by \\citet{schoen00a}.  All the objects in the examples quoted above are large with linear sizes of the order of several hundreds kpc. It would appear that radio sources have to be quite old for activity renewal regardless of their morphological features. The apparent lack of observable double--double or X-shaped structures in small-scale (i.e. young) sources could mean that there are inherent limits in the mechanism(s) of activity re-ignition known to exist in GRGs. This could make it rare if not impossible for such phenomenon to take place in compact objects. Alternatively, the physical conditions inside the cocoon might not favour the development of the inner lobes for a long time after the initial burst of the activity. According to \\citet{kaiser00} that timescale could be up to $\\sim$$5 \\cdot 10^7$ years so, even if the activity in a source younger than that actually ceased and then restarted, such events would remain unobservable. ",
        "conclusions": "The phenomenon of misalignments between milli\\-arcsecond- and arcsecond-scale structures in radio sources is well known and has, in fact, been thoroughly investigated --- see \\citet{appl96} for a review. Based on a large collection of observational material, they have concluded that, although it is possible to apply a helical jet model to individual sources, it is difficult to find a single mechanism responsible for the observed distribution of misalignments indicated by a statistically significant peak at $\\Delta PA\\simeq 90\\degr$. This peak is particularly well defined for BL~Lac objects. One way to interpret the misalignments is to adopt a complicated twisted helical jets model. This has been done for example by \\citet{arw99} for the quasar TXS\\,1055+018 where the jet is pointing exactly perpendicular to the whole structure of the radio source. More recently, \\citet{cass02} carried out VLBI observations of several BL~Lac objects in 1-Jy Sample, in order to investigate the spatial evolution of radio jets from a few tens to hundreds of milliarcseconds, and to search for helical jets in this class of sources. Their EVN and MERLIN observations of bent radio jets in BL~Lac objects do not lend firm support to the explanation of misalignments by means of helical jets.  Thus, although an interpretation of misaligned sources based on a twisted jets concept can be argued, adopting a scenario of re-ignition with a simple change of jet axis leads to a straightforward and natural explanation without complex modelling. Re-ignition becomes even more plausible when the overall linear size of the source in question is small. In the case of TXS\\,1055+018, which has a size of 150\\,kpc\\,$h^{-1}$, a helical model might be viable. However, in the case of J1708+0035 which is 15 times more compact there is simply not enough space for the helix to be squeezed into it.  Finally, we would like to suggest that misaligned sources could actually be compact X-shaped sources in which one `arm' of the cross lies almost in the sky plane and the other `arm', namely the one pertinent to the reborn radio source, is highly beamed toward the observer. But do compact X-shaped sources exist? In fact, such a class has not been recognised so far. Nevertheless, we {\\em have} serendipitously found a map of a compact X-shaped source (TXS\\,0229+132) in the literature \\citep{murphy93}. According to Browne (priv. comm.) it is very unlikely that the strange structure of this source could result from gravitational lensing. It is, however, easily imaginable that if one arm of its cross-like structure was pointed towards the observer, the image ``distorted'' by Doppler boosting would show the pattern we have here: a core-dominated triple. We claim that, taking into account the difficulties of the helical jet model when trying to explain large misalignments \\citep[and references therein]{appl96,cass02}, the restarted activity concept is a very competitive alternative because of its simplicity and plausibility."
    },
    "0410/astro-ph0410381_arXiv.txt": {
        "abstract": "I discuss the effect of the formation of a black hole on a (close) binary and show some of the current constraints that the observed properties of black hole X-ray binaries put on the formation of black holes. In particular I discuss the evidence for and against asymmetric kicks imparted on the black hole at formation and find contradicting answers, as there seems to be evidence for kick for individual systems and from the Galactic $z$-distribution of black hole X-ray binaries, but not from their line-of-sight velocities. ",
        "introduction": "We can study black holes only when they interact with other stars or material. For stellar mass black holes the only way to study them is when they reside in binaries and in particular when they accrete material from their (low-mass) companion stars in black-hole X-ray binaries (BHXBs). These binaries can be used to get information about the formation of black holes, of course with the caveat that strictly speaking they only constrain the formation of black holes \\textit{in binaries}. The formation of black holes from massive stars is poorly understood \\citep[e.g.][]{fk01}. The questions we would like to answer are:\\\\ i) Which stars form black holes? (What is the minimum mass of a main sequence star that forms a black hole, and does this depend on binarity/rotation etc?)\\\\ ii) Do black holes form in either direct collapse, a supernova explosion or either?  \\\\ iii) How much mass is lost during the black hole formation?  (which is important for the chemical enrichment of the interstellar medium) \\\\ iv) Is the black hole formation symmetric or asymmetric?\\\\ Here I will discuss only the last three questions. ",
        "conclusions": "By analyzing the properties of BHXBs a number of things can be learned about the formation of black holes. The first is that some BHs are formed in a supernova, with (quite some) mass loss. Secondly, the $z$-distribution BHXBs and NSXBs are similar, which naively suggests kicks are imparted on black holes as well as on neutron stars but this needs detailed investigation.  The measurement of 3D velocities opens the possibility to study BH formation in \\textit{detail} and the first results indicate that black holes were formed in a supernova and received an asymmetric kick. However, line-of-sight velocities of BHXBs are generally low, suggesting not all black holes receive a kick."
    },
    "0410/astro-ph0410348_arXiv.txt": {
        "abstract": "{We present the first results of a photometric survey of poorly studied nova remnants in the Northern Hemisphere. The main results are as follows: \\object{DM Gem} shows a modulation at 0.123 d (probably linked to the orbit) and rapid variations at $\\sim 22$ min. A moderate resolution spectrum taken at the time of the photometric observations shows intense \\hel{ii}{4686} and Bowen emission, characteristic of an intermediate polar or a \\object{SW Sex} star. Variability at 0.127 d and intense flickering (or quasi-periodic oscillations) are the main features of the light curve of \\object{CP Lac}. A 0.1-mag dip lasting for $\\sim 45$ min is observed in \\object{GI Mon}, which could be an eclipse. A clear modulation (probably related to the orbital motion) either at 0.179 d or 0.152 d is observed in the $B$-band light curve of \\object{V400 Per}. The results for \\object{CT Ser} point to an orbital period close to 0.16 d. Intense flickering is also characteristic of this old nova. Finally, \\object{XX Tau} shows a possible periodic signal near 0.14 d and displays fast variability at $\\sim 24$ min. Its brightness seems to be modulated at $\\sim 5$ d. We relate this long periodicity to the motion of an eccentric/tilted accretion disc in the binary. ",
        "introduction": "Cataclysmic variable stars (CVs) exhibit their most extreme outbursting behaviour in classical novae. These explosive events are the consequence of thermonuclear runaways in the hydrogen envelope of the white dwarf as a result of continuous mass accretion on to its surface \\citep{starrfieldetal76-1}. It takes decades for a nova to return to its quiescent (prenova/postnova) state, where most of them behave like nova-like CVs. Comprehensive reviews on classical novae in particular and CVs in general can be found in \\cite{bode+evans89-1} and \\cite{warner95-1}, respectively.  So far there is a large number of nova remnants that have received little or no attention at all due to their faintness. Because of the lack of study, a crucial parameter such as the orbital period is currently undetermined for the majority of these quiescent novae. Many of these systems might show the hallmarks of the orbital motion in their optical light curves in the form of pure orbital modulations, superhumps, ellipsoidal modulations, eclipses, etc. Thus, time-resolved photometric monitoring provides an easy way to obtain a reliable estimate of their orbital periods.  The main goals of this long-term project are to obtain accurate estimates of the orbital periods of poorly studied novae remnants as the necessary input for efficient time-resolved spectroscopic monitoring, and to search for photometric behaviour characteristic of the different CV sub-classes. In this paper we present the first results for six novae remnants. ",
        "conclusions": "In Fig.~\\ref{fig_allperiodograms} we show all the periodograms presented in the paper (with the exception of \\object{GI Mon}) so that the excess power for a given object can be easily spotted by comparison with the others. The likely periodicities found are summarised in Table~\\ref{table_allperiodograms}.   \\begin{figure} \\centerline{\\includegraphics[width=9cm]{1112fg15.ps}} \\caption[]{\\label{fig_allperiodograms} All Scargle periodograms presented in the paper (with the exception of \\object{GI Mon}) for easy comparison between objects. The peaks corresponding to the possible periodicities listed in Table~\\ref{table_allperiodograms} (excepting the long term variation in \\object{XX Tau}) are marked with vertical lines.} \\end{figure}  \\subsection{\\object{DM Gem}} The light curves provide a modulation at 0.123 d and possible rapid variability at $\\sim 22$ min. The continuum slope of the optical spectrum suggests the presence of an accretion disc in the system. The strength of the \\hel{ii}{4686} emission line is comparable to that of H$\\beta$. The Bowen blend is also intense, indicating the presence of a source of ionising photons. We point out that \\object{DM Gem} might be an IP or a \\object{SW Sex} star. Future  spectroscopic studies must have the appropriate time resolution to sample a possible $\\sim 20$-min modulation in the emission-line flux and enough time coverage to properly deal with a $\\sim 3$-hour orbital period.  \\subsection{\\object{CP Lac}} This nova remnant is characterised by intense short time scale variability superimposed on a 0.127-d wave. We tentatively suggest that shallow eclipses might be present. The intense \\hel{ii}{4686} and Bowen emission in its optical spectrum points towards a magnetic nature. If the 0.127-d periodicity is related to the orbital motion, \\object{CP Lac} will lie within the period range dominated by the \\object{SW Sex} stars.  \\subsection{\\object{GI Mon}} Our data are insufficient to confirm the 0.1802-d periodicity reported by \\cite{woudtetal04-1}. We recorded a shallow dip in one of our light curves which must be confirmed as an eclipse. In the case \\object{GI Mon} is an eclipsing CV, the length of our longest light curve constrains the orbital period to be $> 0.2$ d.  \\subsection{\\object{V400 Per}} Our $B$-band photometry provides a clear modulation at either 4.30 or 3.65 hours, which we link to the orbital motion of the binary.  \\subsection{\\object{CT Ser}} The light curves do not provide a clear picture. A modulation at $\\sim 0.17$ d might be present. Erratic, short time scale variability is also observed.  \\subsection{\\object{XX Tau}} Our results favour a modulation at $\\sim 0.14$ d for this nova remnant. Fast variability at $\\sim 24$ min might also be present, but it is uncertain. The most obvious feature is a brightness modulation at $\\sim 5$ d, similar to other nova-like CVs (particularly the \\object{SW Sex} stars). We attribute this long-period variation to the motion of an eccentric/tilted accretion disc in the binary.  \\begin{table} \\caption[]{\\label{table_allperiodograms}Summary of possible periodicities} \\begin{flushleft} \\begin{tabular}{lll} \\hline \\hline\\noalign{\\smallskip} Object & Periodicity (min) & Remarks\\\\ \\hline\\noalign{\\smallskip} \\object{DM Gem}   & 177.1, 22.6 & Strong \\hel{ii}{4686}. Magnetic? \\\\ \\object{CP Lac}   & 182.9       & Shallow eclipses? SW Sex? \\\\ \\object{GI Mon}   & 259.2       & \\\\ \\object{V400 Per} & 219 or 258  & \\\\ \\object{CT Ser}   & 233.3       & \\\\ \\object{XX Tau}   & 7200, 195.8, 23.7 & Disc precession? \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{flushleft} \\end{table}"
    },
    "0410/astro-ph0410454_arXiv.txt": {
        "abstract": "Hubble Space Telescope imaging of 11 fields in M31 were reduced to color-magnitude diagrams. The fields were chosen to sample all galactocentric radii to 50 kpc ($\\approx 9$ disk scale lengths, or $>$99.9\\% of the total light enclosed). The colors of the red giants at each pointing map to an abundance distribution with errors of order 0.1 dex in [M/H]. The abundance distributions are all about the same width, but show a mild gradient that flattens outside $\\sim$20 kpc. These distributions were weighted and summed with the aid of a new surface brightness profile fit to obtain an abundance distribution representative of the entirety of M31.  Since we expect M31 to have retained most of its congenital gas and subsequently accreted material, and since the present day gas mass fraction is around 2\\%, it must be a system near chemical maturity. This ``observed closed box'' is found to suffer from a lack of metal-poor stars {\\em and} metal-rich stars relative to the simplest closed-box model in the same way as the solar neighborhood.  Models modified to include inhomogeneous chemical enrichment, variable yield, or infall all fit to within the uncertainties.  As noted elsewhere, stars in the outer regions of M31 are a factor of ten more metal-rich than the Milky Way halo, ten times more metal-rich than the dwarf spheroidals cospatial with it, and more metal-rich than most of the globular clusters at the same galactocentric radius. Difficulties of interpretation are greatly eased if we posit that the M31 disk dominates over the halo at all radii out to 50 kpc. In fact, scaling from current density models of the Milky Way, one should not expect to see halo stars dominating over disk stars until beyond our 50 kpc limit. A corollary conclusion is that most published studies of the M31 ``halo'' are actually studies of its disk. ",
        "introduction": "We have knowledge of the spiral Andromeda galaxy (M31; NGC 224) and its elliptical companions M32 (NGC 221) and NGC 205 that rivals and in some ways surpasses our knowledge of the Milky Way.  Comparing the two galaxies can be a stringent test for theories of hierarchical galaxy formation. The Andromeda Galaxy has about the same baryonic-plus-dark mass as the Milky Way \\citep{2000ApJ...540L...9E,2002MNRAS.337...34G} but twice the number of globular clusters. Gas phase material comprises 2\\% of the baryonic total in M31, but about twice this in the Milky Way \\citep{vdb}. Further, the Milky Way is forming stars at a greater rate, with a factor of $\\sim$4 greater mass in ionised hydrogen. Most of the gas in M31 is concentrated in a 10-kpc ``ring of fire'' seen in molecular gas, atomic hydrogen, and dust, although OB associations appear throughout the disk. \\citet{vdk} finds that the bulge of M31 contributes 25\\% of the total light while Milky Way's bulge contributes more like 12\\%. (One should temper this result with the realization that the Milky Way emits more light per unit mass than does M31; in terms of mass the percentage contributions may be more nearly equal. Also, our own $I$-band bulge-disk decomposition discussed below yields a bulge contribution of 12\\%, not 25\\%.)  Globular cluster systems differ in the two galaxies despite both having the same number of clusters per unit mass within a factor of $\\sim$2 \\citep{vdb}. With a dividing line between metal poor and metal rich placed at [M/H]~$=-1.0$, Milky Way metal rich globular clusters are concentrated toward the center of the galaxy.  M31 has a larger number of metal-rich globular clusters than the Milky Way. As in the Milky Way, the metal rich M31 globular clusters lie preferentially toward the center of the galaxy, but unlike the Milky Way, very metal-rich globulars (of nearly solar abundance) are also found quite far from the center. Unlike the Milky Way, a significant subset of M31 globular clusters appear to have strongly disk-like kinematics \\citep{morrison}.  Halo field stars in the Milky Way have an abundance distribution similar to that of the globular clusters spatially coresident there: with a peak at [M/H]~$\\sim -1.5$ and a broad dispersion that matches the simplest of chemical evolution models, the ``Simple model,'' that assumes a closed box, full mixing, instantaneous recycling, and a constant heavy element yield in every generation of stars. Halo field stars in M31, if they do indeed belong to a halo population, are quite different.  Among other authors, \\citet{hh01,brown,rich,durrell,renda}, and \\citet{grill} find from color-magnitude diagrams that almost all of the field stars in the outer regions of M31 are more metal-rich than 47 Tucanae (at [M/H]$\\sim -0.8$), the prototypical metal-rich disk globular cluster, with mean abundances for the giants of [M/H]$\\sim -0.5$.  This is almost alarming: why should a galaxy of about the same mass have a halo abundance a factor of ten higher?  The aforementioned ``Simple Model'' of chemical evolution is a ``straw man'' (easily knocked down) model that applies to a closed box of gas that turns into stars. However, its simplicity, along with the relative complexity of the alternative chemical evolution models, makes it a compelling starting point. Furthermore, taken as a whole, M31 inside a galactocentric radius of 50 kpc must be close to a closed box in that it does not seem reasonable that it {\\em permanently} lost more than a pittance of its total gas mass to intergalactic space. Models support near-total gas retention for high-mass galaxies in non-cluster environments \\citep{gov04}. Given that almost all of its gas is either still in the M31 disk (2\\% by mass) or has been turned into stars (98\\% by mass), then on a global level M31 is a closed box or close enough to a true closed volume that outflow should be negligible.  But M31 is a closed box only in the outflow direction. Infall in M31 could have occurred throughout its history, perhaps episodically where each episode goes nearly to chemical completion before the next wave, or perhaps steadily where the M31 gas fraction has been monotonically decreasing to its minimum at the present epoch. In the latter case, the Simple model might be expected to apply fairly well to the M31 closed box without modification. In the episodic case, one would have a situation usually termed inhomogeneous chemical evolution. In inhomogeneous chemical evolution, different patches of the galaxy mix only slightly or not at all or are allowed to mix only between enrichment episodes.  A final aspect of M31 that affects the discussion of its formation and chemical evolution is the nature of its stellar halo. M31 does possess a system of globular clusters that is spatially and kinematically spheroidal, with a net rotation of only $80\\pm 20$ km s$^{-1}$ \\citep{huchra}. \\citet{perrett} finds a net rotation of $138\\pm 13$ km s$^{-1}$ and velocity dispersion of $156\\pm 6$ km s$^{-1}$ for a sample of 321 clusters. \\citet{racine} finds a globular cluster surface density profile of $R^{-2}$ in the range $6 < R$(kpc) $< 22$, but a steeper drop off, or even a cutoff, beyond this range. \\citet{vdb69} finds that the most metal-rich globular clusters all have velocities of $\\pm 100$ km s$^{-1}$ from the projected disk velocity at their locations, indicating a more disk-like motion for the metal-rich end of the cluster population. More extensive data \\citep{morrison} indicate that a significant fraction of M31 globular clusters have thin-disk kinematics, are sprinkled over the entire projected disk, and have abundances that span the entire metallicity range. The M31 globular clusters have a spread of integrated colors that covers about the same abundance range as Milky Way globular clusters, but contains a more equal number of clusters in each metallicity bin \\citep{reed94}.  The surface brightness of field stars in the outer portions of M31 as inferred from star counting follow, at least roughly, an exponential $R^{1/4}$ profile \\citep{vdb}. This is roughly $\\rho (R) \\propto R^{-5}$ in the outer regions, as opposed to $R^{-3}$ as one would infer for globular clusters \\citep{racine}. But the continuity of the inner and outer parts of the minor axis (illustrated, for example, in \\citet{vdb}, figure 3.7) does not demand a ``halo'' that is separate from an ``outer disk''. In fact there is no evidence of a halo at all from the light profile of the galaxy. Support for this viewpoint can be found from the planetary nebulae kinematics of \\citet{hurley}, who say, ``If M31 has a non-rotating, pressure-supported halo, we have yet to find it, and it must be a very minor component of the galaxy.''  In following sections we discuss HST photometric observations that sample galactocentric radii within 50 kpc, then we present the inferred abundance distributions. (What we refer to as abundance distributions are called by many ``MDFs,'' or metallicity distribution functions, and are simply the fractions of stellar mass observed over the range of abundance.) The observed distributions are then weighted and summed to simulate a globally integrated ``closed box'' abundance distribution for M31, which is then compared to solar neighborhood observations and model predictions. We conclude with a discussion of the constraints imposed upon the chemical evolution and formation of the Andromeda galaxy and a discussion of the relative dominance of disk vs. halo in the outer parts of the galaxy. ",
        "conclusions": "\\subsection{Closed Box Models}  In summary, the comparisons of the observed M31 closed box abundance distribution with chemical evolution models showed that, while the Simple model is too broad, all kinds of slightly modified chemical evolution schemes fit the narrower observed abundance distribution of closed-box M31 stars. Combined with the result of \\citet{wdj} that elliptical (i.e. spheroid-dominated) galaxies also have a narrow abundance distribution compared to the Simple model, the ``narrowness'' of the abundance distribution appears to be a universal phenomenon in all large galaxies, not just for disk galaxies. This is circumstantial evidence that the abundance distribution is generically narrower than the Simple model suggests.  Since outflow is ruled out for closed boxes, the remaining classical solutions to this problem \\citep{at76} are (1) prompt initial enrichment, (2) infall without outflow, (3) dropping the assumption of instantaneous recycling, i.e. allowing for time-delayed yields and allowing inhomogeneous enrichment. Our $Z$-dependent yield is an implementation of (1), and it works fairly well except that the parameterisation is not physical.  Scheme (2), infall, easily matches almost any abundance distribution with suitable fine-tuning of yield, mixing, and infall parameters, as we show with the full-mixing \\citet{hh02} parameterization.  The problem with infall models as applied to galactic closed boxes is that the output abundance distribution is very sensitive to the infall timescale, so that the infall timescale would have to be nearly the same for ellipticals as it is for spirals since they have very similar abundance distributions. This objection is nullified if the infall timescale is always directly proportional to the star formation rate to about a factor of two in all star formation environments.  Scheme (3) was partially explored with inhomogeneous models. The easiest match included substantial mixing of gas during chemical evolution.  Delayed injection of heavy elements serves to add metal-poor stars to the system, and so would make the fit worse if it were included. We are therefore left with a theme often voiced about chemical evolution; that there is a uniqueness problem. Several models with different assumptions can match the data quite well \\citep{tosi}.  \\subsection{Disk Dominance}  As a prelude to further discussion, we ask ``at what radius do we expect M31 halo stars to dominate over disk stars?'' assuming that the M31 halo resembles that of the Milky Way. Disks are usually modelled with an exponential profile of the form $I=I_0\\ {\\rm exp}(-R/h)$, where $h$ is the scale length and $I$ is a surface brightness, not a density. The Milky Way halo is found to follow a power law density $\\rho = \\rho_0 R^{-3.5}$ outside of 5 kpc \\citep{vdb}. Provided we can trust these functional forms, an exponential will fade faster than a power law, so a halo will always dominate outside some radius.  To evaluate the surface density of the halo, normalised to the local surface density of the disk, we use the results of \\citet{sandage}, integrating his Table 1 (kinematically derived densities as a function of height above the disk) and adopting local density ratios for disk:thick disk:halo of 500:30:1 (and for our purposes we simply subsume the thick disk into the disk). The problem is then fully specified except for the disk scale length $h$. Figure \\ref{figratio} shows the prediction for the $V$-band Milky Way scale length $h=4.2$ kpc \\citep{vdb}. Van den Bergh notes that estimates of 2 to 6 kpc for the Milky Way scale length $h$ exist in the literature, part of this spread being a function of the wavelength studied.  The \\citet{chen} models are also shown in Figure \\ref{figratio} exactly as integrated from the published density model (including a shallower halo density power law of $-2.5$ and disk scale length of $h=2.25$ kpc). These results indicate that, in the Milky Way, halo stars become more common than disk stars at galactocentric radii of 35 kpc $< R <$ 45 kpc.  The scale length of M31 is more secure than that of the Milky Way. \\citet{vdb} quotes $h=5.7$ kpc for the $V$-band, and our own $I$-band decomposition gives $h=5.6$ kpc [or $5.7$ kpc if the slightly longer distance of 784 kpc \\citep{stan} is used]. For a scale length of this size, and also doubling the density of the M31 halo because M31 has twice the number of globular clusters, Figure \\ref{figratio} indicates that we should {\\em not} expect to see the halo dominate at any radius that we sample, and that halo stars should become relatively numerous only outside a radius of 50 kpc from the center of M31.  This should be slightly tempered by a geometrical consideration: If the halo is spherical, a line of sight intersects the halo at an impact parameter $R_h \\approx D\\theta$, where $D$ is distance and $\\theta$ is the angle between the line of sight and the M31 nucleus. The impact parameter is the same as the galactocentric radius in the inclined disk ($R_d$) only along the major axis of the isophote. For most lines of sight, the projected disk is foreshortened so that $R_h < R_d$, and the foreshortening is at a maximum along the minor axis of an isophote \\citep{wk87}.  For gauging this possible amplification effect, we list both radii in Table \\ref{tab1} [the amplification of the halo is roughly $(R_d/R_h)^{2.5}$ assuming the halo density profile is $\\rho \\propto R^{3.5}$].  We find that fields near the minor axis can have the halo boosted by a factor of order ten from this effect. Of course, if the halo is flattened instead of spherical (as the Milky Way's is believed to be) the boost rapidly diminishes. And there is evidence from globular clusters that the halo density falls off faster than $\\rho \\propto R^{-3.5}$ outside of 30 kpc \\citep{racine}. It is possible that, with the exception of nearly edge-on spirals where projection effects dominate, the disk stars outnumber the halo stars at all radii. We note that the metallicity remains about [M/H] $\\sim -0.5$ even in fields as far as $R_h = 30$ kpc \\citep{dur04,durrell}.  If we are seeing disk stars in all fields, then explaining the relatively metal-rich mean abundance of [M/H]$\\approx -0.5$ becomes quite natural compared to trying to explain why the halo of M31 is a factor of ten more metal-rich than the Milky Way halo. If we are seeing disk stars, then the M31 halo can be presumed to be alive and well and very similar to the Milky Way halo, which, after all, is supposed to compose less than 2\\% of the mass of the Galaxy. Alive and well, but, as yet, unseen except for the metal-poor-selected stars of \\citet{rg02}.  Many authors have called the outer regions of M31 ``halo,'' but for the most part this is a term of convenience, and not indicative of true pressure-supported halo status since we are almost completely ignorant of the kinematics of these stars as of this writing. Parantheticallly, our ignorance is about to be relieved as large telescopes are beginning to measure many velocities of giant stars in the uncrowded outer regions of M31 [cf. \\citet{rg02}]. It is interesting to wonder how interpretations change if disk status is assumed. \\citet{brown} find an ``intermediate age'' subpopulation in a field at 11 kpc galactocentric radius in the plane of the sky (we derive 11.4 kpc from the coordinates given in their paper plus our assumed distance), or 35.1 kpc galactocentric radius on the extrapolated M31 disk. If this is disk, then finding non-ancient stars seems perfectly normal, and we may wonder why there are no {\\em younger} stars present. Perhaps this is evidence that star formation in the outer M31 disk was quenched after about half a Hubble time. {\\em Preliminary} Keck/DEIMOS velocities of stars around the the \\citet{brown} field indicate a velocity dispersion of 85 km s$^{-1}$ centered on the systemic velocity of M31 (J. S. Kalirai (2005), private communication). Compared to the 156 km s$^{-1}$ line-of-sight velocity dispersion of the M31 globular clusters, this is neither a dynamically cold thin disk nor halo, but probably something akin to the Milky Way's thick disk component and very much in line with our speculation of a dynamically heated ancestral disk. Recently, \\citet{chap} report disklike kinematics between 15 and 40 kpc, but with a much smaller velocity dispersion (30 km s$^{-1}$).  By counting red giants in the vicinity of M31, \\citet{ferg} established the presence of dynamical clumping or streaming at radii out to 55 kpc in an elliptical survey region shaped like an extrapolation of the disk. Some substructure detected strongly resembles merging events, but some of the features inside $\\sim$30 kpc also look like flocculent spiral structure. Companions M32 and NGC 205 project on the sky {\\em inside} these structures. The portions of the \\citet{ferg} Fig. 2 map outside $\\sim$30 kpc look rather uniform (except for the substructures) but this may be due to Galactic foreground stars that were not subtracted. The \\citet{ferg} maps certainly do not rule out that the outer parts of M31 are dominated by a disturbed and tidally shredded disk.  Our synthesis of these considerations is that M31 formed a gaseous, well-mixed disk and a Milky Way-like halo in the first half of a Hubble time. This ancestral disk need not be oriented exactly the same way as the present-day disk. Chemical enrichment proceeded in a ``disky'' way via rotational mixing, galactic fountains, and the occasional supernova wind. Over time, the gas supply ran short so the outer disk stopped forming stars. Additionally, late accreting satellite galaxies made their dynamical presence felt, adding stellar streams and tidally disturbing the disk stars that were already present. The advent of more measurements of stellar kinematics in the outer disk of M31 should reveal the presence of streams from disrupted satellites and also the structure of the postulated disturbed ancestral disk that may have a variable velocity dispersion as a function of galactocentric radius or even gaps and warps caused by strong perturbers.  In particular, if M32 has been orbiting M31 for more than a few orbits it will certainly be well on the way to clearing a gap in the ancestral stellar disk due to tidal resonance effects. Finding the size and shape of such a gap could be used to constrain M32's orbit and infer how long it has been a satellite of M31."
    },
    "0410/astro-ph0410724_arXiv.txt": {
        "abstract": "We have calculated the evolution of cosmic ray (CR) modified astrophysical shocks for a wide range of shock Mach numbers and shock speeds through numerical simulations of diffusive shock acceleration (DSA) in 1D quasi-parallel plane shocks. The simulations include thermal leakage injection of seed CRs, as well as pre-existing, upstream CR populations. Bohm-like diffusion is assumed. We model shocks similar to those expected around cosmic structure pancakes as well as other accretion shocks driven by flows with upstream gas temperatures in the range $T_0=10^4-10^{7.6}$K and shock Mach numbers spanning $M_s=2.4-133$. We show that CR modified shocks evolve to time-asymptotic states by the time injected particles are accelerated to moderately relativistic energies ($p/mc \\gsim 1$), and that two shocks with the same Mach number, but with different shock speeds, evolve qualitatively similarly when the results are presented in terms of a characteristic diffusion length and diffusion time. We determine and compare the ``efficiencies'' of CR acceleration in our simulated shocks by calculating the ratio of the total CR energy generated at the shock to the total kinetic energy that would pass through the shock over time in its initial frame of reference. For these models the time asymptotic value for this efficiency ratio is controlled mainly by shock Mach number, as expected from the aforementioned similarity in CR shocks. In the presence of a pre-existing CR population, shock evolution proceeds similarly to that for higher thermal injection rates compared to thermal leakage CR sources alone. This added contribution has little or no impact on the postshock or CR properties of the high Mach number shocks simulated. The modeled high Mach number shocks all evolve towards efficiencies $\\sim 50$\\%, regardless of the upstream CR pressure. On the other hand, the upstream CR pressure increases the overall CR energy in moderate strength shocks ($M_s \\sim {\\rm a~few}$), since it is a significant fraction of the shock ram pressure. These shocks have been shown to dominate dissipation during cosmic structure formation, so such enhanced efficiency could significantly increase their potential importance as sources of cosmic rays. ",
        "introduction": "Collisionless shocks form ubiquitously in tenuous cosmic plasmas via collective, electromagnetic viscosities. The formation process of such shocks inevitably produces suprathermal particles in addition to thermal particles with Maxwellian velocity distributions \\citep{blaeic87, Jones91, Draine93}. These nonthermal particles can be further accelerated to very high energies through the interactions with resonantly scattering Alfv\\'en waves in the converging flows across a shock, \\ie by diffusive shock acceleration (DSA) \\citep{dru83,blaeic87,maldru01}. Detailed nonlinear treatments of DSA predict that a small fraction of incoming thermal particles can be injected into the CR population, and that a significant fraction of the shock kinetic energy can be transferred to CRs \\citep[\\eg][]{ebj95, kangetal02}.  Simulations of DSA in spherical supernova remnants (SNRs) indicate that CRs can absorb up to 50\\% of the initial blast energies \\citep[\\eg][]{berzvolk97, berzvolk00}.  Support for rapid and efficient DSA in that setting has been provided by recent X-ray observations of young SNRs such as SN1006 and Cas A that indicate the presence of short-lived, TeV electrons emitting nonthermal synchrotron emission immediately inside the outer SNR shock \\citep{koyama95, allen97, bamba03}. On a galactic scale it is well known that the CR energy density is comparable to the gas thermal energy density in the interstellar medium and plays important dynamical roles in the evolution of our Galaxy. Although the Galactic CRs are commonly believed to be accelerated mostly at SNR shocks, CR acceleration is probably important in all shock-heated astrophysical plasmas.  CR populations are also indicated by extended, diffuse nonthermal emissions in some galaxy clusters \\citep[\\eg][and references therein]{min01}. One likely contribution to that CR population is large scale shocks associated with the formation of the clusters. According to hydrodynamic simulations of large scale structure formation in the universe, nonlinear structures such as pancakes, filaments, and knots are surrounded by ``external'' accretion shocks and contain complex webs of ``internal'' flow shocks including, but not limited to, accretion shocks around individual clusters and merger shocks inside clusters \\citep{min00}. Recently, \\citet{ryu03} studied the characteristics of such cosmic shocks and showed that they have a wide range of physical parameters with shock speeds, $u_s$, up to $\\sim 3000\\kms$, preshock temperatures of $10^4< T_0< 10^8$ K, and Mach numbers up to a few 100. They showed that in the present universe the mass inside nonlinear structures has been shocked approximately twice on average over cosmic time, and that shocks with $2 \\la M \\la 4$ have contributed $\\sim 1/2$ of the total energy dissipated at the shocks. Utilizing nonlinear DSA model calculations of \\citet{kangetal02}, Ryu \\etal estimated that the ratio of CR energy to gas thermal energy resulting from dissipation at cosmological shocks is $\\sim1/2$. Hence, the CR energy density could be dynamically important in the intracluster medium, just as in the interstellar medium of our Galaxy and might have had some dynamical influences on large scale structure formation.  The above ``cosmic structure shocks'' are an inevitable consequence of gravitational collapse, and are central to virialization of diffuse baryonic matter. The DSA process behind our present discussion depends, of course, on the existence of a magnetic field in the shock vicinity to mediate Alfv\\'enic turbulence that scatters the CRs. The effectiveness of the scattering at a particular particle momentum, $p=\\gamma \\upsilon m$, can be expressed in terms of the DSA diffusion time, $t_d(p) = \\kappa(p)/u^2_s$, where $\\kappa = (1/3) r_g \\upsilon  (B/\\delta B)^2$ is the spatial diffusion coefficient, $r_g = pc/(eB)$ is the particle gyroradius, $B$ is the magnetic field, and $\\delta B$ characterizes the magnitude of resonant wave magnetic fields. Particles will generally be accelerated from low energies to $E \\sim pc$ on timescales a few times $t_d(p)$ (see equation \\ref{tacc}). Efficient transfer of kinetic energy to CRs typically becomes apparent in DSA simulations by the time transrelativistic CRs are accelerated, establishing the relevant time scale for the introduction of nonlinear feedback into the shock dynamics.  Evidence is strong that magnetic fields exist in clusters at levels exceeding at least a few $\\times 10^{-7}$Gauss \\citep[\\eg][]{kron01}. Fields of that strength lead, with Bohm diffusion ($\\delta B = B$), to diffusion times for transrelativistic CRs, $t_d = (1/3) pc\\upsilon/(eB u_s^2)$ of the order of a few years for $p \\sim mc$ in fast cosmic shocks with speeds $\\sim 10^3\\kms$.  Direct evidence for microGauss-level magnetic fields where most cosmological shocks form, that is, outside individual clusters, is currently quite limited. However, there are good reasons to expect fields adequate for efficient DSA in most cosmic shocks. For example, even in the absence of any primordial seed field, magnetic fields are likely to be spontaneously generated at ``external cosmic shocks'' by some combination of the Biermann battery \\citep{kuls97,rkb98} and the Weibel instability \\citep{sch03}. Whether such seed fields are initially strong enough to reduce diffusion times to cosmologically interesting values or not, quasi-linear plasma theory and simulations show that the streaming motion of suprathermal particles can induce resonant Alfv\\'en wave generation and strong field amplification upstream of collisionless shocks \\citep[\\eg][]{Bell78,Quest88, lucek00, bell01}. \\cite{bell01} have argued that magnetic fields can ultimately be amplified to produce magnetic pressures that are a substantial fraction of the ram pressure through the shock, $\\rho_0 u^2_s$. That limit would correspond to field strengths $\\sim 4\\times 10^{-7} h (1+\\delta)^{1/2} u_{s,3}(1+z)^{3/2}$Gauss, where $h$ is the Hubble parameter in units of 100 km/sec/Mpc, $\\delta=\\delta\\rho/\\rho$ is the local overdensity ratio, $z$ is the redshift of the epoch, and $u_{s,3}$ is the shock speed in units of $10^3\\kms$. A baryon density fraction, $\\Omega_b=0.04$, is assumed. Even magnetic fields only one percent of this value would be sufficient to accelerate CRs to relativistic energies, and, hence establish nonlinear CR feedback on timescales of thousands of years in the shocks of interest.  Once shocks develop nonlinear properties in response to CR feedback, there are several important characteristics that distinguish them from more familiar gasdynamic shocks: 1) CR shocks continue to evolve over relatively long times and broaden as they do so. While full thermalization takes place instantaneously at a simple, discontinuous jump in an ``ideal'' gasdynamic shock, CR acceleration and the corresponding modifications to the underlying flow depend on suprathermal particles passing back and forth diffusively across the shock structure. These processes develop, therefore, on the diffusion time scale, $t_d$, and diffusion length scale, $x_d = t_d u_s$, as alluded to earlier. Generally, these scales are expected to be increasing functions of particle momentum, so CR acceleration and shock evolution rates slow over time. For very high energy particles DSA time and length scales often approach the age and the radius of curvature of astrophysical shocks. 2) CR diffusion upstream of the shock discontinuity leads to strong pressure gradients in a shock precursor, enhancing the total compression through the shock transition over that in the ``viscous'' subshock. The total compression through a high-Mach-number CR shock can greatly exceed the canonical value $r=(\\gamma_g+1)/(\\gamma_g-1)$ for strong gasdynamical shocks \\cite[\\eg][]{berzell99, maldv00} (where $\\gamma_g$ is the gas adiabatic index). 3) Both the effective compressibility of the combined thermal-CR plasmas and the mean CR diffusion coefficient increase over time as relativistic CRs absorb more energy and as escaping CRs remove energy from the shock structure. Details of these properties depend on the CR momentum distribution. Thus, downstream states of the CR modified shock cannot be found from simple ``shock jump conditions\", but rather have to be integrated time-dependently from given initial states or, for steady solutions, found in terms of some predefined limits to the CR spectrum (\\eg an upper momentum cutoff). 4) Energy transfer to the CRs rather than to the thermalized gas in the shock transition reduces the downstream thermal energy of a CR modified shock compared to a gasdynamic shock with the same shock speed and Mach number. The Mach number of the gas subshock is also reduced. Even in a shock with a very large total Mach number, the subshock Mach number and compression ratio are typically $\\sim 3$ \\citep[\\eg][]{mal98,berzell99,kangetal02}.  In order to study these processes closely we recently developed a numerical scheme that self-consistently incorporates a thermal leakage CR injection model based on the analytic, nonlinear calculations of \\cite{mal98} \\citep{gies00} and implemented it into a combined Adaptive Mesh Refinement (AMR) gas dynamics and CR diffusion-convection code \\citep{kangetal02}. As described in the following section, there is only one rather well constrained parameter needed to express this injection model; namely, the strength of nonlinear MHD waves immediately downstream of the plasma subshock with respect to the upstream longitudinal magnetic field strength.  Previously, we applied our new code to a preliminary investigation of cosmic shocks emerging in large scale structure formation of the universe \\citep{kang02, kang03}. Utilizing the Bohm diffusion model in 1D quasi-parallel, plane-parallel shocks we showed that the resulting CR modified shocks evolve roughly ``self-similarly'' if the shock structure is expressed in terms of diffusion time and length scales, and that the time asymptotic states depend mainly on the shock Mach number, with only a weak dependence on the parameter needed in the injection model. We found for strong shocks that about $10^{-3}$ of incoming thermal particles are injected into the CR population with our thermal leakage injection model over the long term, and up to $\\sim 1/2$ of the shock kinetic energy flux measured in the initial shock frame is transferred to CRs. In the present contribution we confirm and extend our previous studies by including simulations of shock models with a broader range of physical parameters suitable for application to cosmic structure formation shocks, and also by including the influence of a pre-existing, upstream CR population, since many of the shocks of interest will probably form in previously shocked plasma. Our numerical simulations here are limited to quasi-parallel shocks in which the mean magnetic field lines are parallel to the direction of shock propagation. In quasi-planar shocks this is not a strong limitation so long as the obliquity of the magnetic field is small enough to keep its intersection subluminal and so long as the magnetic field is too weak to be dynamically important \\citep[\\eg][]{jok82,dru83}. In addition, since streaming motions inside cosmic sheets and filaments appear to stretch the field lines \\citep[\\eg][]{rkb98}, many structure formation ``internal'' shocks are likely to be quasi-parallel. The principle concern may be the influence of obliquity on the injection of low energy CRs at the shocks, as we discuss below.  In the following section we briefly outline our numerical methods in the CR/AMR hydrodynamics code. The simulation results are presented and discussed in \\S 3, followed by a summary in \\S 4. ",
        "conclusions": "Energy dissipation at collisionless shocks involves complex physical processes and plays crucial roles in the evolution of shock-heated astrophysical plasmas. Here our focus is CR acceleration via the first-order Fermi process at cosmic shocks that form in and around nonlinear structures during the formation of large scale structure in the universe. Using our cosmic-ray AMR Shock code \\citep{kangetal02}, we have calculated the CR acceleration at 1D quasi-parallel shocks driven by plane-parallel infall flows with $|u_0|=75-7500 {\\rm km s^{-1}}$ and preshock temperatures of $T_0=10^4-10^{7.6}$K. Mach numbers of the resulting shocks range over $ 2\\le M_s\\le 133$, which covers the characteristic shock Mach numbers found in cosmological hydrodynamic simulations of a $\\Lambda$CDM universe \\citep{ryu03}.  The main purpose of this study is to explore how the CR acceleration efficiency at cosmic structure shocks depends on the preshock temperature, shock speed and pre-existing CRs as well as shock Mach number. This, in turn, affects the thermal and dynamical history of the shocked-heated baryon gas, since the energy transfer to CRs reduces the thermal energy of the gas and increases the compressibility of the gas flow, much like radiative shocks.  Unlike pure gasdynamic shocks, downstream states of the CR modified shocks cannot be found from \"shock jump conditions\", so they have to be integrated time-dependently from given initial states or found from nonlinear analytical methods. Time-dependent simulations are computationally expensive, because the CR diffusion and acceleration processes contain a wide range of length and time scales. Fortunately, there are approximate similarity properties that can be employed to study the time asymptotic behavior of evolved CR modified shocks when they are mediated by Bohm-like diffusion. Firstly, the CR pressure approaches a steady-state value in a time scale comparable to the acceleration time scales for mildly relativistic protons after which the evolution of CR modified shocks becomes approximately ``self-similar''. This feature enables us to predict time asymptotic values of the CR acceleration efficiency, although we followed the CR acceleration up to only moderately relativistic energies (\\ie $p_{\\rm max}/(mc)\\sim 16$). We have confirmed these time asymptotic behaviors through much longer simulations with a new, more efficient, numerical scheme \\citep{twj04}, which were extended to achieve the cut-off momentum $p_{\\rm max}/(mc)\\sim 400$. Secondly, {\\it two models with the same Mach number, but with different accretion speeds} show qualitatively similar evolution and structure when the dynamical evolution is presented in terms of characteristic diffusion scales, $t_{\\rm n}=\\kappa_{\\rm n}/u_{\\rm n}^2$ and $x_{\\rm n}=\\kappa_{\\rm n}/u_{\\rm n}$. Such similarities are only approximate, however, because the partial CR pressure and the Bohm diffusion coefficient for transrelativistic CRs do depend on the momentum $m_pc$. Since the effective injection momentum is $p_{\\rm inj}/m_p c \\propto (u_s/c)$, the initial evolution depends on the shock speed as well as Mach number.  Finally, we state the main conclusions of our study, which can be summarized as follows:  1) Suprathermal particles can be injected efficiently into the CR population at quasi-parallel cosmic shocks via the thermal leakage process. For a given injection parameter defined in the text, $\\epsilon$, the fraction of injected CRs increases with the subshock Mach number, but approaches $\\xi \\sim 10^{-3}$ in the strong shock limit.  2) For a given value of $\\epsilon$, the acceleration efficiency increases with the shock Mach number, but approaches a similar value in the strong shock limit. Time asymptotic values of the ratio of CR energy to inflowing kinetic energy, $\\Phi$, converge to $\\Phi\\approx 0.5$ for $M_s \\gsim 30$ and it is relatively independent of other upstream or injection parameters (see figure 12). Thus, strong cosmic shocks can be mediated mostly by CRs and the gas thermal energy can be up to $\\sim 10$ times smaller than that expected for gasdynamic shocks.  3) For weak shocks, on the other hand, the acceleration efficiency increases with the injection rate (or $\\epsilon$) and the pre-exiting $P_c$ (see figure 12). For cosmologically common $M_s=3$ shocks, for example, $\\Phi \\sim 0.04$ for $\\epsilon=0.2$ and $\\Phi \\sim 0.2$ for $\\epsilon=0.3$ in the absence of pre-existing CRs. The presence of a pre-existing CR population acts effectively as a higher injection rate than the thermal leakage alone, leading to greatly enhanced CR acceleration efficiency in low Mach number shocks. We found even for weak shocks of $M_s=3$, that up to 40 \\% of the total energy flux through the shocks can be transferred to CRs, when the upstream CR pressure is comparable to the gas pressure in the preshock flow."
    },
    "0410/astro-ph0410512_arXiv.txt": {
        "abstract": "CCD photometry on the intermediate-band $uvbyCa$H$\\beta$ system is presented for the metal-deficient open cluster, NGC 2243. Restricting the data to probable single members of the cluster using the CMD and the photometric indices alone generates a sample of 100 stars at the cluster turnoff. The average $E(b-y)$ = 0.039 $\\pm$0.003 (s.e.m.) or $E(B-V)$ = 0.055 $\\pm$0.004 (s.e.m.), where the errors refer to internal errors alone. With this reddening, [Fe/H] is derived from  both $m_1$ and $hk$, using $b-y$ and H$\\beta$ as the temperature index. The agreement among the four approaches is excellent, leading to a final weighted average of [Fe/H] = --0.57 $\\pm$0.03 (s.e.m.) for the cluster, on a scale where the Hyades has [Fe/H] = +0.12. Using a combination of photometric and spectroscopic data, 27 probable cluster members are identified and used to delineate the red giant branch and a well-defined clump at $V$ = 13.70, while eliminating the so-called second clump at $V$ = 14.1. Interpolation between isochrones of appropriate [Fe/H] leads to an apparent modulus of $(m-M)$ = 13.15 $\\pm$0.1 and an age of 3.8 $\\pm$0.2 Gyr. A differential CMD comparison with Berkeley 29, a cluster with a galactocentric distance almost twice that of NGC 2243, constrains Berkeley 29 to be at least as young and as metal-rich as NGC 2243. ",
        "introduction": "This is the fifth paper in an extended series detailing the derivation of fundamental parameters in star clusters using precise intermediate-band photometry to identify probable cluster members and to calculate the cluster's reddening, metallicity, distance and age. The initial motivation for this study was provided by \\citet{tat97}, who used a homogeneous open cluster sample to identify structure within the galactic abundance gradient. The open clusters appear to populate a bimodal abundance distribution with the clusters interior to a galactocentric distance of $\\sim$10 kpc averaging [Fe/H] $\\sim$ 0.0 while those beyond this boundary have average [Fe/H] values of $\\sim -0.3$. The dispersion about each mean is less than 0.1 dex and implies that within each zone, no statistically significant abundance gradient exists, in contrast with the traditional assumption of a linear gradient over the entire disk. This structure has since been corroborated through the use of Cepheids by \\citet{an12,an02,lu03} and, most recently, with OB stars \\citep{da}. Further evidence for the lack of  a simple linear gradient across the entire galactic disk, particularly in the outer portion, is found in the spectroscopic data for Berkeley 29 and Saurer 1 \\citep{gc04}, a point we will return to in the last section of this paper. The exact origin and reason for the survival of the discontinuity remain unresolved issues, though the radial variation in the impact of spiral structure on star formation and stellar dynamics appears to be a likely contributor \\citep{sc01,mi02,lp03}.  Detailed justifications of the program and the observational approach adopted have been given in previous papers in the series \\citep{AT00a,AT00b,TW03,at04} (hereinafter referred to as Papers I through IV) and will not be repeated. Suffice it to say that the reality of the galactic features under discussion will remain questionable unless the error bars on the data are reduced to a level smaller than the size of the effect being evaluated or the size of the sample is statistically enhanced. The overall goal of this project is to do both.  An equally important aspect of this research is detailed testing of stellar evolution theory as exemplified by comparisons to stellar isochrones based upon models derived under a variety of assumptions. The agreement with (or the deviation from) the predicted distribution and location of stars within the color-magnitude diagram (CMD) has consistently provided a lever for adjusting our degree of confidence in the specifics of stellar interiors as a function of mass and age. A valuable illustration of this can be found in the discussion of the red giant branch distribution (first-ascent and clump giants) within NGC 3680, in comparison to NGC 752 and IC 4651, clusters of comparable age (Paper IV).  The focus of this paper is the older, metal-deficient open cluster, NGC 2243. The key role of this cluster in the sequential development of the project is as a representative of the moderately metal-poor class of open clusters found in the galactic anticenter, a class that contributes significantly to the definition of the galactic abundance gradient beyond the solar circle. In our context, it typifies the metal-poor population found beyond the discontinuity in the galactic abundance gradient at a galactocentric radius between 10 and 11 kpc, a population that, to date, {\\it has no cluster counterpart within the solar circle}, in sharp contrast with the field star sample \\citep{tat97}.  The early photometric attempts to understand the nature of NGC 2243 from CMD and color-color diagrams were based upon a mixture of photoelectric and photographic data \\citep{ha, vb}. While agreeing that the cluster was metal-deficient relative to the sun, the disagreement over the exact value of the metallicity, and most of the other key cluster parameters, was significant. There have since been three CCD-based studies, \\citet{be,bo,ka} (hereinafter referred to as BE, BO, and KA, respectively). The first two used $BV$ photometry, while the last used $VI$. We will defer discussion of their conclusions, as well as those of other investigators, regarding the cluster parameters to the sections of the paper where they are derived within the present analysis.  Section 2 contains the details of the $uvbyCa$H$\\beta$ CCD observations, their reduction and transformation to the standard system, and a search for photometrically anomalous stars. In Sec. 3 we discuss the CMD and begin the process of identifying the sample of probable cluster members. Sec. 4 contains the derivation of the fundamental cluster parameters of reddening and metallicity. In Sec. 5, these are combined with broad-band data to derive the distance and age through comparisons with theoretical isochrones and to allow a differential comparison with the most distant open cluster, Berkeley 29 (hereinafter Ber 29). Sec. 6 summarizes our conclusions regarding NGC 2243, Ber 29 and their role with respect to the abundance gradient of the galactic disk. ",
        "conclusions": "The specific goal of this investigation has been the derivation of the key cluster parameters of reddening, metallicity, distance, and age for the metal-poor anticenter cluster, NGC 2243.  We have attempted to use the various photometric indices to optimize the sample of single-star cluster members near the turnoff with high precision photometry in determining $E(B-V)$ = 0.055 $\\pm$ 0.004 and [Fe/H] = --0.57 $\\pm$ 0.03, with excellent agreement among the various combinations of indices used to derive [Fe/H]. This point is crucial because $m_1$, $hk$,  $b-y$, and H$\\beta$ are tied to the standard system independently, {\\it i.e.}, each has its own transformation from instrumental to standard system. If there were significant errors in the zero-points of the indices, the variation in the dependence of each color index on reddening and metallicity would conspire to generate much greater differences among the derived metallicities from the four techniques than are found in the analysis. In fact, the majority of the scatter among the derived [Fe/H] values can be explained solely by the internal scatter within the photometry for the various indices.   Given the reddening and metallicity, one should be able to obtain the distance and age by comparison to appropriate isochrones. Using color-color diagrams coupled with the modest amount of membership information supplied by high quality radial velocities, we were able to isolate highly probable members from the extended cluster area rather than just the core. This proved valuable in that it excluded a large number of field giants that confused the location of the subgiant branch and eliminated the controversial second clump from the discussion as a random projection of field stars. The note of caution that underlies this claim is that two of these stars have radial velocities consistent with membership, but are excluded because their abundances are inconsistent with those of the cluster. If some form of anomalous evolution has the ability to make normal cluster stars appear metal-rich because of mass transfer, mixing, or some unknown phenomenon, these stars should be revisited. However, for our purposes, their exclusion is reasonable.  Comparison of the cluster CMD to isochrones that bracket the derived [Fe/H] produces a set of parameters and discrepancies indicating that neither [Fe/H] is ideal, but that NGC 2243 lies at an [Fe/H] intermediate to that of the sets tested. The final, best fit parameters are $(m-M)$ = 13.15 $\\pm$0.10 at an age of 3.8 $\\pm$0.2 Gyr. As always, it cannot be overemphasized that these estimates are tied to a specific set of isochrones generated using a specific conversion between the theoretical and observational plane. In a relative sense, the data confirm past analyses using cluster morphology that NGC 2243 is approximately the same age as M67 but younger than Mel 66. The data should be on a the same scale as the clusters analyzed in previous papers in this series.   Given our results for NGC 2243, the cluster is located at a galactocentric distance interior to 11 kpc, but at the edge of the galactic discontinuity. A differential comparison between NGC 2243 and Ber 29 has been undertaken because the latter cluster now represents the most distant open cluster with a supposedly reliable [Fe/H] based upon spectroscopy. Unfortunately, the spectroscopic abundances differ by 0.3 dex, despite having overlapping one-sigma error bars, making the interpretation of the distant galactic gradient an exercise in personal bias. If one believes in a uniform gradient over a distance of 17 kpc, the high [Fe/H], when coupled with Saurer 1, forces the beholder to dismiss the clusters as anomalies unrepresentative of the real disk \\citep{gc04}. Adoption of the low metallicity is consistent with the claim of a linear galactic gradient \\citep{to}, but leaves no explanation for Saurer 1. An even simpler solution that is consistent with the higher metallicity for Ber 29 and Saurer 1 is that there is no significant gradient beyond 10.5 kpc, as found by \\citet{tat97}. The current reality is that if these two cluster have abundances typical of clusters near 11 kpc, it is intriguing, but hardly definitive evidence for any claim regarding the nature of the disk at galactocentric distances near 20 kpc. More and better data are needed both at large distance and in the key region bracketing 10 kpc.  Given the similar morphology of NGC 2243 and Ber 29, we can place some constraints on the differential properties of the two clusters due to the significant improvement in reddening estimation for Ber 29. The reddening maps of \\citet{sc98} coupled with the spectroscopic data of \\citet{gc04} imply that Ber 29 has $E(B-V)$ below 0.094. If we adopt this limit as the true value for the cluster and assume that  Ber 29 has [Fe/H] identical to NGC 2243, a differential comparison of the two clusters' CMDs implies $(m-M)$ = 16.15 for Ber 29 and an age slightly younger than NGC 2243. Any attempt to make Ber 29 significantly more metal-poor than NGC 2243 would require that Ber 29 be older than NGC 2243 and/or that the reddening estimate for Ber 29 be raised above the limit. The derived parameters for Ber 29 are in surprisingly good agreement with those of \\citet{to} and, when coupled to the success of the early analysis of NGC 2243 by \\citet{bo} in comparison with the current investigation, the convergence of the results is growing evidence of the power of the technique of CMD synthesis in generating reliable results in the absence of more direct observational data. Adjusting the age estimate for Ber 29 if [Fe/H] = --0.44 leads to 2.8 Gyr.  However, spectroscopic data \\citep{gc04} indicate that Ber 29 does not have scaled-solar abundances and, unlike NGC 2243, has some light-alpha-enhanced abundances. If we choose to compare the cluster to the available isochrones with the closest match to this effect, the age estimates increase by approximately 10\\% but, more importantly, the agreement between the location and luminosity of the giant branch and clump for [Fe/H] = --0.38 is excellent, in sharp contrast with the results for NGC 2243. The apparent modulus for Ber 29 rises slightly to $(m-M)$ = 16.25. On the basis of the internal consistency of the comparisons made above, it would be difficult to assign an [Fe/H] as low as --0.8 to Ber 29; all of our results indicate that the value of --0.5 $\\pm$0.1 is the most probable."
    },
    "0410/astro-ph0410038_arXiv.txt": {
        "abstract": "The cosmological utility of type Ia Supernovae (SNe) prompted numerous studies of these events, and they are now well characterized observationally, both as individual objects and as a population. In contrast, all other types of SNe (i.e., core-collapse events) are not as well observationally characterized. While some individual events have been studied in great detail (e.g., SN 1987A or SN 1998bw), the global properties of the core-collapse SN population are little known. However, in recent years, major drivers for change have emerged, among them the verification of the connection between core-collapse SNe and long-duration Gamma-Ray Bursts (GRBs), the possible utility of some core-collapse SNe (type II-P) as independent cosmological probes, and studies of core-collapse SNe as high redshift targets for missions like SNAP and JWST. The Caltech Core-Collapse Project (CCCP) is a large observational program using the Hale 200'' and the robotic 60\" telescopes at Palomar observatory to obtain optical photometry, spectroscopy and IR photometry of $\\sim50$ nearby core-collapse SNe. The program is designed to provide a complete sample of core-collapse events, with well-defined selection criteria and uniform, high-quality optical/IR observations, as well as radio and X-ray light curves for some events. We will use this sample to characterize the little-studied properties of core-collapse SNe as a population. The sample will be used as a comparison set for studies of SNe associated with Gamma-Ray Bursts, to promote and calibrate the use of SNe II-P for cosmography, and to set the stage for investigations of SNe at high-z using coming space missions such as SNAP and JWST. ",
        "introduction": "For many decades Supernovae (SNe) have been a focus of scientific interest, as sources of heavy elements and cosmic rays, a fundamental process shaping gas dynamics in galaxies, and, more recently, as cosmological distance indicators. SN explosions can be roughly divided into two main categories: type Ia events (SNe Ia) are believed to be thermonuclear explosions of white dwarf stars, while all other SNe result from core-collapse of massive progenitor stars. The great cosmological utility of SNe Ia (e.g., Riess et al. 1998; Perlmutter et al. 1999; Riess et al. 2004) spurred intensive studies of this SN sub-type. Both detailed studies of individual nearby events, and numerous works characterizing the properties of the SN Ia population, have been published. This is not the case, however, for core-collapse SNe. While some individual core-collapse events are very well studied (e.g., SN 1987A) the properties of the population of these events are not well known - e.g., the average peak magnitude and light curve shape, the distribution around this average, the typical spectroscopic properties and their evolution, etc.  Recently, several possible drivers for change have emerged. The connection between hydrogen-deficient core-collapse SNe (SNe Ib/c) and Gamma-Ray Bursts (GRBs; Stanek et al. 2003; Hjorth et al. 2003; Matheson et al. 2003) resulted in great interest in these events. The possible usefulness of hydrogen-rich core-collapse SNe II-P as distance estimators (e.g., Hamuy \\& Pinto 2002) may offer an independent method to test type Ia cosmology. The need to know the average properties of core-collapse SNe in order to interpret the results from surveys for SNe at high redshift has been demonstrated by several works (e.g., Dahlen \\& Fransson 1999; Sullivan et al. 2000; Gal-Yam, Maoz and Sharon 2002; Sharon 2004; Dahlen et al. 2004) and will become even more urgent when larger planned surveys with the SNAP and JWST space missions are carried out. ",
        "conclusions": ""
    },
    "0410/astro-ph0410506_arXiv.txt": {
        "abstract": "{We present VLT FORS1 and FORS2 spectra of 39 candidate high-redshift supernovae that were discovered as part of a cosmological study using Type Ia supernovae (SNe~Ia) over a wide range of redshifts. From the spectra alone, 20 candidates are spectrally classified as SNe~Ia with redshifts ranging from $z=0.212$ to $z=1.181$. Of the remaining 19 candidates, 1 might be a Type II supernova and 11 exhibit broad supernova-like spectral features and/or have supernova-like light curves. The candidates were discovered in 8 separate ground-based searches. In those searches in which SNe~Ia at $z \\sim 0.5$ were targeted, over 80\\% of the observed candidates were spectrally classified as SNe~Ia. In those searches in which SNe~Ia with $z > 1$ were targeted, 4 candidates with $z > 1$ were spectrally classified as SNe~Ia and later followed with ground and space based observatories. We present the spectra of all candidates, including those that could not be spectrally classified as supernova. ",
        "introduction": "Over the past decade, observations of SNe~Ia have played a leading role in measuring the expansion history of the Universe and in constraining cosmological parameters. It was through these observations that we discovered that the expansion is currently accelerating and that the Universe is presently dominated by an unknown form of dark energy with a negative equation of state (Perlmutter et al. \\cite{Perlmutter98}; Garnavich et al. \\cite{Garnavich98}; Schmidt et al. \\cite{Schmidt98}; Riess et al. \\cite{Riess98}; Perlmutter et al. \\cite{Perlmutter99}; Tonry et al. \\cite{Tonry03}; Knop et al. \\cite{Knop03}, Riess et al. \\cite{Riess04}; Barris et al. \\cite{Barris04};  for a review, see Perlmutter and Schmidt \\cite{Perlmutter03}).  When these results are combined with the results that have been derived from the fluctuations in the cosmic microwave background (Jaffe et al. \\cite{Jaffe01}; Bennett et al. \\cite{Bennett03}; Spergel et al. \\cite{Spergel03}), the properties of massive clusters (Allen, Schmidt \\& Fabian \\cite{Allen02}; Borgani et al. \\cite{Borgani01}) and the large scale structure of galaxies (Hawkins et al. \\cite{Hawkins02}), a picture of a flat Universe dominated by dark energy emerges.  Considerable effort has been directed towards extending the redshift range over which SNe~Ia are observed.  The Hubble diagram of SNe~Ia with $z \\sim 0.5$ is degenerate to a linear combination of $\\Omega_{\\rm M}$ and $\\Omega_{\\Lambda}$. Hence, an independent determination of these two parameters from SNe~Ia at $z \\sim 0.5$ is not possible.  However, observations of SNe~Ia over a wide range of redshifts and, in particular, very distant ($z \\ga 1$) SNe~Ia can break this degeneracy (Goobar and Perlmutter \\cite{Goobar95}). With this aim in mind, and following a highly successful pilot search (Aldering \\cite{Aldering}), the Supernova Cosmology Project (SCP) started a program to discover, spectrally confirm and photometrically monitor a substantial number of SNe~Ia with redshifts beyond one.  In this paper, we present VLT FORS1 and FORS2 spectra of 39 candidate high redshift supernovae. We present all spectra, including those spectra for which a secure spectroscopic classification could not be made.  The results of the photometric follow-up, the derived apparent magnitudes and the implications these measurements have for cosmology will be reported elsewhere. ",
        "conclusions": "In terms of classifying candidates from the spectra alone, there is a clear correlation between the number of candidates that are classified as SNe~Ia and the redshift at which SNe~Ia were targeted. In searches 1, 2, 3 and 5, (See Table. \\ref{tab:campaigns}), where SNe~Ia at $z\\sim0.5$ were targeted for VLT spectroscopic follow-up, 13 out of 16 candidates (excluding T02-049) are classified as SNe~Ia.  In search 8, where SNe~Ia with $z > 1$ were targeted for VLT spectroscopic follow-up, 4 out of 13 candidates are classified as SNe~Ia.  There are multiple reasons for the large difference. The aim of search 8 was to find several $z > 1$ SN Ia and the strategy of the spectroscopic follow-up was tuned to make the best use of the time that was available. In general, each candidate was first observed for one hour. Candidates that were found to have $z < 1$ were no longer observed. This included SuF02-002, SuF02-005, \\object{SN 2002km} and SuF02-028. In one case (\\object{SN 2002km}) a secure classification could be made. In the other three cases, a supernova might have been identified if we had chosen to integrate longer. Alternatively, if the candidate showed evidence for broad features or if the redshift from host galaxy lines (in particular [OII]) placed the host at $z >1$, the candidates were re-observed during later nights. As the amount of allocated time was limited, not all promising candidates could be followed. These factors led to a lower overall yield at $z < 1$, but they also enabled us to confirm several $z>1$ SNe~Ia and to obtain their redshifts.  Nevertheless, the spectroscopic confirmation of $z>1$ SNe Ia is challenging.  At $z \\sim 1$, SNe Ia are about 1.5 magnitudes fainter than at $z \\sim 0.5$.  Additionally, the spectral features that one uses for classification shift further and further into the red where sky subtraction can be difficult because of variable night-sky emission and detector fringing. This can be partially compensated by integrating longer and using instruments and telescopes that are efficient in the 600 to 1000 nm spectral region. Although the spectra of SNe~Ia show significant features short-ward of the broad CaII feature at 3900~\\AA\\ that could be used to aid the classification, the lack of good quality UV spectra for nearby supernovae of all types means that these features cannot be used without using features that are further into the red.  For $z>1$, which have peak magnitudes near $\\rm I\\sim25$, an additional source of ambiguity appears. Given the typical signal-to-noise that one can achieve with state-of-the-art instrumentation, one can sometimes match the spectra equally well with  SNe~Ia at two different redshifts. Fortunately, host galaxy lines, either [OII] or H and K or sometimes all three, can be used to measure  a precise redshift in most cases. In this paper, SNe 2002fe, 2002gi, 2002kn, 2002kr and 2002ks fall into this category. However, in other cases, such as SuF02-007 and \\object{SN 2002lc}, there are no clear galaxy lines, even though the spectra of these candidates show broad features.  In terms of classifying candidates from the spectra alone, there is also a correspondence between the selection criteria that are used to select candidates and the percentage of candidates that are spectrally identified as SNe~Ia.  In the rolling search with the CFHT, none of the 5 candidates could be spectrally confirmed as a SN~Ia. For comparison, in the Spring 2002 search with CTIO, all four candidates (excluding T02-049) were confirmed as SNe~Ia.  Although the numbers are small, they are significant. The search area and candidate selection criteria of the rolling search were such that the search was also sensitive to relatively fainter supernovae (Type II or SN~1991bg-like supernovae) on relatively brighter hosts. The spectra confirm this as many of the candidates from the rolling search are dominated by the light of the host galaxy.  The contribution from the host galaxy can be approximately quantified with the percentage increase in the flux of the candidate between the search and reference images. A small increase usually means a significant amount of host contamination. A very large or formally infinite increase usually means little or no host contamination. The flux is measured over a fixed aperture whose diameter depends on the seeing. The signal-to-noise ratio of the detection over the same fixed aperture provides a measure of the significance of the detection. A low signal-to-noise ratio usually means that the candidate is faint, and this could mean that the candidate, if it is a SN~Ia, is either very distant or has been caught very early. In Fig.  \\ref{fig:increase} the signal-to-noise ratio of the detection is plotted against the percentage increase for candidates that were brighter  than ${\\rm I}=24.7$ at the time of discovery.\\footnote{For the surveys done in R-band, we use ${\\rm R}=24.7$} Candidates from the CFHT 2002 search are highlighted with large circles.  The figure shows that classification from spectroscopy is generally not successful if the percentage increase is below $\\sim 25\\%$.  The boundary of this region is marked with a dashed line in Fig.  \\ref{fig:increase}. Candidates in which the percentage increase is less than 25\\% are thus noted in Table \\ref{tab:results}.  \\begin{figure} \\centering \\includegraphics[width=8cm]{1504fi4.eps} \\caption{The percentage increase versus the  signal-to-noise ratio for candidates that were brighter than ${\\rm I}=24.7$ at the time of discovery. Candidates that were classified as Ia (Ia$^*$) are plotted with solid circles (squares), C02-031 - a possible SN~II - is plotted as a star and unclassified candidates are plotted with open circles. Candidates from the CFHT 2002 search are highlighted with large circles.  Candidates with arrows have percentage increases that are greater than 1000\\%, which means that the host was considerably fainter than the candidate and perhaps undetected.  The dashed line marks the region where the percentage increase is 25\\% or less. Candidates in this region are difficult to classify spectrally.} \\label{fig:increase} \\end{figure}  Not one of the candidates has broad emission lines that would indicate that it is an AGN. This demonstrates that our selection strategy, which selects against candidates having small intensity variations that are also precisely centered on the host galaxy, is quite effective in rejecting AGNs.  In this paper, we have strictly used only the spectra for classification purposes and all the classifications listed in Table \\ref{tab:results} are based on the spectra alone.  However, in searches like the CFHT Spring 2002 search and the Subaru Fall 2002 search, where spectroscopic confirmation is difficult because the candidates are near relatively bright hosts or because the candidates are relatively faint, additional information such as the light curve or the colour of the candidate can become part of the criteria used for classification. The strategy of these two searches was such that most of the candidates were also monitored during the subsequent weeks and months. Of the 18 candidates that were observed in these two surveys, 5 were spectrally classified as either Ia, Ia$^*$ or II?. Of the remaining 13 candidates,  10 were followed with sufficient coverage (more than 4 light curve points) and 9 have supernova-like light curves. This includes \\object{SN 2002fr}, \\object{SN 2002fm}, \\object{SN 2002kq}, SuF02-007, \\object{SN 2002lr},  SuF02-026, \\object{SN 2002kz}, SuF02-051 and \\object{SN 2002kx}. These cases are noted in Table \\ref{tab:results} and in the comments on individual candidates."
    },
    "0410/astro-ph0410730_arXiv.txt": {
        "abstract": "We discuss the internal structure of stars in the mass range 1.5 to 4\\,$M_\\odot$ from the PMS to the subgiant phase with a particular emphasis on the convective core and the convective superficial layers. Different physical aspects are considered such as overshooting, treatment of convection, microscopic diffusion and rotation. Their influence on the internal structure and on the photospheric chemical abundances is briefly described.  The role of binarity in determining the observed properties and as a tool to constrain the internal structure is also introduced and the current limits of theories of orbital evolution and of available binary data--sets are discussed. ",
        "introduction": "The theoretical evolution of A stars is extremely simple to compute if one ignores complex phenomena such as gravitational settling, radiative forces, rotation, turbulent mixing, magnetic fields, binarity... We shall briefly introduce the infuence of different treatments of convection, diffusion and rotation after a presentation of a \"conservative\" situation in which the only mixings come from convection and overshooting. Binarity as a tool to constrain stellar models is then discussed. ",
        "conclusions": ""
    },
    "0410/astro-ph0410440_arXiv.txt": {
        "abstract": "{We derive basic stellar parameters (angular diameters, effective temperatures, metallicities) and interstellar reddening for all $\\beta$ Cephei stars observed during the IUE satellite mission, including those belonging to three open clusters. The parameters are derived by means of an algorithmic procedure of fitting theoretical flux distributions to the low-resolution IUE spectra and ground-based spectrophotometric observations. Since the metallicity has a special importance for pulsating B type stars, we focus our attention in particular on this parameter. ",
        "introduction": "$\\beta$ Cephei variables are a well-known group of early B-type pulsating stars. Their pulsations are driven by the classical $\\kappa$-mechanism, operating in the layer of metal opacity bump at $T\\approx 2\\cdot 10^5~{\\rm K}$ caused by a huge number of the absorption lines of the iron-group elements. The explanation of oscillations in these stars was possible only after the publication of the new opacity tables by Iglesias et al. (\\cite{iglesias92}), and was given by Cox et al. (\\cite{cox}), Moskalik \\& Dziembowski (\\cite{moskalik}), Kiriakidis et al. (\\cite{kiriakidis}) and Dziembowski \\& Pamyatnykh (\\cite{dziembowski}). As was shown by Pamyatnykh (\\cite{pamyatnykh99}), the extent of the theoretical instability domain in the H-R diagram is very sensitive to the heavy element abundance, $Z$. The instability domain decreases with decreasing value of $Z$, and for $Z\\approx 0.01$ there are no unstable modes in the observed $\\beta$ Cephei domain. Therefore, the determination of metallicities is the first step of testing the theory of pulsational instability in massive Main Sequence stars. It would be important to fix the minimum observed value of [m/H] for which the oscillations are still present in these stars. Most $\\beta$ Cep pulsators belong to our Galaxy, and they are supposed to have the metal abundance close to the solar value. Recently, Pigulski \\& Ko{\\l}aczkowski (\\cite{pigulskikol02}) discovered the first extragalactic $\\beta$ Cep stars in LMC, where the average metallicity is much below the solar value. All LMC $\\beta$~Cep stars are in or close to young associations and they have probably higher metal abundance than average.  Moreover, the question \"which pulsational parameters are correlated with the metal abundance ?\" is still open.  In this paper we analyze the IUE ({\\it International Ultraviolet Explorer}) data combined with ground-based spectrophotometric observations of $\\beta$ Cephei stars. The ultraviolet (UV) part of the spectra of main-sequence B-type stars is very rich in lines of the iron-group elements. Because the greatest amount of energy for these objects is emitted in the spectral region between 1216 and 10000~\\AA, even the low-resolution spectra can give credible values of stellar parameters such as the effective temperature, $T_{\\rm eff}$, surface gravity, $\\log g$, microturbulent velocity, $v_t$, metallicity, [m/H], angular diameter, $\\theta$, and the color excess, $E(B-V)$. The usefulness of the low-resolution IUE spectra for this purpose has already been demonstrated by many authors. Fitzpatrick \\& Massa (\\cite{fitzpatrick99}) obtained the mean stellar parameters for fourteen main sequence stars in the spectral range O9.5 -- A1. Morossi et al. (\\cite{morossi}) used IUE spectra, combined with the visual spectrophotometric observations, to derive [m/H], $T_{\\rm eff}$ and $\\theta$ for 53 objects in the solar neighborhood. In Niemczura et al.(\\cite{niemczura02}) we presented the results for some objects from the paper of Code et al. (\\cite{code}) and for the Magellanic Clouds B type stars, obtained by the least-squares optimization algorithm. In a previous paper (Daszy\\'nska et al. \\cite{daszynska02}), using the above mentioned method, we determined the values of the metal abundance parameter in the atmospheres of 31 field $\\beta$ Cephei stars and 16 stars belonging to three open clusters NGC~3293, NGC~4755 and NGC~6231. Here we use the same data as in Daszy\\'nska et al. (\\cite{daszynska02}) but we derive the stellar parameters with the improved procedure which allows, in particular, for a more objective error estimation. This method was applied recently by Niemczura (\\cite{niemczura03}) to determine the metallicities, effective temperatures, stellar diameters and interstellar reddenings for Slowly Pulsating B-type stars.  In Sect.~2, all stars we analyze are presented and spectrophotometric observations together with their calibration are shown. The procedure of fitting the theoretical spectra to the observations and the method of bootstrap resampling adopted for determining the errors of parameters are briefly described in Sect.~3. The results and interparametric correlations are presented in Sect.~4. In Sect.~5, we discuss in detail the values of metallicities we obtained. Conclusions are given in Sect.~6. ",
        "conclusions": "We derived basic stellar parameters such as the effective temperature, gravity, angular diameter, metallicity and interstellar reddening for all $\\beta$ Cep stars monitored by the International Ultraviolet Explorer satellite. Among them there were field $\\beta$ Cep stars and variables belonging to three young open clusters: NGC 3293, NGC 4755 and NGC 6231. The mean values of the [m/H] parameter for the $\\beta$~Cephei stars are equal to $-0.13\\pm0.03$~dex and $-0.14\\pm0.03$~dex for IUE/NEWSIPS and IUE/INES data, respectively. The mean metallicity for the field $\\beta$~Cephei variables is equal to $-0.14\\pm0.03$~dex (IUE/NEWSIPS) and  $-0.13\\pm0.03$~dex (IUE/INES), while the mean values of [m/H] for the cluster stars obtained from the analysis of IUE/NEWSIPS spectra are equal to $+0.05\\pm0.06$~dex for NGC~3293, $-0.43\\pm0.05$~dex for NGC~4755 and $-0.01\\pm0.06$~dex for NGC~6231. From the IUE/INES spectra, we obtained nearly the same values: $+0.08\\pm0.04$~dex for NGC~3293, $-0.40\\pm0.06$~dex for NGC~4755 and $-0.05\\pm0.05$~dex for NGC~6231. These values are consistent with our previous results (Daszy\\'nska et al. \\cite{daszynska02}), obtained with the less accurate method.  The values of the metallicity parameter we obtained, are not correlated with any pulsational parameter or the projected rotational velocity. However, we have noticed that, on the whole, multiperiodic variables have higher values of metallicity than the monoperiodic ones. This is consistent with pulsation theory.  The abundance of metals, especially that of iron, in $\\beta$~Cephei variables is one of the fundamental parameters. In the case of main-sequence B stars, metallicities, [m/H], obtained from the UV spectra give mainly an estimation of [Fe/H].  The value of [Fe/H] in the Kurucz's models is larger than in the widely used in the pulsation theory OPAL and OP data. One can therefore expect that [Fe/H] determined using the ATLAS9 models should be lower in comparison with [Fe/H] obtained on the basis of the OPAL or OP data (see also Fitzpatrick \\& Massa \\cite{fitzpatrick99}). To compare observational metallicities of the stars with the values resulting from the theory, we have to include the correction for [m/H] of about 0.12~dex. The metallicities of the $\\beta$~Cep stars predicted by the theory of pulsations are not in contradiction with most values determined by us. Pamyatnykh (\\cite{pamyatnykh99}) showed that the $\\beta$~Cephei instability strip vanishes for $Z \\approx 0.01$, corresponding to [m/H] $\\approx -0.30$ dex, and to [m/H] $\\approx -0.42$~dex from the Kurucz's fluxes. There are several stars with metallicity parameter lower than this limit. But we have to remember that our values give information about photospheric metal abundances. Some other effects which we did not take into account such as diffusion or element mixing may also be important. Besides, we used a standard mixture of elements. Nevertheless, our results provide important information about the metallicity range for these pulsating stars. In particular, because there is a problem with the estimation of the [m/H] parameter for early type stars from photometry, making use of low-resolution ultraviolet spectra seems up to now the best and fastest way to determine the metallicity of B-type pulsating variables.  However, for asteroseismological purposes, there is a need for a detailed analysis of chemical composition, because oscillation frequencies are very sensitive to the adopted mixture, and this is the aim of our future work.\\\\   \\noindent {\\it Acknowledgements} We thank prof. Dziembowski and prof A. Pamyatnykh for permission of using the evolution code and non-adiabatic pulsation code."
    },
    "0410/astro-ph0410489_arXiv.txt": {
        "abstract": "We have obtained Johnson V and Gunn i photometry for a large number of Local Group galaxies using the Isaac Newton Telescope Wide Field Camera (INT~WFC).  The majority of these galaxies are members of the M31 subgroup and the observations are deep enough to study the top few magnitudes of the red giant branch in each system.  We previously measured the location of the tip of the red giant branch (TRGB) for Andromeda~I, Andromeda~II and M33 to within systematic uncertainties of typically $< 0.05$~mags (\\citealt{mcconnachie2004a}).  As the TRGB acts as a standard candle in old, metal poor stellar populations, we were able to derive distances to each of these galaxies.  Here we derive TRGB distances to the giant spiral galaxy M31 and 13 additional dwarf galaxies - NGC~205, NGC~185, NGC~147, Pegasus, WLM, LGS3, Cetus, Aquarius, And~III, And~V, And~VI, And~VII and the newly discovered dwarf spheroidal And~IX.  The observations for each of the dwarf galaxies were intentionally taken in photometric conditions.  In addition to the distances, we also self-consistently derive the median metallicity of each system from the colour of their red giant branches. This allows us to take into account the small metallicity variation of the absolute I magnitude of the TRGB.  The homogeneous nature of our data and the identical analysis applied to each of the 17 Local Group galaxies ensures that these estimates form a reliable set of distance and metallicity determinations that are ideal for comparative studies of Local Group galaxy properties. ",
        "introduction": "The Local Group is home to some 35 galaxies, the majority of which are suspected to be satellites of  the two most massive bodies, the Galaxy and         M31        (for         recent         reviews,        see \\citealt{mateo1998a,lyndenbell1999,vandenbergh1999}).    Many  of  the dwarfs, such as M32, have  been known about for decades, while others, such   as   And   V,   VI,   VII  \\&   IX   are   recent   discoveries (\\citealt{armandroff1998,armandroff1999,karachentsev1999,zucker2004}). As the  discovery of  And  IX by  \\cite{zucker2004}  demonstrates, it  is unlikely that  even now we have  a full inventory of  the Local Group, and no  doubt some  faint bodies still  await discovery.   Interest in Local  Group galaxies  stems from  their proximity  and representative nature, since they are relatively easily resolved and include examples over a wide range of luminosities.  With so many galaxies known in  the Local Group it has become possible to conduct comparative studies  of the different population types such as  the dwarf  irregulars  and dwarf  spheroidals.  However,  accurate studies  of stellar  populations  requires accurate  knowledge of  the distance to each galaxy, and this becomes significantly more important when  populations  in different  galaxies  are  to  be compared.   The distances to the nearby satellites of the Milky Way are believed to be known to  reasonable accuracy and are  based on a  variety of methods. However, this is not always the  case for the more distant Local Group galaxies, such as the satellites of M31.  Furthermore, the majority of distance  determinations  to these  galaxies  have  been conducted  as independent   studies  by   different  groups.   It   is  nevertheless advantageous to acquire a  homogeneous set of distance determinations. This  then removes  systematic  uncertainties that  may exist  between measurements  that  could  be  due to  the  data  aquisition/reduction process,  the standard  candle  employed, and  the  algorithm that  is applied.   Reliable relative  distances to  the galaxies  can  then be obtained,  and  provides  much  of  the  motivation  behind  the  work presented in this paper.  All of  the Local Group  galaxies contain a significant  population of old,  relatively metal-poor  stars (Population  II). For  M31  and the Galaxy, these stars are found predominantly in the stellar halo, while for the dwarfs, Population~II stars tend to be the dominant component. A lack of Cepheid variables in  this population means that in order to estimate  distances to  them other  indicators have  to be  used.  One obvious  candidate  is  based on  the  Tip  of  the Red  Giant  Branch (TRGB). This point in stellar evolution marks the onset of core helium burning in  Red Giant  Branch (RGB) stars  and observation  and theory show  it to  occur  at  a relatively  constant  I-band magnitude  (eg. \\citealt{dacosta1990,salaris1997}).   \\cite{barker2004}  have recently conducted  a  detailed study  of  the reliability  of  the  TRGB as  a distance  indicator for  stellar populations  with a  variety  of star formation histories and find that  it is reliable for any system which is not  dominated by  a metal-rich population  ([Fe/H]$ > -0.3$)  or a population   with   a    substantial   young   ($\\lesssim   1.7$~Gyrs) component. This makes it ideal  for calculating distances in the Local Group.  \\cite{lee1993a}   were  influential   in  developing   a  quantitative technique, based upon a  Sobel edge detection algorithm, for detecting the  location of  the TRGB  in a  photometric dataset.   It  was later adapted   and  refined   by  \\cite{sakai1996}   and,   more  recently, \\cite{mendez2002} who also developed  a maximum likelihood approach to estimate  the TRGB  location.  Edge  detection  algorithms essentially assume  that the  luminosity function  in the  region of  the  tip are step-like,  and  their  simplest  application looks  for  the  largest absolute  change  in star  counts  between  neighbouring  bins in  the luminosity  function.  As  such, they  can be  easily affected  by the noise features which are often prevalent in real luminosity functions on which TRGB analyses are to be conducted.  In a  previous paper (\\citealt{mcconnachie2004a},  hereafter Paper~I), we developed a  new technique to allow for  the accurate determination of the TRGB as represented  in a photometric dataset. We advocated the use  of Wide  Field Cameras  for such  a study  as this  maximises the number of stars that contribute to the luminosity function, decreasing the effects  of poisson noise on  the measurement and  ensuring that a fair sample of the brightest RGB stars are present. In Paper~I we also introduced,  tested   and  implemented   a  new  algorithm   for  TRGB determination which uses a  data-adaptive slope to fit the `luminosity probability distribution'  (LPD) in  the region of  the tip  and hence derive a distance modulus for the system. The LPD is given by  \\begin{equation} \\label{smooth} \\wp\\left(m\\right)=\\sum_{i=1}^{N_\\star} \\frac{1}{\\sqrt{2\\pi}\\sigma_i}\\exp\\left(-\\frac{\\left(m_i-m\\right)^2}{2\\sigma_i}\\right), \\end{equation}  \\noindent where $m_i$ and $\\sigma_i$ are the magnitude and photometric error of the $i^{th}$ star  from a sample of $N_\\star$. This technique was shown  to have systematic uncertainties usually  of order \\mbox{$< 0.05$~mags.}  We went  on to calculate the distances,  as derived from data  taken  with  the   Isaac  Newton  Telescope  Wide  Field  Camera (INT~WFC), to  the dwarf galaxies And~I  \\& II, and  the spiral galaxy M33.  The  reader is referred  to this paper, and  references therein, for a discussion of the practicalities involved with this technique.  In this paper, we significantly extend the number of Local Group galaxies for which we have measured locations of the TRGB and distances, based on a large homogeneous photometric database of Local Group galaxies. This data has been taken with the INT WFC, a four-chip EEV 4K x 2K CCD mosaic camera with a $\\sim 0.29 \\square^\\circ$ field of view (\\citealt{walton2001}).  Over the past three years we have been using this instrument in conjunction with the Canada-France-Hawaii 12K Camera to conduct a photometric survey of M31 (\\citealt{ibata2001a,ferguson2002,mcconnachie2003}, Irwin et al., {\\it in preparation}), which is now complete over an area of 40 $\\square^\\circ$.  We have also presented photometry for M33, And~I, And~II (Paper I; see also Ferguson et al., {\\it in preparation}), and here we present colour magnitude diagrams (CMDs), metallicity information and distance determinations to And~III, And~V, And~VI, And~VII, And~IX, LGS3, Cetus, WLM, Pegasus, Aquarius, NGC~147, NGC~185, NGC~205 and M31. These galaxies constitute the majority of Local Group members visible from the northern hemisphere.  M32 is not included in this list due to severe crowding problems and its awkward projected location close to the centre of M31. Our data for IC10 is likewise not presented here as this object's close proximity to the Galactic plane causes a serious amount of differential reddening which complicates its analysis.  For the rest, the resulting set of homogeneous distance estimates is the largest that has been measured for the Local Group: the data for each galaxy were taken using the same telescope and the same instrument, through the same filters for similar exposure times.  The data for the dwarf galaxies was taken in conjunction with the main M31  and  M33  surveys.   Observations  for  each  dwarf  galaxy  were restricted to  photometric conditions simplfying  cross-calibration of all the data to the same photometric system.  All data taken were then reduced and calibrated in exactly  the same way with the same pipeline processing  (\\citealt{irwin2001}).  The  subsequent analysis  for each galaxy was also identical, with  the same algorithm applied to each to derive  the  distance  estimates.   As such,  differential  systematic errors  are  minimised  and  the  entire dataset  is  thus  ideal  for comparative studies of these galaxies.  Section 2 provides an overview of our photometry, discusses the reddening estimates, the assumed absolute I magnitude of the TRGB, the metallicity determination and how we deal with foreground and background contamination.  In Section~3 we briefly review the current knowledge of the stellar content of each of the galaxies and calculate the median metallicity and distance to each from the INT WFC data.  We finish in Section~4 with a discussion of our results and a comparison to previous work. ",
        "conclusions": "Table~1 lists the  galaxies analysed as part of  this study, including those from Paper~I, along with their positions, the adopted values for the reddening, the median metallicities as determined by the colour of the red giant  branch and the adopted value  of $M_I^{TRGB}$ using the modified  calibration  of  \\cite{bellazzini2004}.   Table~2  list  the measured position of  the TRGB, the distance modulus  and distance for each galaxy.   The three galaxies originally analysed  in Paper~I have had their distances recalcuated  to correct for the slight metallicity dependancy that was not previously addressed in Paper~I. Additionally, the TRGB location  in M33 has been rederived  using a local foreground correction  (in  the same  way  as for  M31)  and  using only  objects that lie within $2\\,-\\,\\sigma$ of the stellar locus. Also  included for comparision  are a selection  of previous  distance estimates  to each system from the literature.  Figure~\\ref{comparison} shows a graphical comparison  between our  results and  these earlier  measurements.  In general, good agreement is  observed and no obvious systematic offsets or trends  are visible.  For  those cases where Cepheid  distances are available       (M31:      \\citealt{joshi2003,freedman1990};      M33: \\citealt{lee2002}) our results match  these estimates to within better than a few tens of kiloparsecs.  For a  few of the fainter objects  in our sample (eg.  And~IX) we have relatively few  bright RGB stars available  to us, due to  the lack of such  stars in  these  systems. The  distance measurements  implicitly assume that the brightest RGB  stars in the system are good indicators of the  actual position of the TRGB.  For the case of  And~IX, the RGB luminosity  function  shows  a  steep  rise and  is  clearly  defined, compensating  in  part for  the  fewer  stars  available to  us.  Thus relatively  accurate  TRGB  measurements  can be  obtained  for  these systems,  even although the  number of  bright RGB  stars is  far less than, say, NGC~205 or M31.  We conclude with a list of  some of the main results evident from this study:  1. The distance to the  newly discovered dwarf spheroidal companion to M31, Andromeda~IX, is  measured to be $765 \\pm  25$~kpc. This is in good  agreement to  the distance  calculated  by \\cite{zucker2004}. Assuming   And~IX   does  not   have   a  strong   $\\alpha$-element enhancement, its metallicity is measured to be \\mbox{[Fe/H] $\\simeq -1.5$}. Figure~4 of \\cite{caldwell1999} shows an empirical relation that  is   found  to   exist  between  integrated   luminosity  and metallicity  for  dwarf  spheroidal  galaxies,  such  that  fainter galaxies   have  lower   metallicity.    \\cite{zucker2004}  measure $M_{tot,V} \\simeq -8.3$\\,mags which  would imply that And~IX should have  a  metallicity of  approximately  \\mbox{[Fe/H] $\\sim  -2.2$.} And~IX  thus appears  to  be more  metal-rich  than this  empirical relation would suggest.  2. And~V is measured to have \\mbox{[Fe/H] = -1.6}, again assuming no $\\alpha$-enhancement.  This disagrees with the recent measurement by \\cite{davidge2002} of \\mbox{[Fe/H] = -2.2}.  It does agree, however, with the earlier measurement by \\cite{armandroff1998}. Both this latter measurement and our measurement indicate that, like And~IX, And~V does not lie on the expected relation between integrated luminosity and metallicity. These results suggest that And~I, II, III, V, VI, VII \\& IX all have comparable metallicities in the range \\mbox{$-1.7 \\le$ [Fe/H] $\\le 1.4$}, with an estimated accuracy on the measurements of $\\pm\\,\\sim 0.1\\,-\\,0.2$\\,dex.  3. Aquarius is confirmed to lie $\\sim 1$~Mpc from the Milky Way. Early estimates  of   its  distance  had  placed  it   well  outside  the zero-velocity  surface  of  the  Local Group  until  \\cite{lee1999} showed that it was substantially  closer than this. We note however that the distance to Aquarius  is the least certain from this study for the reasons discussed in Section~3.12. It could potentially lie closer to us than we have derived here.  4. Pegasus is measured  to lie at $919 \\pm  30$~kpc.  This measurement agrees with earlier work done by \\cite{aparicio1994}. A more recent study  by  \\cite{gallagher1998}  places  this object  some  160~kpc closer -  this distance  is required  by their study  so as  to fit self-consistent  stellar population  models based  upon  the Geneva stellar evolutionary  tracks. \\cite{gallagher1998} also  require to adopt  a  larger reddenning  value  than  is  normally assumed  for Pegasus  (${\\rm E(B  -  V)} \\simeq  0.14$)  in order  to match  the colours  of  the  main  sequence.   Even by  adopting  this  larger reddening value, the TRGB  distance appears incompatible with their measurement of the distance.  5. NGC~185  and NGC~147 have  previously been  suspected of  forming a binary  system  due  to   their  small  angular  seperation  ($\\sim 1^\\circ$)     and    similar     line     of    sight     distances (\\citealt{vandenbergh1998}). The distances that we derive for these objects suggest  that they are likely to  be gravitationally bound, separated in  line of sight  by only $\\sim 60$\\,kpc.   $1^\\circ$ at the distance of  this system corresponds to $\\sim  11$\\,kpc, and so their physical  seperation is  also of order  60 kpc,  although the uncertainties  in each of  their distances  means that  the precise value may be substantially different to this.  The 17 Local Group galaxies listed in Tables~1~\\&~2 were observed with the same telescope, instrument and filters for similar exposure times and the resulting data were reduced using the same pipeline processing.  An identical analysis was then conducted on each to calculate the position of the TRGB, its median metallicity and its distance.  The use of a Wide Field Camera has maximised the number of stars observed per pointing.  This reduces the effect of Poisson noise on our measurements and ensures that the luminosity function that we derive is as accurate a representation of the intrinsic luminosity function of the galaxy as is possible in the region of the tip. The resulting set of metallicity and distance estimates have thus had systematic uncertainties minimised and are therefore ideal for any study of Local Group galaxy properties."
    },
    "0410/astro-ph0410676_arXiv.txt": {
        "abstract": "{ We have imaged the \\hx (8.3 GHz), \\hu (15 GHz), and H166$\\alpha$ (1.4 GHz) Radio Recombination Lines (RRLs) from NGC 253 at resolutions of 4.5 pc (0.4\\asecc) , 2.5 pc (0.2\\asecc) and 53 pc (4.5\\asecc) respectively. The \\hx line arises from individual compact sources, most of which possess radio continuum counterparts. The line widths range from $\\sim$200 km \\s for the sources near the radio nucleus to 70-100 km \\s for the extranuclear ones. These lines are emitted by gas at a density $\\sim$10$^4$ \\ccc. The remainder of the cm-wave RRLs arise in lower density gas ($\\sim$500 \\ccc) with a higher area filling factor and with ten times higher mass. A third component of higher density gas ($>$10$^4$ \\ccc) is required to explain the mm-wave RRLs. ",
        "introduction": "NGC 253 is a nearby (D=2.5 Mpc; 1\\asecc=12 pc) highly inclined spiral galaxy, undergoing a vigorous starburst (L$_{\\rm FIR}$=3$\\times$10$^{10}$ \\lsunn; \\citealt{th80}) in the central $\\sim$100 pc region. The ongoing star formation is concentrated along an inclined ring. A number of compact near and mid-IR (NIR and MIR) sources have been discovered in this region (Forbes, Ward, and Depoy 1991\\nocite{forbes91}; Forbes et al. 1993\\nocite{forbes93}; Pi\\~na et al. 1992\\nocite{pina92}; Keto et al. 1993\\nocite{keto93}; Sams et al. 1994\\nocite{sams94}). Optical imaging by Watson et al. (1996)\\nocite{watson96} revealed the presence of four star clusters, which were identified with individual MIR/NIR sources (see also Forbes et al. 2000\\nocite{forbes00}). The morphology of the radio continuum emission, which is unaffected by dust extinction, is quite different. Turner and Ho (1985)\\nocite{th85} discovered a linear chain of compact sources at 15 GHz, which were studied in detail by Antonucci and Ulvestad (1988)\\nocite{au88} and Ulvestad and Antonucci (1997; hereafter UA97)\\nocite{ua97}. The latter were able to classify the brighter compact sources as synchrotron or free-free dominated, based on the derived spectral indices. Kalas and Wynn-Williams (1994)\\nocite{kww94} and Sams et al (1994)\\nocite{sams94} showed that most of the infrared sources are regions of low dust extinction and are not coincident with the radio sources. Hence, the radio sources and the infrared-optical sources seem to trace very different types of objects. In particular, the peak of the NIR emission is offset by about 3.5\\asec southwest from the peak of the radio continuum emission (the radio nucleus). There is no associated radio emission towards the NIR peak and only weak NIR emission is observed near the radio nucleus. High resolution radio continuum observations with the Very Large Array (VLA) suggest that the radio nucleus hosts an Active Galactic Nucleus (AGN; \\citealt{th85}, UA97\\nocite{ua97}), which was subsequently confirmed (\\citealt{mohan02}, \\citealt{weaver02}). The NIR continuum peak, identified as a region of active star formation, also coincides with the position of the emission peaks of various ionised gas tracers ([NeII]: B\\\"oker, Krabbe and Storey 1998\\nocite{boker98}; Keto et al. 1999\\nocite{keto99}; \\brg: Forbes et al. 1993\\nocite{forbes93}; optical continuum and line emission: Engelbracht et al. 1998\\nocite{engel98}; Watson et al. 1996\\nocite{watson96}; Forbes et al. 2000\\nocite{forbes00}).   The immediate product of the starburst process, with its associated high star formation rate and efficiency, is a young massive stellar population. This population can studied through its surrounding ionised gas, the properties of which are determined partly by the ionising stars themselves, and partly by gas dynamics. Recent star formation is well traced by NIR emission, except in regions of high dust extinction. However, such highly obscured star formation can be probed by radio continuum and hydrogen radio recombination line (RRL) emission (except in cases where the continuum optical depth due to free-free emission is much larger than unity). RRL emission is insensitive to dust extinction, arises only in thermal gas, and can be imaged at arcsecond resolution. NGC 253 was the first extragalactic object (excluding the Magellanic Clouds) to have been detected in RRLs \\citep{sb77}. This line emission was initially modelled as being stimulated by the background nuclear emission \\citep{shaver78,mebold80}. The first interferometric observations of RRLs from NGC 253 were carried out by \\citet{ag90} at higher frequencies (the H166$\\alpha$, H110$\\alpha$ and the H92$\\alpha$ lines at 1.4 GHz, 4.8 GHz and 8.3 GHz respectively) using the VLA. Models indicated that these lines arise due to internal emission from high density gas (\\nee$\\sim$10$^4$ \\ccc; similar to compact HII regions). \\citet{puxley97} observed mm-wave RRLs from this galaxy and suggested that the origin of the RRL emission at all frequencies is spontaneous emission from optically thin HII regions. A total ionisation rate of (3.7 $\\pm$ 0.8)$\\times$10$^{53}$~\\s was derived, an order of magnitude higher than that derived from infrared emission lines. The ionisation budget and hence the star formation rate as determined by RRLs depends on the nature of the line emission, which is not yet fully understood. High resolution observations ($\\sim$1.3\\asec or 16 pc) of the 8.3 GHz RRL from NGC 253 were carried out by \\citet{ag96}. They showed that the line emission peaks at the radio nucleus (as do the higher frequency RRLs as well; \\citealt{zhao01}). They detected much weaker line emission near the NIR peak, which is unexpected, since the infrared peak is probably the most active star forming site in this galaxy. Using the high resolution \\hx and \\hu data presented in this work, the RRL emission from the radio nucleus alone was studied separately and the results were presented by \\citealt{mohan02}. Based on the results of the line emission models, it was concluded that the ionising source for the RRL emitting gas is probably a central low luminosity AGN with weak X-ray emission, which was subsequently confirmed through $CHANDRA$ observations \\citep{weaver02}.   In order to investigate the nature of RRL emission in this galaxy, and hence the nature of obscured star formation, we have performed high resolution observations of RRLs in NGC 253 at multiple frequencies using the VLA. The 8.3 GHz H92$\\alpha$ line and the 15 GHz H75$\\alpha$ line were observed using the VLA in the A configuration, which constitute the highest resolution RRL observations to date (0.4\\asec and 0.2\\asec respectively). These line emissions from the radio nucleus were studied separately and have been described in detail in \\citet{mohan02}. The current paper deals with the remainder of this cm-wave RRL emission. We have carried out these observations in order to (1) locate the sites of high frequency RRL emission and identify radio continuum counterparts to the line emitting sources, (2) to investigate the nature of the radio nucleus and the IR peak, based on their RRL emission, and (3) separate the diffuse and compact RRL emission in an attempt to construct a multi-density model of the ionised gas in NGC 253, using multi-frequency RRL data, similar to Arp 220 \\citep{anan00}. The 1.4 GHz H166$\\alpha$ line emission was observed at a resolution of $\\sim$4\\asec using the VLA in the BnA configuration and the C configuration data presented in \\citet{ag90} were re-analysed. The aim was to better localise the large-scale line emission and hence (1) constrain the projected lateral size of the line emitting gas and thus obtain more accurate models, (2) investigate the relative importance of externally stimulated emission versus internal (spontaneous) emission at 1.4 GHz compared to higher frequencies and (3) check the reality of the observed position offset between the low frequency and high frequency RRL emission found by \\citet{ag90} using the C and D configuration data. ",
        "conclusions": "The 1.4~GHz line emission is detected with a resolution of $\\sim$50 pc. The 8.3~GHz \\hx and 15~GHz \\hu line emissions have also been imaged, at high resolutions ($\\sim$4 pc and $\\sim$2 pc respectively). The total \\hx emission, which is about half the value measured at low resolution, consists of an extended region of emission surrounding the radio nucleus (the `15~pc region') and four compact sources in the outer disk of NGC 253. The line emission in the 15~pc region can be decomposed into four additional compact sources. The peak of the line emission is towards the radio nucleus. The line widths of the outer sources is $\\leq$100 km~\\s whereas the line widths of the sources in the inner 15~pc are about 200 km~\\sss. \\hu emission has been detected from only two of these sources. Not all the \\hx line sources have identifiable continuum counterparts.  In contrast to existing interpretations, we show that multi-wavelength RRL emission from NGC 253 does not arise in optically thin gas in LTE. The four extranuclear \\hx sources as well as the four compact sources in the 15~pc region are modelled individually and arise in gas at densities close to 10$^4$ \\ccc. The total ionisation and gas mass in this high density Component~I are $\\sim$4$\\times$10$^{52}$ \\s and 5000-10000 \\msunn. The total size subtended by Component~I is $<$10 pc. This component produces a negligible fraction of the 1.4~GHz line emission.  A second component of ionised gas, Component~II, is required to explain all of the observed H166$\\alpha$ line at 1.4~GHz and the remainder of the cm-wave RRL emissions. This gas is a lower density of a few 100~\\ccc. The The typical transverse size is a few tens of pc. A third component of gas, whose density must be higher than 10$^4$ \\ccc, is also needed to explain the observed H40$\\alpha$ emission at 99~GHz, but has not been modelled in this work.  In conclusion, \\begin{itemize} \\item The ionised gas component in NGC 253 is decomposed into three distinct components: a low density component ($\\sim$500 \\ccc) which contributes most of the observed thermal radio continuum emission and almost all of the 1.4~GHz RRL; a higher density component ($\\sim$10$^4$ \\ccc) which is more localised and exhibits supersonic line widths, and a third component at an even much higher density, responsible for mm-wave RRLs. \\item The nature and location of the IR peak is uncertain and probably contains low density gas ionised by a star cluster which is much less obscured than the central 15~pc region. \\item The ionisation of the line emitting high density thermal gas in the 15~pc region is probably determined by obscured star clusters and the dynamics could be partly  determined by the AGN, like the Galactic Centre. Alternatively, some of the compact line sources could be line of sight enhancements in emission measure due to complex dynamics in the region, rather than individual HII regions surrounding massive star clusters. In either case, we have discovered further examples of supernebulae with large line widths sustained by internal gas dynamical processes. \\end{itemize}"
    },
    "0410/astro-ph0410074.txt": {
        "abstract": "This work represent the first major study of the optical and IR characteristics of the mass donor companions to the X-ray pulsars in the Small Magellanic Cloud (SMC).  In this work several new counterparts have been identified, and possible ones confirmed, as companions to X-ray pulsars in the SMC giving a total of 34 such objects now identified. In addition this work presents three new binary periods and confirms two X-ray periods using optical data for objects in this group. This homogeneous sample has been studied as a group to determine important general characteristics that may offer insight into the evolution of such systems. In particular, the spectral class distribution shows a much greater agreement with those of isolated Be stars, and appears to be in some disagreement with the galactic population of Be stars in Be/X-ray binaries. Studies of the long term optical modulation of the Be star companions reveal an extremely variable group of objects, a fact which will almost certainly make a major contribution to the pronounced X-ray variability. The spatial distribution of these systems within the SMC is investigated and strongly suggests a link between massive star formation and the HI density distribution.  Finally, studies of the circumstellar disk characteristics reveal a strong link with optical variability offering important clues into the long-term stability of such disks. ",
        "introduction": "The Be/X-ray systems represent the largest sub-class of massive X-ray binaries.  A survey of the literature reveals that of the 115 identified massive X-ray binary pulsar systems (identified here means exhibiting a coherent X-ray pulse period), most of the systems fall within this Be counterpart class of binary.  The orbit of the Be star and the compact object, presumably a neutron star, is generally wide and eccentric.  X-ray outbursts are normally associated with the passage of the neutron star close to the circumstellar disk (Okazaki \\& Negueruela, 2001). A recent review of these systems may be found in Coe (2000).  A Be star is a non supergiant hot star, with a B spectral type, whose spectrum has, or had at some time, one or more Balmer lines in emission.  The optical light from a Be/X-ray binary is dominated by the mass donor star in the blue end of the spectrum, but at the red end there is normally a significant contribution from the circumstellar disk. Long term optical observations provide valuable insights into the behaviour of the circumstellar disk, and hence into some of the details of the binary interactions within the system.   X-ray satellite observations have revealed that the SMC contains an unexpectedly large number of High Mass X-ray Binaries (HMXB). At the time of writing, 47 known or probable sources of this type have been identified in the SMC and they continue to be discovered at a rate of about 2-3 per year, although only a small fraction of these are active at any one time because of their transient nature.  Unusually (compared to the Milky Way and the LMC) all the X-ray binaries so far discovered in the SMC are HMXBs, and equally strangely, only one of the objects is a supergiant system, all the rest are Be/X-ray binaries.   The X-ray properties of many of the SMC Be/X-ray binaries were presented in the works of Majid, Lamb \\& Macomb (2004) and Haberl \\& Pietsch (2004). In those paper the authors summarised the known data on the high energy emission from the accreting pulsars. Haberl \\& Pietsch (2004) also proposed several optical counterparts based upon positional coincidences between the X-ray and optical catalogues. In this work several different/new counterparts are identified and the detailed optical properties of most of the sample of 48 objects are presented. Comparisons with the galactic sample of such objects are also made. Since the extinction and distance to the SMC is well documented, this provides an excellent opportunity to study an homogeneous population of objects in a different environment to that of the galaxy.  A list of all known X-ray pulsars may be found in Table~\\ref{tab2}. To minimise confusion arising from long, similar source names only the short SXP identities will be used in this paper. This identity is created simply from the the acronym SXP ({\\bf S}mall magellanic cloud {\\bf X}-ray {\\bf P}ulsar) followed by the pulse period in seconds to three significant figures.   % Converted from Microsoft Word to LaTeX % by Chikrii Softlab Word2TeX converter (version 3.0) % Copyright (C) 1999-2003 Chikrii Softlab. All rights reserved. % http://www.chikrii.com % mailto: info@chikrii.com  % Warning: You are using Chikrii Softlab Word2TeX in TRIAL mode! %          In TRIAL mode some restrictions will apply. %          For more information please visit http://www.chikrii.com % YOU CAN USE THIS FILE WITH THE SOLE PURPOSE OF EVALUATING Word2TeX.  %\\documentclass{article} %\\usepackage{latexsym}  %\\begin{document}  \\begin{table*} \\begin{center} \\begin{tabular}{|l|l|l|l|} \\hline \\textbf{Short ID}& \\textbf{RA (2000)}& \\textbf{Dec (2000)}& \\textbf{Full or alternative name} \\\\ \\hline & & & \\\\ \\hline SXP0.09& 00:42:35.0& -73:40:30.0& AX J0043-737  \\\\   SXP0.72& 01:17:05.2& -73:26:36.0& SMC X-1  \\\\  SXP0.92& 00:45:35.0& -73:19:02.0& PSR J0045-7319 \\\\  SXP2.16& 01:19& -73:12& XTE SMC2165 \\\\  SXP2.37& 00:54:34.0& -73:40:43.0& SMC X-2  \\\\  SXP2.76& 00:59:12.8& -71:38:44.0& RX J0059.2-7138  \\\\  SXP3.34& 01:05:02.0& -72:11:00.0& AX J0105-722, RX J0105.1-7211, [MA93]1506 \\\\  SXP4.78& 00:52& -72:19& XTE J0052-723 \\\\  SXP6.85& 01:01& -72:43& XTE J0103-728 \\\\  SXP7.78& 00:52:07.7& -72:25:43.7& SMC X-3  \\\\  SXP8.02& 01:00:42.8& -72:11:32.0& CXOU J010042.8-721132  \\\\  SXP8.80& 00:51:52.0& -72:31:51.7& RX J0051.8-7231,1E0050.1-7247, [MA93]506  \\\\  SXP9.13& 00:49:13.6& -73:11:37.0& AX J0049-732, RX J00492-7311  \\\\  SXP15.3& 00:52:15.3& -73:19:14.0& RX J0052.1-7319, [MA93]552 \\\\  SXP16.6& 00:50& -73:16& XTE J0050-732 \\\\  SXP18.3& 00:55& -72:42& XTE SMC pulsar XTE J0055-727 \\\\  SXP22.1& 01:17:40.5& -73:30:52.0& RX J0117.6-7330,X Nova92, [MA93]1845  \\\\  SXP31.0& 01:11:09.0& -73:16:46.0& XTE J0111.2-7317 \\\\  SXP34.1& 00:55:27.9& -72:10:58.0& CXOU J005527.9-721058, RXJ0055.4-7210 \\\\  SXP46.4& & & XTE SMC pulsar  \\\\  SXP46.6& 00:53:53.8& -72:26:35.0& 1WGA 0053.8-7226,XTE J0053-724 \\\\  SXP51.0& & & XTE SMC pulsar  \\\\  SXP59.0& 00:54:57.4& -72:26:40.3& RX J0054.9-7226,XTE J0055-724, [MA93]810 \\\\  SXP74.7& 00:49:04.6& -72:50:53.0& RX J0049.1-7250, AX J0049-729 \\\\  SXP82.4& 00:52:09& -72:38:03& XTE J0052-725, MACS J0052-726\\#004 \\\\  SXP89.0& & & XTE SMC pulsar  \\\\  SXP91.1& 00:50:55.8& -72:13:38.0& AX J0051-722 ,RX J0051.3-7216, [MA93]413 \\\\  SXP95.2& 00:52& -72:45& XTE SMC95 pulsar \\\\  SXP101& 00:57& -73:25& AX J0057.4-7325, RX J0057.3-7325 \\\\  SXP138& 00:53:23.8& -72:27:15.0& CXOU J005323.8-722715, [MA93]667 \\\\  SXP140& 00:56:05.2& -72:22:00.0& XMMU J005605.2-722200,2E0054.4-7237, [MA93]904 \\\\  SXP144& & & XTE SMC pulsar  \\\\  SXP152& 00:57:50.3& -72:07:56.0& CXOU J005750.3-720756, [MA93]1038 \\\\  SXP165& & & XTE SMC pulsar \\\\  SXP169& 00:52:54.0& -71:58:08.0& XTE J0054-720, AX J0052.9-7158, [MA93]623 \\\\  SXP172& 00:51:50.0& -73:10:40& AX J0051.6-7311, RX J0051.9-7311, [MA93]504 \\\\  SXP202& 00:50:20.8& -72:23:16.0& 1XMMU J005920.8-722316 \\\\  SXP264& 00:47:23.7& -73:12:27.0& XMMUJ004723.7-731226, RXJ0047.3-7312,[MA93]172 \\\\  SXP280& 00:58:08.0& -72:03:30.0& AX J0058-72.0, [MA93]1036  \\\\  SXP293& 00:50& -73:06& XTE J0051-727  \\\\  SXP304& 01:01:02.7& -72:06:58.0& CXOU J010102.7-720658, [MA93]1240,RXJ0101.0-7206 \\\\  SXP323& 00:50:44.8& -73:16:06.0& AX J0051-73.3,RXJ0050.7-7316, [MA93]387 \\\\  SXP348& 01:03:13.0& -72:09:18.0& SAX J0103.2-7209, RX J0103-722, [MA93]1367 \\\\  SXP452& 01:01:20.5& -72:11:18.0& RX J0101.3-7211, [MA93]1257 \\\\  SXP504& 00:54:55.6& -72:45:10.0& CXOU J005455.6-724510,RXJ0054.9-7245,AXJ0054.8-7244, [MA93]809 \\\\  SXP565& 00:57:36.2& -72:19:34.0& CXOU J005736.2-721934, [MA93]1020 \\\\  SXP701& 00:55:17.9& -72:38:53.0& XMMU J005517.9-723853 \\\\  SXP756& 00:49:40.0& -73:23:17.0& AX J0049.4-7323, RX J0049.7-7323, [MA93]315 \\\\  \\end{tabular}  \\caption{List of known X-ray pulsars in the SMC. The [MA93] object numbers come from Meyssonnier \\& Azzopardi, 1993. Detailed lists of published references for many of the sources may be found in Haberl \\& Pietsch (2004).} \\label{tab2} \\end{center} \\end{table*}  %\\end{document} ",
        "conclusions": "The X-ray pulsar population of the SMC is providing an excellent homogeneous sample of such objects at the same distance and reddening. Consequently it is possible to investigate the physical characteristics of such high mass X-ray binaries and put the conclusions in the context of a differening environment to the galactic population. From this work the following conclusions have been drawn:  \\begin{itemize}  \\item The spectral class distribution shows a much greater agreement with those of isolated Be stars, and appears to be in some disagreement with the galactic population of Be stars in Be/X-ray binaries. Future accurate spectral measurements will provide better determinations of the spectral class than the photometry presented in this work.  \\item Studies of the long term optical modulation of the Be star companions reveal an extremely variable group of objects, a fact which is almost certainly linked to the observed large X-ray variability. The common underlying link is the rate of mass outflow from the Be star and its contribution to the circumstellar disk.  \\item The spatial distribution of these systems within the SMC strongly suggests a link between massive star formation and the HI density. Though the line of sight density needs to be interpretted in three dimensional terms, the results do point to a clear preference for massive star formation under specific conditions.  \\item Finally, studies of the circumstellar disk characteristics (equivalent width, optical variability, infra-red colours) reveal a strong link with the degree of optical variability offering potential clues into the long-term stability of such disks.  \\end{itemize}  As further observations increase the population of these objects in the SMC these initial conclusions can be further examained. Given the current growth rate of discoveries, the final population numbers may well exceed the galactic population by a factor of ten, and hence they will becoming increasingly important as reference population for understanding the evolution of such massive/neutron star binaries."
    },
    "0410/astro-ph0410395_arXiv.txt": {
        "abstract": "We study neutrino emission from a normal, ungapped, quark phase in the core of a compact star. Neutrino emission from non-interacting quark matter leads to an emissivity that scales as $\\epsilon\\sim T^7$. We show that the emissivity is enhanced by a combination of Fermi liquid and non-Fermi liquid effects. Fermi liquid effects lead to an emissivity that scales as $\\epsilon\\sim \\alpha_s T^6$, as originally shown by Iwamoto. We demonstrate that non-Fermi liquid effects further enhance the rate, leading to $\\epsilon\\sim \\alpha_s^3 T^6\\log(m/T)^2$, where $m$ is the electric screening scale and $m\\gg T$ under the conditions found in compact stars. We show, however, that combined with non-Fermi liquid effects in the specific heat the enhancement in the emissivity only leads to a modest reduction in the temperature of the star at late times. Our results confirm existing bounds on the presence of ungapped quark matter in compact stars. We also discuss neutrino emission from superconducting phases with ungapped fermionic excitations. ",
        "introduction": "\\label{sec_intro}  Compact stars provide a unique opportunity to study cold and dense strongly interacting matter and its rich phase structure. An important source of information about the structure of compact stars is their cooling behavior \\cite{Pethick:1991mk,Lattimer:2004pg,Yakovlev:2004iq,Weber:2004kj,Page:2004fy}. For the first $\\sim 10^5$ years after the star is born neutrino emission from the bulk is the most efficient energy loss mechanism. Since the matter in the interior of the star is almost degenerate the exact nature of the cooling mechanism is very sensitive to the structure of the low energy excitations. As a result the cooling behavior places important constraints on the phase diagram of dense matter.  Cooling mechanisms are generally grouped into fast processes with an emissivity $\\epsilon\\sim T^6$ and slow processes with $\\epsilon\\sim T^8$. Slow processes include the modified Urca process \\cite{Friman:zq} and neutrino bremsstrahlung. Examples for fast mechanisms are the direct Urca process \\cite{Lattimer:ib}, neutrino emission from pion or kaon condensates \\cite{Maxwell:1977,Thorsson:1995rk}, and neutrino emission from an ungapped quark phase \\cite{Iwamoto:eb}. In a fully gapped phase of nuclear or quark matter the emissivity is exponentially small. Near the critical temperature for superconductivity the emissivity is dominated by pair breaking and recombination with a temperature behavior that is intermediate between fast and slow mechanisms \\cite{Flowers:1976,Jaikumar:2001hq}.  In this paper we wish to focus on the emissivity of a possible quark matter phase. At asymptotically high baryon density quark matter is in the color-superconducting CFL phase \\cite{Alford:1998mk}. In this phase all quark excitations are gapped and the emissivity is dominated by exponentially small processes involving massive pseudo-Goldstone modes \\cite{Jaikumar:2002vg,Reddy:2002xc,Kundu:2004mz}. At densities that are relevant to compact stars distortions of the CFL due to the non-zero strange quark mass cannot be neglected and the phase structure is much more complicated. Most of the phases that have been proposed in the literature involve quarks with vanishing or very small gaps. Examples are the 2SC phase \\cite{Bailin:1984bm,Alford:1998zt,Rapp:1998zu} in which one of the colors remains ungapped, single flavor spin-one superconductors that have very small gaps \\cite{Schafer:2000tw,Schmitt:2003aa}, and gapless CFL or 2SC phases \\cite{Alford:2003fq,Shovkovy:2003uu}.  In the following we will concentrate on the neutrino emissivity of ungapped quark matter. This problem is important both in connection with the (almost) gapless phases mentioned above and as a benchmark in order to exclude the presence of ungapped quark matter in compact stars. The emissivity of a normal quark phase was first studied by Iwamoto \\cite{Iwamoto:eb}, see also \\cite{Iwamoto:1982,Burrows:1980ec,Burrows:1979uj}. Iwamoto showed that the direct Urca process is strongly suppressed in non-interacting quark matter, but that Fermi liquid corrections in interacting matter lead to a fast rate $\\epsilon\\sim \\alpha_s T^6$, where $\\alpha_s$ is the strong coupling constant.  It is also known, however, that unscreened color magnetic interactions lead to a breakdown of the Fermi liquid description at temperatures $T_{nfl}\\sim m \\exp(-9 \\pi/(4\\alpha_s))$, where $m^2= N_F \\alpha_s \\mu^2/ \\pi$ is the electric screening scale \\cite{Holstein:1973,Baym:uj}. This scale is very small compared to ordinary QCD scales, but it is large compared to the temperature of neutron stars after the first minute or so. For $\\mu=500$ MeV we have $T_{nfl}\\sim 500$ keV. It was recently shown that non-Fermi liquid effects lead to large corrections of the specific heat of degenerate quark matter for $T\\ll T_{nfl}$ \\cite{Ipp:2003cj}. In the present work we compute non-Fermi liquid corrections to the emissivity and study the cooling behavior of degenerate ungapped quark matter. ",
        "conclusions": "\\label{sec_sum}  In this work, we have discussed the influence of non-Fermi liquid effects on the cooling behavior of compact stars containing quark matter in the normal phase. Non-Fermi liquid effects lead to a logarithmic enhancement in both the neutrino emissivity and the specific heat. The net result of these two effects is a mild logarithmic enhancement in the cooling rate. As our rate is even larger than the Iwamoto rate we confirm and sharpen earlier bounds on the existence of ungapped quark matter in neutron stars \\cite{Blaschke:2000dy,Page:2000wt,Prakash:2000jr,Slane:2002ta}. More quantitative statements will require detailed studies of realistic models in which the quark core is in contact with a hadronic phase or an atmosphere. This is beyond the scope of our investigation.  \\begin{table} \\begin{center} \\begin{tabular}{|c|c|c|} \\hline phase &  emissivity  & specific heat \\\\ \\hline\\hline normal phase &  $\\epsilon\\sim \\alpha_s^3 T^6\\log^2(T)$  & $c_v \\sim \\alpha_s \\mu^2 T\\log(T)$  \\\\ \\hline CFL   &  $\\epsilon \\sim T^{3/2} \\exp(-m_K/T)$    & $c_v \\sim T^3$  \\\\ \\hline 2SC   &   $\\epsilon\\sim \\alpha_s T^6$   & $c_v \\sim \\mu^2 T$  \\\\ \\hline 2SC+1SC   &   $\\epsilon\\sim \\alpha_s T^6$   & $c_v \\sim \\mu^2 T$  \\\\ \\hline $[1{\\rm SC}]^3$   &  $\\epsilon\\sim \\alpha_s T^6$   & $c_v \\sim \\mu^2 T$  \\\\ \\hline gCFL  & $\\epsilon\\sim \\alpha_s T^{5.5}$  & $c_v \\sim \\mu^2 \\sqrt{\\Delta T}$  \\\\ \\hline gCFLK & $\\epsilon \\sim  \\alpha_s T^8$ & $c_v \\sim \\mu^2 T$   \\\\ \\hline g2SC  &  $\\epsilon\\sim \\alpha_s T^6$   & $c_v \\sim \\mu^2 T$  \\\\ \\hline \\end{tabular} \\vspace{0.2cm} \\caption{Dominant cooling rates in different phases of quark matter. This table is a summary of the discussion in Sect.~\\ref{sec_sum}. In the entry $[1{\\rm SC}]^3$ we have assumed that ungapped modes exist and the Urca process is allowed. This is not true in the CSL phase. In the gCFL phase we have used the estimate of Kouvaris et al.~and in the gCFLK phase we have assumed that only slow processes such as neutrino bremsstrahlung are allowed. In the g2SC phase we have assumed that the Fermi liquid direct Urca rate can be used. In both the gCFL and g2SC phase we have ignored possible instabilities signaled by a negative current-current correlation function. } \\label{tab_cool} \\end{center} \\end{table}  Let us now discuss the importance of non-Fermi liquid effects in partially gapped color superconducting phases, see Table \\ref{tab_cool}. The simplest case is the 2SC phase. The 2SC phase can arise when the difference between the Fermi momenta of the strange quark and the up and down quarks is too large for strange-non-strange pairing to occur \\cite{Schafer:1999pb,Alford:1999pa}. There also regions in the phase diagram where $us$ (dSC) or $ds$ (uSC) pairing might occur \\cite{Iida:2003cc,Fukushima:2004zq,Ruster:2004eg}. The 2SC phase is characterized by a partial Higgs mechanism. Color $SU(3)$ is broken according to $SU(3)\\to SU(2)$ and five out of eight gluons acquire a mass. The gapless fermions of the third color interact via screened gauge bosons. As a consequence, there are no non-Fermi liquid corrections in either the specific heat or the neutrino emissivity and the standard result of Iwamoto applies \\cite{Iwamoto:eb}.  If the 2SC phase occurs in quark matter with three flavors then the unpaired quark flavor can form a spin one condensate. Spin one condensates are interesting from the point of view of cooling because the typical gaps are much smaller than the spin zero gap. In particular, we can have $T_c \\simeq T_{nfl}$ and non-Fermi liquid effects in the specific heat are important. Whether or not gapless modes occur below $T_c$ depends on the exact nature of the spin one condensate \\cite{Schafer:2000tw,Schmitt:2003aa}. In a phase with CSL pairing in both $LL$ and $LR$ chirality channels the order parameter is isotropic and no gapless modes exist. In all other phases not all fermions are gapped and the order parameter has nodes on Fermi surface. If spin one pairing takes place in all flavor sectors (referred to as $[1{\\rm SC}]^3$ in Table \\ref{tab_cool}) these modes will dominate the cooling behavior.  The cooling behavior of the recently proposed gapless CFL \\cite{Alford:2003fq} and 2SC \\cite{Shovkovy:2003uu} phases is a more difficult question. In the gCFL phase there are gapless charged and neutral modes and the direct Urca process is possible. Alford et al.~\\cite{Alford:2004zr} have argued that the (almost) quadratic dispersion relation of one of the gapless modes leads to an enhancement of the cooling rate by a factor $\\sqrt{T}$. A similar enhancement occurs in the specific heat \\cite{Alford:2003fq}. Kryjevski and Sch{\\\"a}fer studied quark modes in a kaon condensed CFL phase \\cite{Kryjevski:2004jw}. They find at most one gapless mode. As a consequence, the dominant cooling mechanism is expected to be a slow process such as neutrino bremsstrahlung. However, there are several light modes that will contribute to neutrino transport at MeV temperatures. In the g2SC phase there are gapless up and down quark excitations and the direct Urca process is possible.  We should stress that our discussion of neutrino emission from the g2SC and gCFL phase ignores potential instabilities related to a negative current-current correlation function \\cite{Wu:2003,Huang:2004bg,Casalbuoni:2004tb}. Wu and Yip as well as Giannakis and Ren \\cite{Wu:2003,Giannakis:2004pf} have argued that these instabilities indicate that the correct ground state is an inhomogeneous superconductor of the type discussed by Larkin, Ovchninkov, Fulde and Ferrell \\cite{Larkin:1964,Fulde:1964,Alford:2001ze}. The LOFF phase also contains gapless fermions and the calculation of the neutrino emissivity from this phase in an important problem for future studies.  {\\bf Note added in proof:} \\\\ Since the first version of this work the scale in the logarithm of equ. (\\ref{sig_1l}) has been computed by Gerhold and Rebhan \\cite{Gerhold:2005uu} and the result $\\Lambda_\\Sigma \\approx 8 m$ turned out to be substantially larger than the corresponding scale in the specific heat $\\Lambda_c \\approx 0.28 m$ we have employed as an estimate for the scale in the emissivity. This leads quantitatively to a further enhancement of the neutrino emissivity but does not change the cooling behavior significantly."
    },
    "0410/astro-ph0410110_arXiv.txt": {
        "abstract": "Motivated by the recent interest in cosmologies arising from energy density modifications to the Friedmann equation, we analyse the scaling behaviour for a broad class of these cosmologies comprised of scalar fields and background barotropic fluid sources. In particular, we determine the corresponding scalar field potentials which lead to attractor scaling solutions in a wide class of braneworld and dark energy scenarios. We show how a number of recent proposals for modifying the Friedmann equation can be thought of as being dual to one another, and determine the conditions under which such dualities arise. ",
        "introduction": "The Friedmann equation is one of the cornerstones of modern cosmology relating the expansion of the Universe to the total energy density within it. It forms the starting point for almost all investigations in cosmology. However, over the past few years, possible corrections to the Friedmann equation have been derived or proposed in a number of different contexts, generally inspired by braneworld investigations. These corrections are often of a form that involves the total energy density $\\rho$, and are such that they tend to play a role early in the history of the Universe, fading away as we enter the late--time, post--nucleosynthesis era (although that is not always the case as we shall see). Up to now, the different models have been presented in the literature without any attempt to relate them. In this paper, by introducing a generalised form for the correction, we will provide a formalism which allows us to relate a large class of modified Friedmann cosmologies. Assuming the total energy density to be comprised of a canonical scalar field $\\phi$ with potential $V(\\phi)$, together with some form of barotropic fluid, we will demonstrate how the existence of scaling solutions determines the form of $V(\\phi)$, and in doing so we will establish a direct relation between the form of the potential and the functional form of the modification to the Friedmann equation. Scaling (attractor) solutions in cosmology are very important because they allow one to understand the asymptotic behaviour of a particular cosmology and to determine whether such behaviour is stable or not. They have also been advocated by a number of authors as a way of establishing the behaviour of general scalar fields in a cosmological setting  both in the context of conventional Friedmann cosmologies \\cite{Ratra_Peebles,Steinhardt,trac,CLW,ferreira,Liddle,moduli_stabilization,sahni, delaMacorra:1999ff,Ng:2001hs} and in particular classes of modified Friedmann cosmologies \\cite{Copeland:2000hn,Huey:2001ae,Sahni:2001qp,Huey:2001ah,Majumdar:2001mm,Nunes:2002wz,vandenhoogen,Lidsey:2003sj,Tsujikawa:2003zd,Sami:2004ic,Savchenko:2002mi,Calcagni:2004bh,Tsujikawa:2004dp}.  In section II we present the equations of motion arising out of the modified Friedmann equations and introduce variables which allow the scaling solutions to be determined. The general conditions for scaling behaviour are then established in Section III and we show that these can be written as a closed form relationship between the scalar field and the functional form of the modification to the Friedmann equation. In section IV we demonstrate the existence of duality symmetries between different scaling solutions, and determine the conditions which must be satisfied in terms of the modification to the Friedmann equation for such duality properties to be obtained. Section V then applies the results of the previous sections to a general class of models, evaluating the scaling potentials (and their duals), as well as the explicit evolution of the scale factor in the scaling regimes. In particular, we apply the technique to a number of recently investigated cosmologies of relevance both to braneworld and dark energy scenarios, including the Randall--Sundrum \\cite{R_S}, Shtanov--Sahni \\cite{Shtanov:2002mb} and Cardassian \\cite{Freese:2002sq} models. The Dvali-Gabadadze-Porrati \\cite{Dvali:2000hr} braneworld scenario is investigated in Section VI and we summarize our results in section VII. Throughout units are chosen such that $\\hbar = c =1$. ",
        "conclusions": "In this paper we have brought together a number of recent approaches to cosmology which involve modifications of the Friedmann equation. By introducing the general function $L(\\rho)$ as the way of parametrising the modification, we have been able to establish the conditions under which the new system enters scaling solutions. Considering the case where the energy density is comprised of a scalar field and background barotropic fluid, we have obtained the general relationship that would have to be satisfied between the  evolving scalar field and  $L(\\rho)$. In particular, we have obtained the corresponding potential $V(\\phi)$ which leads to scaling solutions and, for a rather general class of functions of $L(\\rho)$, we have shown that there exist dual solutions which also exhibit similar scaling behaviour. This has allowed us to relate solutions which would otherwise appear quite distinct, including those involving collapsing and expanding cosmologies.  Moreover, the duality can directly relate singular and non--singular cosmologies. To illustrate this property, consider a particular scaling solution $a_+(\\phi )$ that is singular in the sense that the scale factor vanishes at $a_+(0) =0$. (The value of the scalar field can be chosen to be $\\phi =0$ without loss of generality). Suppose, however, that the logarithmic  derivative of the scale factor with respect to the field is non--zero at this point, $d \\ln a_+/d\\phi |_0 \\ne 0$, and furthermore, that the scale factor is a monotonic function with a finite first derivative for all physical (non--zero) values of the field. These properties are satisfied, for example, in the R-S II and Cardassian models.  The qualitative behaviour of the dual solution, $a_-(\\phi )$, is then determined from Eq. (\\ref{dualscale}). If we define a new parameter $\\epsilon (\\phi )\\equiv [d \\ln a /d\\phi ]^{-1}$, the value of the scale factor is simply given by the area under the curve $\\epsilon_+ (\\phi )$, where the field evolves from zero to some value $\\phi$. (We are assuming implicitly that $\\dot{\\phi} >0$ and $p<0$ again without loss of generality). However, due to the exponential nature of Eq. (\\ref{dualscale}) the initial value of the dual scale factor is {\\em non--zero} and this results in a non--singular background.  On the other hand, the time reversal of the seed solution $a_+(\\phi )$ would result in a collapsing dual model where the limits in the integral (\\ref{dualscale}) are taken from $\\phi$ to zero. Consequently, the dual solution can be analytically continued through $\\phi =0$ into a contracting phase. In this sense, therefore, any singular (expanding) scaling solution satisfying the above (very weak) conditions can generate a non--singular bouncing cosmology, where the latter is associated with a combination of the seed solution and its time reversal. For fixed values of $\\{ p, \\lambda_{\\pm} \\}$, the collapsing phase of the bouncing solution will be unstable if the expanding phase is stable, and vice--versa. In principle, however, different seed solutions may be employed to generate distinct and stable collapsing and expanding branches that can be smoothly joined at the bounce.  This opens up the possibility that such dualities will allow us to relate singular cosmologies to non-singular bouncing cosmologies, a topic presently of considerable interest in cosmology.  Finally, we have argued that the type of correction given by Eq. (\\ref{L_form}) arises in a number of particle physics motivated models. Further examples arise in the limit where $A\\rho^{\\nu} \\gg 1$. In particular, the case $\\nu =-1/3$ corresponds to the high--energy limit of the Gauss--Bonnet braneworld \\cite{gb}. In this model, the R-S II scenario is generalized to include a Gauss--Bonnet combination of curvature invariants in the five--dimensional bulk action. More generally, effective Friedmann equations of the form $H^2 \\propto \\rho^{\\nu}$, where $\\nu$ is arbitrary \\cite{chung}, can arise in models based on Ho\\v{r}ava--Witten theory compactified on a Calabi--Yau three--fold \\cite{H_W}. Generalized scaling solutions driven by corrections of this form were recently investigated for a variety of scalar field models \\cite{Tsujikawa:2004dp}.  Future directions involving the use of the duality properties of these models would include an extension of our analysis to negative potentials, thereby allowing us to link these classes of solutions with those arising in the cyclic/ekpyrotic scenario \\cite{Steinhardt:2001st}. On the other hand, as we have seen, the more general corrections proposed in \\cite{Dvali:2000hr,Deffayet:2000uy} that arise due to modifications of gravity on large scales can lead to an explanation of the present cosmic acceleration without introducing dark energy \\cite{Lue:2004za}. It would be interesting to investigate the impact that the duality transformations we have described have on such a scenario."
    },
    "0410/astro-ph0410326_arXiv.txt": {
        "abstract": "{ We use a complete AGN sample (Ho et al. 1997) to study the environment of Seyferts, LINERs and HII galaxies. For each AGN we search for companions in the UZC redshift catalogue and compute local as well as large scale neighbour density and distance to the nearest neighbour. We find that on small scale ($\\sim$ 0.2 $h^{-1}$ Mpc) LINERs exhibit denser environments than Seyferts and HII galaxies at 3$\\sigma$ significance level. The same result is not confirmed when densities are computed on large scales. LINERs also exhibit closer nearest neighbours than Seyferts and  HII galaxies (at 2$\\sigma$ c.l.). However, when excluding AGNs in early type galaxy hosts, the neighbour density characteristics of LINERs Seyferts and HII galaxies turn out to be similar, a result that confirms that the excess of neighbours around LINERs is most likely due to a morphology-density effect. } \\authorrunning{V.Zitelli et al.} \\titlerunning{AGNs and environment} ",
        "introduction": "It is not yet fully understood whether the dynamical effect exerted by interaction with other galaxies is one of the major cause of the efficient gas fueling observed in AGNs. Indeed, the claimed excess of companions in Seyfert galaxies (Dahari 1984, Laurikainen \\& Salo 1995) is not confirmed \\citep{Bushouse87,DeRob98} at better than 95\\% significance level.  Examination of the frequency of the occurrence of Seyferts in pairs and groups (Kelm et al. 1998)  has further shown that Seyferts constitute a $\\sim$2\\% in all galaxy systems, whereas the fraction of Seyferts in UZC-Compact Groups (CGs) turns out to be 3 times as large as among isolated UZC galaxies (Kelm et al. 2004). Interestingly, no significant differences in dynamical properties are found, when comparying galaxy systems which host, or not, a Seyfert member. For Seyferts, it has also been shown (Malkan et al. 1998) that there is little direct evidence for unusually high rates of interaction; less than 10\\% show tidal feature or multiple nuclei. And Seyferts in UZC-CGs do not seem to be more likely than other CG galaxies to display major interaction patterns (or a bar), which  is consistent with the results of earlier papers (Keel 1996, Ho et al . 1997) all finding no clear relationship between the presence of AGNs and detailed morphological properties.  Concerning environmental differences between LINERs and Seyferts Kauffmann et al. (2004) have used a complete sample of galaxies drawn from the SDSS to suggest that the fraction  of low luminosity AGN (LINERs) depends very little on local density, whereas the fraction of strong AGN in massive galaxies decreases as a function of density. And Schmitt (2001) has shown that LINERs exhibit a higher percentage of companions than Seyferts and HII galaxies, but that differences disappear when considering only galaxies of similar morphological type. \\section {The sample} The AGN sample we use here is selected from the master list of 503 nearby galactic nuclei by Ho et al. (1997), obtained from the revised version of the Shapley-Ames Catalog of Bright Galaxies (RSA) and Second Reference Catalog of Bright galaxies (RC2). It is a flux limited catalog of galaxies with magnitude $B_T$ $\\leq$ 12.5. Our AGN subsample includes all AGN with radial velocity in the range [1500$\\div$3000] km$s^{-1}$: 48 LINERs (including transition objects), 15 Seyfert galaxies and 52 HII galaxies.  For each AGN we have automatically identified neighbours in the UZC catalog (Falco et al. 1999) which is complete for galaxies brighter than $m_{zw}$=15.5 The local density of each AGN is estimated by counting all neighbours within a cylinder 0.2 $h^{-1}$ Mpc in projected radius and $\\pm$1000 km/s in depth. The large scale density is estimated by counting all neighbours within a cylinder 1 $h^{-1}$ Mpc in projected radius and $\\pm$1000 km$s^{-1}$ in depth. The distance to the nearest among these neighbours defines the nearest neighbour distance. \\begin{figure} \\centering \\includegraphics[clip=true,angle=0,width=0.55\\hsize]{zitelli.fig1.ps} \\caption{Local density distributions (neighbours within 200 $h^{-1}$kpc) of LINERs, Seyferts and HII galaxies. LINERs exhibit significantly larger local density than Seyfert and HII galaxies. } \\label{fit1} \\end{figure} \\begin{figure} \\centering \\includegraphics[clip=true,angle=0,width=0.55\\hsize]{zitelli.fig2.ps} \\caption{Large scale density distributions of LINERs, Seyferts and HII galaxies. The hyphothesis that LINERs, Seyfert and HII galaxies are drawn from the same population cannot be rejected at the 1$\\sigma$ significance level          } \\label{fit2} \\end{figure} \\section {Neighbour density} In fig.\\ref{fit1} we show the distribution of AGNs as a function of the local neighbour density (neighbours within 0.2 $h^{-1}$Mpc) for Seyferts, Liners and HII galaxies respectively. Distributions show some difference by eye. The U-test indicates,  at a significance level 3$\\sigma$, that the local density distributions of LINERs and Seyferts are drawn from different parent distributions. Significant differences (3$\\sigma$) are also found when comparing (U-test) the local density of LINERs and HII galaxies, whereas Seyferts and HII galaxies do not appear to differ significantly.  In fig.\\ref{fit2} we show the distribution of AGNs as a function of the large scale neighbour density (neighbours within 1 $h^{-1}$Mpc). The large scale density distributions of LINERs, Seyferts and HII galaxies appear similar. The U- test indicates that differences between population are not significant. ",
        "conclusions": ""
    },
    "0410/astro-ph0410175_arXiv.txt": {
        "abstract": "We study post weak decoupling coherent active-sterile and active-active matter-enhanced neutrino flavor transformation in the early universe.  We show that flavor conversion efficiency at Mikheyev-Smirnov-Wolfenstein resonances is likely to be high (adiabatic evolution) for relevant neutrino parameters and energies.  However, we point out that these resonances cannot sweep smoothly and continuously with the expansion of the universe.  We show how neutrino flavor conversion in this way can leave both the active and sterile neutrinos with non-thermal energy spectra, and how, in turn, these distorted energy spectra can affect the neutron-to-proton ratio, primordial nucleosynthesis, and cosmological mass/closure constraints on sterile neutrinos.  We demonstrate that the existence of a light sterile neutrino which mixes with active neutrinos can change fundamentally the relationship between the cosmological lepton numbers and the primordial nucleosynthesis $^4${He} yield. ",
        "introduction": "\\label{sec-1} If light sterile neutrinos exist we will be forced to re-think the role of the weak interaction in the early universe, primordial nucleosynthesis, and cosmology.  Though light sterile neutrinos which mix with active neutrinos long have been a subject of theoretical speculation, the LSND and KARMEN experiments \\cite{LSND} gave rise to particular interest in the mass-squared difference range $0.2\\,{\\rm eV}^2<\\delta m^2_{\\rm as}<100\\,{\\rm eV}^2$.  Here we study for this range of $\\delta m^2_{\\rm as}$ the cosmological lepton number-driven conversion of active neutrinos, $\\nu_\\alpha$ (and/or $\\bar\\nu_\\alpha$) with $\\alpha={\\rm e},\\mu,\\tau$, to a singlet, \\lq\\lq sterile\\rq\\rq\\ neutrino species $\\nu_s$ (or $\\bar\\nu_s$) in the epoch of the early universe after decoupling of the weak interactions, when neutrino spectral distortions are likely to persist.  A positive signal in the on-going mini-BooNE experiment \\cite{miniB}, {\\it i.e.}, confirming the interpretation of the LSND result in terms of vacuum neutrino mixing, sets up an immediate crisis in neutrino physics. Such a result, when combined with the already well established evidence for neutrino mixing at mass-squared differences associated with the atmospheric ($\\delta m^2 \\sim 3\\times{10}^{-3}\\,{\\rm eV}^2$) and solar neutrino ($\\delta m^2 \\sim 7\\times{10}^{-5}\\,{\\rm eV}^2$) anomalies, would suggest the existence of three independent neutrino mass-squared differences which would, in turn, require four neutrino species. Given the $Z^0$-width limit on the number of flavors of neutrinos with standard weak interactions (3), a fourth neutrino would have to be \\lq\\lq sterile,\\rq\\rq\\ with sub-weak interaction strength, {\\it e.g.}, perhaps an SU(2) singlet. The only alternative to this line of reasoning and to this conclusion is the possibility of CPT violation~\\cite{cpt1}.  However, there is no consistency of the neutrino oscillation data with a CPT-violating three-neutrino model at a 3-$\\sigma$ level~\\cite{cpt2}.  Hand in hand with this particle physics dilemma, evidence for a singlet neutrino that mixes with active neutrinos in this mass-squared range also confronts cosmology with a curious and vexing problem. In the standard cosmological model with zero or near-zero net lepton numbers one would expect that matter-supressed neutrino oscillations in the channel $\\nu_\\alpha \\rightleftharpoons\\nu_s$ or in $\\bar\\nu_\\alpha \\rightleftharpoons \\bar\\nu_s$ (where $\\alpha=e,\\mu,\\tau$) proceeding in the regime above weak interaction decoupling ($T > 3\\,{\\rm MeV}$), would efficiently populate seas of singlet neutrinos~\\cite{aba}. The significant additional energy density in these sterile neutrino seas would engender a faster expansion rate for the universe and a consequently higher temperature for Weak Freeze-Out (where the initial isospin of the universe, the neutron-to-proton ratio is set).  A higher Weak Freeze-Out temperature would result in more neutrons and, hence, a higher yield of {$^4${He}}.  A higher predicted abundance of $^4$He {\\it arguably} may be in conflict (or is close to being in conflict) with the observationally inferred upper limit on the primordial helium abundance.  Depending on the helium abundance inferred from compact blue galaxies, an increase in the predicted Big Bang Nucleosynthesis (BBN) $^4$He yield may or may not be disfavored \\cite{aba,cyburt}.  However, the primordial helium abundance is notoriously difficult to extract from the observational data and recent studies point to a fair range for the observationally inferred primordial mass helium fraction: $23\\%$ to $26\\%$ \\cite{OS}. The upper limit of this range is provocatively close to the standard BBN $^4$He mass fraction yield prediction, $24.85\\pm 0.05\\%$, as computed with the deuterium-determined or CMB (Cosmic Microwave Background) anisotropy-determined baryon density.  Additionally, it has been suggested \\cite{murayama} that a fully populated sea of sterile neutrinos and antineutrinos with rest masses $\\sim \\sqrt{\\delta m^2_{\\rm as}}$ could be in conflict with neutrino mass bounds derived from CMB anisotropy limits and large scale structure considerations \\cite{WMAP}. There is a recent analysis of constraints from measurements of galaxy bias stemming from galaxy-galaxy lensing and the inferred linear matter power spectrum derived from the Lyman alpha forest in the Sloan Digital Sky Survey (SDSS) \\cite{Sloan}.  This analysis specifically considers a so-called \\lq\\lq 3+1\\rq\\rq\\ neutrino mass hierarchy, {\\it i.e.,} the scheme which is appropriate for constraining sterile neutrinos. The neutrino mass constraint so derived is somewhat less stringent than constraints in schemes with three neutrinos with degenerate masses. However, the central conclusions of Ref. \\cite{murayama} survive.  Should we someday be confronted with a positive indication of neutrino flavor mixing with mass-squared scale consistent with the range for $\\delta m^2_{\\rm as}$, we will have a problem that would call for modification either of our notions of basic neutrino physics or of the standard cosmological model.  There have been a number of ways proposed to get out of these cosmological difficulties.  For instance, if neutrinos do not acquire mass until after the BBN epoch, as may occur via a late-time phase transition \\cite{flypaper,latetime}, the singlet states will not be populated via oscillations during the BBN era. In addition to schemes involving the epoch of neutrino mass generation \\cite{Chacko:2003dt}, annihilation \\cite{neutrinoless} or decay \\cite{latetime} of the singlet neutrinos (when $m_s \\sim T$) may alleviate or avoid the CMB and large scale structure constraints.  However, chief among the mechanisms proposed to escape the cosmological difficulties associated with singlet neutrinos, is the the invocation of a significant net lepton number \\cite{FV}.  The idea is that the net lepton number gives active neutrinos larger effective masses in medium in the early universe, thereby driving them further off-resonance in the epoch prior to weak decoupling ({\\it i.e.}, $T>3\\,{\\rm MeV}$) and reducing their effective matter mixing angles with the singlet neutrino. In turn, smaller effective matter mixing angles would imply a suppressed production of singlet neutrinos and, hence, a reduced population of the singlet neutrino sea.  This lepton number-induced suppression of active-sterile mixing at high temperature is why we assume here that there is no initial population of the sterile neutrino sea.  The lepton number residing in the sea of $\\nu_\\alpha$ and $\\bar\\nu_\\alpha$ neutrinos ($\\alpha=e,\\mu,\\tau$) is defined in analogy to the baryon number $\\eta \\equiv (n_b-n_{\\bar b})/n_\\gamma$, \\begin{equation} L_{\\nu_{\\alpha}}= {{n_{\\nu_\\alpha}-n_{\\bar\\nu_\\alpha}}\\over{n_\\gamma}} \\label{lepno} \\end{equation} where $n_\\gamma = \\left(2\\zeta(3)/\\pi^2\\right) T^3_\\gamma$ is the proper photon number density at temperature $T_\\gamma$, and where $n_{\\nu_\\alpha}$ and $n_{\\bar\\nu_\\alpha}$ are the number densities of $\\nu_\\alpha$ and $\\bar\\nu_\\alpha$ neutrinos, respectively, at this epoch. After the epoch of $e^\\pm$ annihilation the baryon number is $\\eta \\approx 6\\times10^{-10}$, whereas at earlier epochs it is roughly two and half times larger. We will consider here net lepton numbers which are vastly larger than $\\eta$, so its precise value is of no consequence for our results.  We can insure that the effective matter mixing angles for the oscillation channel $\\nu_\\alpha \\rightleftharpoons\\nu_s$ (or $\\bar\\nu_\\alpha \\rightleftharpoons\\bar\\nu_s$) are sufficiently small to suppress singlet neutrino production if the Mikheyev-Smirnov-Wolfenstein (MSW) \\cite{MSW} resonance temperature is less than the weak decoupling temperature, $T_{\\rm res} < T_{\\rm dec}$. This implies that the lepton number associated with any of the active neutrino flavors should satisfy, \\begin{equation} L > {{10^{-3}}\\over{\\epsilon}} \\left({{2}\\over{N_{\\rm degen}}}\\right) {\\left({{3\\,{\\rm MeV}}\\over{T_{\\rm dec}}}\\right)}^4 \\left({{\\delta m^2_{\\rm as} \\cos2\\theta}\\over{1\\,{\\rm eV}^2}}\\right) \\label{suppress} \\end{equation} where $\\theta$ is the vacuum mixing angle characteristic of $\\nu_\\alpha\\rightleftharpoons\\nu_s$ oscillations, $N_{\\rm degen}$ is the number of neutrino species possessing this lepton number, and where $\\epsilon \\equiv E_\\nu/T$. For neutrinos with typical energies in the early universe ({\\it i.e.}, $\\epsilon \\sim 1$), suppression of singlet neutrino production would require lepton numbers ranging from $L > {10}^{-4}$ for $\\delta m^2_{as} =0.2\\,{\\rm eV}^2$ to $L > 5\\times{10}^{-3}$ for $\\delta m^2_{as} =10\\,{\\rm eV}^2$. Current limits on lepton numbers are $\\vert {L_{\\nu_\\alpha}}\\vert < 0.1$ \\cite{abb} (and possibly even weaker by a factor of two or so if allowance is made for another source of extra energy density in the early universe \\cite{barger,Kneller:2001cd}). Therefore, this avenue for escape from the sterile neutrino conundrum appears to be allowed, albeit at the cost of a huge disparity between the lepton and baryon numbers.  However, this argument overlooks an important point. Though the large lepton number supresses the effective matter mixing angle for $\\nu_\\alpha\\rightleftharpoons\\nu_s$ during the epoch of the early universe where active neutrinos are thermally coupled ($T>T_{\\rm dec}$), it can cause coherent matter-enhancement of this channel at lower temperatures where the active neutrinos rarely scatter and are effectively decoupled. Resonant MSW transformation of active neutrinos to singlets in the channel $\\nu_\\alpha\\rightleftharpoons\\nu_s$ is, however, self limiting. This is because as the universe expands and the resonance sweeps from low toward higher neutrino energy, the conversion of $\\nu_\\alpha$'s decreases the lepton number which, in turn, causes the resonance sweep rate to increase, eventually causing neutrinos to evolve non-adiabatically through resonance and so causing flavor transformation to cease.  At issue then is how many active neutrinos can be converted to sterile neutrinos prior to or during the epoch where the neutron-to-proton ratio is set (\\lq\\lq Weak Freeze-Out\\rq\\rq). If there is a significant conversion, the resultant non-thermal active neutrino energy spectra can cause an increase or decrease (if $\\bar\\nu_\\alpha \\rightleftharpoons \\bar\\nu_s$ is enhanced) in the $^4$He yield and call into question the viability of invoking a large net lepton number to reconcile neutrino physics and BBN. Other but related aspects of transformation-induced nonthermal neutrino spectra effects on primordial nucleosynthesis have been studied in Ref.s~\\cite{Abazajian:1999yr, Kirilova:2001ab}. Likewise, previous studies have considered other aspects of the relationship between sterile neutrinos and BBN \\cite{new}, as well as constraints on sterile neutrinos without a primordial lepton number \\cite{Cirelli:2004cz}.  In any case, non-thermal energy distribution functions for $\\nu_e$ and/or $\\bar\\nu_e$ change the relationship between the BBN $^4${He} yield and the neutrino chemical potentials.  In section II we discuss the physics of active-sterile neutrino flavor transformation in the early universe and point out a key issue in how the MSW resonance sweeps through the neutrino energy distribution functions as the universe expands. The generally high adiabaticity of neutrino flavor evolution is also pointed out in this section.  Simultaneous active-active and active-sterile neutrino flavor conversion, and \\lq\\lq synchronization\\rq\\rq\\ are also discussed in this section.  Possible multi-neutrino mass level crossing scenarios in the early universe are discussed in this section. Sterile neutrino contributions to closure, constraints on this from large scale structure and Cosmic Microwave Background radiation considerations, as well as other sterile neutrino sea population constraints are examined in section III.  In section IV we describe how distorted $\\nu_e$ and/or $\\bar\\nu_e$ distribution functions impact the rates of the lepton capture reactions that determine the neutron-to-proton ratio and the $^4$He yield in BBN. This is then applied in various initial lepton number and neutrino conversion scenarios. Finally, in section V we give conclusions and speculations regarding the neutrino mass and cosmological lepton number insights that would follow in the wake of an experimental signature for a large neutrino mass-squared difference of order the range given for $\\delta m^2_{as}$. Appendix A provides an exposition of the lepton capture rates on free nucleons when, as appropriate, $\\nu_e$ or $\\bar\\nu_e$ energy distribution functions are zero up to some energy, and thermal/Fermi-Dirac at higher energies. ",
        "conclusions": "The most general conclusions that can be drawn from this work is that the existence of one or more light sterile neutrinos could (1) alter the relationship between neutrino chemical potential and primordial nucleosynthesis yields, and (2) leave both active and sterile neutrinos with nonthermal, distorted energy spectra. Attempts to constrain an \\lq\\lq LSND\\rq\\rq\\ sterile neutrino based on conventional degenerate primordial nucleosynthesis considerations, as well as attempts to reconcile this neutrino with BBN limits via a primordial lepton number are now suspect. However, obtaining the detailed relationship between lepton numbers, active-sterile neutrino mixing parameters and light element nucleosynthesis yields that would be required to effect reliable constraints is beyond the scope of the work presented in this paper.  This is mostly a consequence of another discovery made in this work: MSW resonances cannot sweep smoothly and continuously beyond $\\epsilon_{\\rm max}$. We showed that one way the resonance condition and adiabatic conversion criterion can be met beyond this point is if the resonance skips to higher energies a number of times until the initial lepton number is depleted. Partly for this reason resonance evolution is likely to be quite complicated. Complication will also arise because active-active neutrino mixing can be efficient and can occur simultaneously with active-sterile transformation.  In any case, neutrino flavor conversion is almost inevitable once the resonance condition is met. This is on account of another insight presented in this paper: the highly adiabatic nature of neutrino flavor evolution through MSW resonances for the neutrino mass/mixing parameters of most interest and for the conditions in the post Weak Decoupling early universe.  If the mini-BooNE experiment sees a positive signal, confirming the existence of light sterile neutrinos, we will be forced to confront the problems posed in this paper. Likewise, future progress in improving the precision and confidence in the observationally inferred primordial helium abundance coupled with CMB and large scale structure-derived limits on neutrino collisionless damping scales could give us hints about active-sterile neutrino mixing physics in the early universe."
    },
    "0410/astro-ph0410496_arXiv.txt": {
        "abstract": "We present the results of a spatio-kinematic study of the planetary nebula Menzel\\,1 using spectro-photometric mapping and a 3-D photoionization code. We create several 2-D emission line images from our long-slit spectra, and use these to derive the line fluxes for 15 lines, the H$\\alpha$/H$\\beta$ extinction map, and the [SII] line ratio density map of the nebula. We use our photoionization code constrained by these data to derive the three-dimensional nebular structure and ionizing star parameters of Menzel\\,1 by simultaneously fitting the integrated line intensities, the density map, and the observed morphologies in several lines, as well as the velocity structure. Using theoretical evolutionary tracks of intermediate and low mass stars, we derive a mass for the central star of 0.63\\,$\\pm$0.05M$_{\\odot}$. We also derive a distance of 1050\\,$\\pm$150pc to Menzel\\,1. ",
        "introduction": "The Planetary Nebula (PN) Menzel\\,1 (Mz\\,1) or G322.4-02.6 (15$^h$ 34$^m$ 16$^s$.7 -59$^o$ 08' 59'' 2000.0) is a bright object with a bipolar morphology and a prominent central ring of enhanced emission. H$\\alpha$ and [OIII] narrow band images of Mz\\,1 have been published by \\cite{SCM92}. Being bright has not resulted in Mz\\,1 being well-studied; only a few papers have been dedicated to the object.   H$_2$ emission has been detected from Mz\\,1 by \\cite{W88} and the morphology in this molecular line is similar to that of our optical image, shown in Fig.\\,\\ref{pic}.  A detailed kinematical study of Mz\\,1 was done by \\cite{MB98} using high resolution echelle spectroscopy. Using the intensity ratio of the {[}SII{]} lines they computed densities of 1700 $cm^{-3}$ for the main ring and 400 $cm^{-3}$ for the bipolar part. They show that the gas in the main part of the nebula is expanding at 23\\,$km \\, s^{-1}$, and that the velocity structure is consistent with a cylindrical model with expansion velocities proportional to the radial distance from the center. They also show that the dynamical age of the ring in Mz1 is of the order of 7000\\,yrs and estimate a mass of about 0.5M$_{\\odot}$ for the nebula using dynamical arguments and their computed distance of 2.0\\,kpc.  Distances have been determined by several authors such as \\cite{VZ95} with 2.53 kpc, \\cite{CKS92} with 2.28 kpc and \\cite{Z95} with 2.85 kpc, all using statistical methods and based on the same data. Using the same formalism as \\cite{VZ95} and a new value for the radius of the nebula, Marston et al. (1998) obtained a distance of 2.0 kpc with a claimed uncertainty in this estimate of about 30\\%. \\cite{AO92} lists 10 distances of which 8 are statistical and 2 individual determinations. Given that all the above authors use a filling factor, $\\epsilon$\\,=\\,0.5, any such method has an inherent and large uncertainty in the distances they determine. We compute 2.0\\,$\\pm$0.5\\,kpc (adjusted standard deviation) from all 14 literature distances without any weighting factor. Note that each individual distance can be in error by a large factor (see the appendix), and the standard deviation computed from the literature values is not very robust.  An effective temperature of 139\\,kK and a luminosity of 147\\,L$_{\\odot}$ were determined by \\cite{SCS93} with a distance of 1.8\\,kpc. Mz\\,1 is clearly under-luminous for a typical PN.  In this work we present observations of, and a 3-D photoionization model for Mz\\,1, and derive the 3-D structure of the PN constrained by observed fluxes and morphologies in many emission lines, using the same method as \\cite{MSGH04} applied to NGC\\,6369. By determining the 3-D structure of nebulae, the large uncertainty involved in all classical statistical distance determination methods is eliminated. Assuming an arbitrary filling factor, constant ionized mass or diameter, mass-radius relationship etc. is not needed here: we {\\it determine} what the structure and ionized mass are, and can therefore derive distances to much greater accuracy than has been previously possible.  In summary, we obtain the 3-D spatial structure of the nebula along with its chemical composition, ionizing source temperature, luminosity, and mass, as well as an {\\it independent} distance, in a self-consistent manner. In \\S2 we discuss the observations and briefly explain the basic reduction procedures, including our image reconstruction technique used to obtain the emission line intensity images. In \\S3 the results obtained from these images are discussed: the pixel by pixel reddening correction of the images, the integrated line fluxes, and the computed temperature and density maps. In \\S4 we present the model results generated by our 3-D photoionization code, and we discuss the derived quantities. In \\S5 we give our conclusions, and explain in detail how our method works in the appendix. ",
        "conclusions": "We present spectrophotometric maps of Mz\\,1, giving spatially resolved information for many emission lines and precise total fluxes for the observed part ($\\ge$95\\% of the total flux) of the nebula. The images produced with this technique were used to determine properties of the nebula using the H$\\alpha $/H$\\beta$ ratio map for the de-reddening, the {[}SII{]} line ratios for the density, and the {[{NII{]} line map for the temperature.  The H$\\alpha $/H$\\beta$ map shows little structure. The prominent ring feature of the nebula does not show significant differences in extinction when compared to the other regions.  This may imply that most of the reddening is due to the foreground, and not intrinsic to the nebula.  We also show the density map obtained from the observations indicating the presence of a dense waist ring and a bipolar structure of lower density. Based on this map we propose for Mz\\,1 a 3--D hour-glass structure with a waist whose density decreases smoothly from the equator to the poles.  Using a photoionization code and the proposed structure we obtained a complete 3--D physical model for Mz\\,1. The fitted model line intensities show excellent agreement --well within the computed errors-- with the observed values. The obtained distance of d=1050$\\pm$150\\,pc falls near the lower end of the error range determined from 14 literature values for the distance. This makes sense as the angular radius of 25\\arcsec~used in the statistical methods is smaller than the true size of the nebula of 38\\arcsec~ as measured by \\cite{SCM92}, and our distance is also smaller.  The luminosity of 164\\,$\\pm$\\,25\\,L$_{\\odot}$ and temperature of 120\\,$\\pm$\\,16\\,kK we have determined for the central star are within their errors equal to those found by \\cite{SCS93}, taking into account the difference in distance and total flux. Our errors have been conservatively estimated to be similar to the cumulative observational errors we computed and used to constrain the model. This luminosity is low for a PN CS, and indicates that the star is evolved and well down its Sch\\\"{o}nberner track, as confirmed by its position in Fig.\\,\\ref{evolMZ1}.  \\clearpage  \\begin{figure}[h] \\includegraphics[width=\\columnwidth]{f11.eps} \\caption{Comparison of our model temperature and luminosity obtained for the central star with model tracks calculated by \\cite{B95}. \\label{evolMZ1}} \\end{figure}  \\clearpage  We determine the mass of the CS to be 0.63\\,$\\pm$0.05M$_{\\odot}$ by fitting to the theoretical evolutionary tracks of \\cite{B95} , using the errors estimated from our observational data and propagated to the CS luminosity and temperature.  Due to the fact that in this part of the HR diagram, the tracks lie close together, we cannot say with any precision what the precursor mass of this star was, except to place a lower limit on this mass of about 1\\,M$_{\\odot}$ using the added track from \\cite{VW94} in Fig.\\,\\ref{evolMZ1}. The likely value for the CS precursor mass is about 3M$_{\\odot}$, based on the most probable CS mass of 0.63\\,M$_{\\odot}$. From our model output we compute the ionized nebular mass to be 0.14\\,$\\pm$0.03\\,M$_{\\odot}$.  An independent check is the time scale of expansion of Mz1. From \\cite{MB98} and scaling to our distance we get an expansion time of 3500\\,yrs for the ring which can be longer or shorter depending on what the history of the expansion has been -decreasing velocity due to energy conservation in a wind blown cavity would shorten the time and increasing velocity due to expansion into a medium with decreasing density would lengthen it-so a range of about 2000-5000\\,yrs.  For the outer parts of Mz1 this time would be about 4700-12000\\,yrs as the material has had to travel 2.4 times further away from the CS (38\\arcsec~instead of 16\\arcsec). From our luminosity range and computing Bl\\\"{o}cker track ages we obtain 4500 to 10000\\,yrs, quite compatible with the outer nebular expansion time, thus confirming our distance determination.  Since the sum of the ionized nebular mass and the CS mass is 0.77\\,M$_{\\odot}$, we expect that there possibly is more than 2\\,M$_{\\odot}$ of neutral matter near Mz\\,1. Clearly, for the lower CS precursor mass limit of 1\\,M$_{\\odot}$, the neutral mass estimate is a factor of ten smaller. In any case, there is likely more neutral than ionized mass in the system. Note, however, that \\cite{HBCF96} estimate a molecular mass of 0.0086\\,M$_{\\odot}$ and an ionized mass of 0.027\\,M$_{\\odot}$.  They used a radius for Mz1 of 12.6\\arcsec~instead of the 38\\arcsec~optical radius, so naively assuming that the mass scales as the third power of the radius would increase their ionized mass by a factor of up to 9 to 0.25\\,M$_{\\odot}$, which is much closer to our computed mass of 0.14\\,M$_{\\odot}$.  We have put together the Spectral Energy Distribution (SED) from 0.36 to 100\\,$\\mu$m for Mz\\,1, using data from the literature listed in SIMBAD. The plot of $\\lambda\\,F(\\lambda)$ is shown in Fig.\\,\\ref{sed}. Note that the 100\\,$\\mu$m IRAS point is an upper limit. The double peaked distribution is typical for a PN.  \\clearpage  \\begin{figure} \\includegraphics[width=\\columnwidth]{f12.eps} \\caption{The SED for Mz\\,1 between 0.44\\,$\\mu$m and 100\\,$\\mu$m using flux values taken from the literature.\\label{sed}} \\end{figure}  \\clearpage  Integrating our SED yields a luminosity of about 60\\,L$_{\\odot}$ which includes the usual correction proposed by \\cite{M87} of a factor of 1.5. This implies that about 60\\% of the UV radiation from the CS escapes from the nebula. Given that the structure we found is open toward the poles, and is likely clumpy, this is a reasonable value. The fact that we estimate that there is between 0.2 and 2\\,M$_{\\odot}$ of neutral matter around Mz\\,1, some of this may have, in fact, become ionized due to this escaping radiation. Searching for a large, low surface brightness halo around Mz\\,1 may therefore be useful.  Perhaps most important is the fact that {\\it our simple method can be applied to any spatially resolved emission line nebula, yielding the complete 3--D structure and an accurate distance}. The time consuming multiple position long slit observations can be dispensed with since Integral Field Spectrographs are now becoming available at major observatories, making the method not just simple but also efficient. See the detailed explanation of our method in the appendix."
    },
    "0410/astro-ph0410343_arXiv.txt": {
        "abstract": "The final merger of a pair of massive black holes in a galactic nucleus is compelled by gravitational radiation. Gravitational waves from the mergers of black holes of masses $(10^5-10^7)(1+z)^{-1}M_\\odot$ at redshifts of $1-20$ will be readily detectable by the {\\it Laser Interferometer Space Antenna}, but an electromagnetic afterglow would be helpful in pinpointing the source and its redshift. Long before the merger, the binary ``hollows out'' any surrounding gas and shrinks slowly compared to the viscous timescale of a circumbinary disk. The inner gas disk is truncated at the radius where gravitational torque from the binary balances the viscous torque, and accretion onto the black holes is diminished.  Initially, the inner truncation radius is able to follow the shrinking binary inward. But eventually the gravitational radiation timescale becomes shorter than the viscous timescale in the disk, leading to a merged black hole surrounded by a hollow disk of gas. We show that the subsequent viscous evolution of the hollow, radiation pressure-dominated disk will create an $\\sim 10^{43.5}(M/10^6M_\\odot)\\mbox{ergs s}^{-1}$ X-ray source on a timescale $\\sim  7(1+z)(M/10^6M_\\odot)^{1.32}\\,\\mbox{yr}$. This justifies follow-up monitoring of gravitational wave events with next-generation X-ray observatories.  Analysis of the detailed light curve of these afterglows will yield new insights into the subtle physics of accretion onto massive black holes. ",
        "introduction": "\\label{sec:intro}  Evidence is mounting that galactic spheroids contain massive black holes (MBHs) in their nuclei.  When two galaxies merge, their MBHs form a binary in the nucleus of the new galaxy \\citep{Begelman:80}.  The binary interacts with its stellar and gaseous environment and with other MBHs that collect in the same nucleus in multiple mergers.  These interactions torque the binary and extract its orbital angular momentum and energy.  The binary can thus be rendered so compact that gravitational radiation (GR) carries away the remaining orbital energy, inducing coalescence.  For binaries with masses $M=M_1+M_2\\lesssim10^7 M_\\odot$, the gravitational slingshot ejection of stars is sufficient to guarantee coalescence in a Hubble time \\citep{Milosavljevic:03}.  GR emitted by these binaries shortly before and during coalescence will be detected by the space-based gravitational wave detector {\\it Laser Interferometer Space Antenna} ({\\it LISA}).\\footnote{http://lisa.jpl.nasa.gov.} The coalescence of isolated black holes in general relativity is not accompanied by observable electromagnetic (EM) emission.  We here show that circumbinary gas can lead to a delayed EM afterglow.  Of interest to {\\it LISA} are mergers of binaries with masses $M<10^7 M_\\odot$ of arbitrary mass ratio $q\\leq 1$.  There is abundant evidence for dense gas in galactic nuclei.  Geometrically thin molecular disks have been detected in water maser emission in Seyfert galaxies (e.g., \\citealt{Miyoshi:95,Greenhill:03}).  The Galactic nucleus contains a $4\\times10^6M_\\odot$ MBH surrounded by an $\\sim10^4M_\\odot$ molecular gas torus (e.g.,~\\citealt{Jackson:93}).  Massive accretion disks must be present in quasars and narrow-line Seyfert I nuclei to explain their luminosities.  In general, the specific angular momentum of inflowing gas exceeds that of the binary, while the gas temperature is below the virial temperature $GM\\mu m_{\\rm p}/2ak$, where $\\mu$ is the mean molecular weight in units of the proton mass $m_{\\rm p}$, $a$ is the binary's semimajor axis, and $k$ is the Boltzmann constant.  The sub-virial gas settles into a rotationally supported, geometrically thin circumbinary disk.  If the disk is inclined with respect to the binary's orbital plane, the quadrupole component of the binary's gravitational potential causes differential precession and the warping of the disk.  As in the \\citet{Bardeen:75} mechanism, the warp dissipates viscously, resulting in a disk in the binary's orbital plane \\citep{Larwood:97,Ivanov:99}.  The disk is truncated at an inner edge where gravitational torques and viscous stresses balance (e.g.,~\\citealt{Artymowicz:94}). Surface density in the hollowed region is much smaller than in the disk (e.g., \\citealt{Guenther:04}). While material does peel off the inner edge and flow across the hole \\citep{Artymowicz:96}, the accretion rate for moderate mass ratios ($q\\gtrsim 0.01$) is only a fraction ($\\lesssim 10\\%$) of the rate at which the disk would be accreting without being torqued by the binary (\\citealt{Lubow:99}). As the binary's semimajor axis decays owing to stellar or gas dynamical processes (including angular momentum extraction by the circumbinary disk itself), the disk's inner edge spreads inward viscously while maintaining an approximately constant ratio of inner edge radius to semimajor axis, $r_{\\rm edge}/a\\sim2$. In the final stages of in-spiral, however, the time for decay of the semimajor axis is set by GR and decreases rapidly with $a$ (equation \\ref{eq:grtime}). When this becomes shorter than the (viscous) time for the inner edge of the disk to spread inward, the binary continues toward merger while the disk structure remains frozen \\citep{Armitage:02}. This was also noted by \\citet{Liu:03}, who suggested that the formation of an inner hole in the disk preceding MBH coalescence causes the interruption of jet activity in double-double radio galaxies.  Vertical support in the central parts of MBH accretion disks is dominated by radiation pressure, while opacity is dominated by electron scattering.  The \\citet{Shakura:73} $\\alpha$-disks in this regime are thermally and viscously unstable \\citep{Lightman:74} if viscous stresses scale with total pressure. However, they are stable if the stresses scale with the gas pressure alone \\citep{Sakimoto:81}.  Such a scaling can occur \\citep{Turner:04} if the distance photons diffuse per orbit is about equal to the scale of magnetic fields produced by the magnetorotational instability (MRI; \\citealt{Balbus:91}) that facilitates angular momentum transport.  The radiation then partly decouples from the turbulence driven by the MRI \\citep{Turner:03}. These disks are clumpy and porous to the radiation produced within (\\citealt{Begelman:02} and references therein).  We here describe the observable signatures of such circumbinary disks in MBHs.  In \\S~\\ref{sec:model}, we present a scenario for the combined evolution of an MBH binary and its circumbinary accretion disk immediately preceding coalescence.  In \\S~\\ref{sec:surveys}, we discuss observable signatures. ",
        "conclusions": ""
    },
    "0410/astro-ph0410669_arXiv.txt": {
        "abstract": "We have cross-correlated X-ray catalogs derived from archival $Chandra$ ACIS observations with a Sloan Digital Sky Survey (SDSS) Data Release 2 (DR2) galaxy catalog to form a sample of 42 serendipitously X-ray detected galaxies over the redshift interval $0.03 < z < 0.25$.  This pilot study will help fill in the ``redshift gap\" between local X-ray-studied samples of normal galaxies and those in the deepest X-ray surveys. Our chief purpose is to compare optical spectroscopic diagnostics of activity (both star-formation and accretion) with X-ray properties of galaxies.  Our work supports a normalization value of the X-ray-star-formation-rate (X-ray-SFR) correlation consistent with the lower values published in the literature.  The difference is in the allocation of X-ray emission to high-mass X-ray binaries relative to other components such as hot gas, low-mass X-ray binaries, and/or AGN.  We are able to quantify a few pitfalls in the use of lower-resolution, lower signal-to-noise optical spectroscopy to identify X-ray sources (as has necessarily been employed for many X-ray surveys).  Notably, we find a few AGN that likely would have been misidentified as non-AGN sources in higher-redshift studies. However, we do not find any X-ray hard, highly X-ray-luminous galaxies lacking optical spectroscopic diagnostics of AGN activity.  Such sources are members of the ``X-ray Bright, Optically Normal Galaxy\" (XBONG) class of AGN. ",
        "introduction": "}  X-ray emission in galaxies has at a basic level two origins:  gravitational potential energy release from accretion processes and thermal emission from hot gas.  Accretion processes may be divided among three main systems:  where a lower mass ($\\simlt1$~\\msun), and correspondingly longer-lived ($>1$~Gyr), star accretes onto either a neutron star or black hole (low-mass X-ray binary; hereafter LMXB), where a higher mass ($\\simgt5$~\\msun) and thus shorter-lived ($<10^{7}$~yr)  star accretes onto a neutron star or black hole (high-mass X-ray binary; hereafter HMXB), and finally where accretion occurs onto a nuclear supermassive (10$^6$--10$^{8}$~\\msun) black hole (active galactic nucleus; hereafter AGN).   Hot gas is produced through a combination of stellar mass loss and supernovae, and this phase is often the dominant one in the interstellar media (ISM) of galaxies.   Indeed, one of the more notable results in X-ray studies of galaxies was finding the expected interstellar-medium (ISM) component of elliptical galaxies \\citep{Forman79}. The ISM in ellipticals is largely invisible at radio and optical wavelengths and had for a time presented a mass-conservation problem given the expected output from stellar winds. The hot ISM in galaxies with younger stellar populations may be highly extended, as supernova and massive stellar winds give rise to large-scale outflows \\citep[superwinds;][]{Heckman90}. For the moment ignoring AGN, it is thus expected that in early-type galaxies, with older stellar populations, we will find that X-ray emission is dominated by LMXBs and hot ISM.  In late-type galaxies, we expect a mix of emission from hot gas, LMXBs, and HMXBs with starburst galaxies having a great amount of both hot gas and HMXBs.   The hot gas component is typically X-ray soft (dominating at $\\simlt1$~keV) whereas the X-ray binary emission is more X-ray hard (dominating the total X-ray emission at $>2$~keV).  It appears that, as expected, the properties of X-ray binaries are correlated with star formation rate (hereafter SFR) in many nearby ($\\simlt 50$~Mpc) galaxies studied intensively with Chandra \\citep[e.g., the X-ray Luminosity Functions of the binaries appear to obey scaling laws based upon SFR; ][]{Kilgard02,ColbertXLF2003}. The total X-ray emission from binary systems may be written as a linear combination of the young population associated with SFR and the older population associated with stellar mass \\citep[i.e.,$L_{\\rm X}=\\alpha*SFR + \\beta*M$; ][]{David92,ColbertXLF2003}. In galaxies with a higher ratio of SFR to stellar mass, where the LMXB component is expected to be negligible, the total X-ray luminosity may be dominated by HMXBs and it may be possible to use the total X-ray luminosity, neglecting other contributions, as an SFR indicator \\citep[e.g.,][]{Gilfanov2004SFR}.   Indeed, there have been indications that the total X-ray luminosity of a galaxy may serve as a star-formation diagnostic at higher redshift ($z\\simgt0.3$, $d\\simgt1700$~Mpc) \\citep[e.g.,][]{GrimmSFR,BauerXII}.  Of cosmological interest, it also appears that X-ray emission from distant Lyman Break galaxies may be similar to galaxies in the nearby Universe \\citep[e.g.,][]{BrandtLyBreak,Seibert2002,Nandra02,Reddy2004}.  The X-ray band holds much promise as an independent probe of SFR.  For instance, it has been shown that one of the chief uncertainties in the determination of the power-law index of the Initial Mass Function (hereafter IMF) arises mostly from unresolved binary star systems \\citep{Kroupa2001}.  Theoretical models of the production of accreting X-ray binary systems, which are probed through X-ray observations of star-forming galaxies, are sensitively dependent upon properties of the binaries in galaxies \\citep[e.g.][]{Sipior2004}.  It is thus possible that as both the models and the empirical X-ray/SFR relations improve we may develop a probe that is particularly sensitive to the binary populations in galaxies. Interesting conclusions regarding the evolution of the ``binary IMF\" have been discussed in \\cite{NormanXLF}.  Constraining the X-ray/SFR relation requires understanding the X-ray components of galaxies that are not directly associated with current star formation.  There have been fairly different approaches to this problem.   For instance, \\cite{GrimmSFR} select galaxies with high ratios of SFR to stellar mass (hereafter SFR/M$_{*}$) claiming that they can then reject the LMXB component more easily.  \\cite{GrimmSFR} claim that the ${L_{\\rm X}}$-${\\rm SFR}$ relation is linear only at higher SFR ($\\simgt 5$~\\msun~yr$^{-1}$) where high-mass X-ray binaries (HMXBs) dominate the X-ray emission in galaxies. At lower SFR, they claim that the HMXB  ${L_{\\rm X}}$-${\\rm SFR}$ relation becomes non-linear due to small number statistics \\citep{Gilfanov2004stats}. Four other groups have found that the ${L_{\\rm X}}$-${\\rm SFR}$ relation is linear from low to high SFR \\citep{BauerXII,Ranalli03,ColbertXLF2003,Persic2004}. Both \\cite{Persic2004} and \\cite{ColbertXLF2003} include galaxies with lower SFR/M$_{*}$ values and make direct accommodation for the LMXB component. Additionally, \\cite{ColbertXLF2003} uses only the resolved point source population so that the hot gas and nuclear (AGN) component may be rejected. \\cite{Persic2004} use X-ray spectral fits to reject emission from hot gas and AGN.   \\cite{BauerXII}, whose work is at higher redshift, select optical emission-line galaxies which do not show signs for AGN activity. The result is that there is also quite a spread in the {\\it normalization} of the ${L_{\\rm X}}$-${\\rm SFR}$ relation; published values differ by approximately a factor of five.  This difference appears to lie with the way in which one treats the LMXB and AGN X-ray ``contamination\" in the total X-ray luminosity of galaxies.  Understanding this relationship is also quite important for the study of distant galaxies, where AGN contamination may be much more difficult to constrain due to the dilution of AGN optical spectral features by the host galaxy light \\citep[e.g.,][]{Moran2002}. High-$z$ classifications are also more difficult due to the inaccessibility of the H$\\alpha$ emission line, a useful discriminator between true absorption-line dominated (passive) galaxies and those which contain highly obscured AGN \\citep[e.g.,][]{Goto2003}.  There is even an extreme class of objects referred to as ``X-ray Bright, Optically-Normal Galaxies\" (XBONGS) which have optical absorption-line dominated spectra but very high hard X-ray luminosities \\citep{Comastri02}. Any census of AGN activity in the optical waveband must be concerned with the frequency of XBONG-type objects.  Systematic methods for separating X-ray-detected galaxies dominated by AGN emission from those dominated by starburst emission have recently seen a breakthrough with e.g., the Bayesian analysis of \\cite{NormanXLF}.  This work has allowed the first normal/starburst galaxy X-ray luminosity function to be constructed at $z\\approx0.3$ and $z\\approx0.7$ \\citep{NormanXLF}. There is clear (and expected) evolution in this XLF, which has a lognormal shape, similar to FIR luminosity functions.  The evolution appears to be consistent with the expected $(1+z)^{2.7}$ evolution of SFR.  As such high-$z$ normal galaxy XLF studies continue to improve, we may hope to detect subtle evolutionary differences between e.g., the X-ray and FIR in order to constrain e.g., X-ray binary evolution over long timescales.  Such work critically depends upon a thorough understanding of the relative contributions of X-ray emission associated with SFR, AGN or quiescent populations in galaxies.  There remains a problem:  there still is no suitable normal galaxy sample constructed in the nearby ($z\\simlt0.1$) Universe with which to make comparisons with high-redshift X-ray work \\citep[e.g.,][]{HornOBXF,NormanXLF}. The observational expense of studying extremely local ($d<100$~Mpc) galaxies to sufficient depth to probe both diffuse X-ray gas and the binary LF (down to $\\sim10^{38}$~\\lumin, thus reaching neutron-star binaries) is great;  exposures of $~50$~ks are typically required. Currently there are no sufficiently large ($>30$ galaxies) samples complete in any particular observational property (optical magnitude and/or morphology, starburst properties like the ratio of SFR/mass, etc.). Also, while the range of SFR sampled in X-ray-SFR studies to date has been quite large, it has in some cases been heavily biased towards higher star-formation rates (SFR).  The studies aforementioned have typical average SFRs of $\\approx10$~M$_\\odot$~yr$^{-1}$, going up to $\\approx100$~M$_\\odot$~yr$^{-1}$.  This makes probing the possible transition to non-linear ${L_{\\rm X}}$-${\\rm SFR}$ quite difficult. Also, there have been heterogeneous methods used to measure stellar mass and SFR.  For instance, \\cite{GrimmSFR} were forced to use approximately 10 separate FIR studies of nearby galaxies in order to derive SFRs, and for the same galaxies the FIR flux estimate often differed by factors of several.   Reliable UV, H$\\alpha$, and radio SFR estimates are not always available for the same galaxy.  This is expected to change dramatically with well-observed nearby-Universe samples such as the Spitzer Infrared Nearby Galaxies Survey \\citep{KennicuttSINGS}. Finally, there is some concern about ``distance-distance\" biases introduced by the large range in distances sampled by some of these studies.  For instance, many studies place high-$z$ galaxies ($z>0.8$) from the $Chandra$ Deep Fields on the same correlation fits as galaxies in the nearby Universe ($d<50$~Mpc).  However, strong bias is only expected when the range in flux is quite small whereas the distance range is large;  for most studies this is not a serious problem.  It is thus critical to observe galaxies with moderate SFR values ($\\sim1$--10~M$_\\odot$~yr$^{-1}$) as well as galaxies with no evidence for current star formation.   Since the observational challenge at higher redshift is separating the type-2 AGN from the normal/star-forming galaxies, low-to-moderate luminosity type 2 AGN should also be included.  The sample should have {\\it uniformly} measured SFRs and stellar masses and some highly accurate diagnostic of accretion activity.  With these requirements in mind, we have turned to the Sloan Digital Sky Survey (SDSS) Data Release 2 (DR2) catalog and the $Chandra$ X-ray Observatory archive to search for a sample of X-ray-detected galaxies near $z\\approx0.1$.   DR2 is the second major survey-quality data release of the SDSS \\citep[April 2004; ][]{Abaz2003}.  At $z\\simgt0.05$ the space density of galaxies with $R\\simlt18$ reaches $\\approx100$~deg$^{-2}$, so that one would expect $\\approx7$ galaxies in each $Chandra$ ACIS-I field of view ($\\approx16^{\\prime}\\times16^{\\prime}$).  The SDSS is covering $\\approx10$,000 square degrees to this optical depth.  The typical SDSS galaxy redshifts ($z\\approx0.1$) are local enough to provide a calibration sample for the high-redshift studies of galaxies.  We are working towards providing the missing calibration of the X-ray properties of galaxies at $z\\approx0.1$ by studying a uniform sample of SDSS galaxies with serendipitous $Chandra$ detections.  The major advantage of the SDSS is its uniformity; all galaxies have been observed with the same telescope and spectrograph; there is no ambiguity about optical filters or other calibration issues.  The spectroscopic coverage of galaxies in the SDSS DR2 over its $\\approx2$,600 square degree area is complete down to $R\\approx17.8$ \\citep{Stoughton02}. $Chandra$ has the sensitivity to reach galaxies with $\\log{({{f_{\\rm X}}\\over{f_{\\rm R}}})} \\approx -2$ (typical of e.g., luminous dusty starbursts, Alexander et al. 2002) in the 0.5--2~keV band at $r\\approx18$ with 40~ks of ACIS exposure.\\footnote{The $r$ magnitude here is the SDSS $r$-band filter.} Thus the depths of coverage of many $Chandra$ ACIS observations should be sufficient to place significant constraints on the galaxies with SDSS spectra. \\nocite{davoXI}  Our work on individually X-ray detected SDSS galaxies is quite complementary to existing work. There have been some initial promising results on spiral galaxies from statistical analysis of galaxies (not individually detected) using a combination of shallow ($\\simlt 5$~ks) $XMM-Newton$ observations and the 2dF redshift survey \\citep{Georgakakis03}. \\cite{Georgakakis04needles} have also carried out a focused program to identify optically bright, X-ray faint objects which are typically quiescent galaxies \\citep[e.g.][]{HornOBXF}, reporting on constraints on X-ray galaxy number counts at bright X-ray fluxes using SDSS and $XMM$.   Our study only selects against Type-I AGN and unresolved sources, we make no exclusion of sources with higher X-ray-to-optical flux ratios (which may include some dusty starburst galaxies and higher-mass elliptical galaxies). Our focus on {\\it optically}  selected, individually X-ray detected galaxies also complements other serendipitous $Chandra$ programs which are driven by X-ray selection and naturally favor AGN \\citep[e.g., the $Chandra$ Multiwavelength Project, ChaMP; ][]{Green03}.  The purpose of this paper is to compare optical activity diagnostics such as emission-line ratios with the X-ray properties of galaxies.  The ultimate goal of our studies will be to assemble a complete field galaxy sample; this paper is a pilot study which works towards that ultimate goal (in fact, although we avoid cluster centers, a significant fraction of the galaxies are in cluster environments).  Our chief emphasis is on the X-ray/SFR relation, but we also explore the X-ray properties of quiescent galaxies and AGN.  A future paper will cover the normal galaxy SDSS XLF.  The cosmology assumed in this paper is $H_0=70$~km~s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.30$, $\\Omega_{\\Lambda}=0.70$. ",
        "conclusions": "We have conducted a pilot study of 42 SDSS DR2 galaxies serendipitously detected in the X-ray band.  The high signal-to-noise SDSS spectra afford detailed activity diagnostics that are often missing in the deepest surveys.  Among X-ray detected galaxies which show no signs for AGN activity (15 absorption-dominated galaxies and 7 star-forming galaxies), the X-ray emission is naturally explained by a combination of LMXBs, HMXBs, and hot ISM.  Thus, we do not find evidence for AGN that are missed in the SDSS spectroscopic analysis.  Our work thus supports that of others \\citep[e.g.,][]{Moran2002} that fairly simple observational effects may make X-ray detected AGN appear as XBONGS at higher-redshift. Among the 19 SDSS galaxies which have AGN signatures in their optical spectra, the $L_{\\rm X}$-$L_{[OIII]}$ ratios are found to be coarsely consistent with Seyfert 2's in the nearby Universe.  We do find four instances where galaxies might have been misclassified in lower signal-to-noise optical spectra.  A NLS1's line widths clearly reveal the presence of an AGN while its emission-line ratios are similar to that of a starburst galaxy.   Three other galaxies have pure absorption-dominated optical spectra blueward of H$\\alpha$ but exhibit H$\\alpha$ and [NII] emission typical of AGN.  All three of these galaxies would not have been identified as AGN above $z\\approx0.5$.  The detection of X-ray emission from a post-starburst galaxy indicated that the burst mass fraction calculated from the optical spectroscopic models may be in agreement with the expected X-ray emission from a starburst.  However, the likely evolved nature of the post-starburst (the age of which is poorly constrained) indicates that the X-ray luminosity is quite excessive given the magnitude of the starburst. This one object is not sufficient to test the models but is strongly suggestive that a large sample of post-starburst galaxies (perhaps with stronger constraints on the burst mass fractions) may have a strong impact on the modeling of accreting binary systems.  It may also be an excellent way to search for young ULX sources as this galaxy appears to harbor a variable X-ray source.  We are able to rule out that this post-starburst galaxy is an AGN to very strong limits from optical spectroscopy.  We have made an effort to compare the various $L_{\\rm X}$-SFR relationships in the literature. We find that differences among e.g., bolometric corrections and IMF account for up to a 50\\% difference in the SFR relationships.    Any difference greater than this level must then be attributed to e.g., a different treatment of the contributions of AGN vs. star-formation. Our analysis supports the lower normalization of the $L_{\\rm X}$-SFR relation after proper account of AGN activity and contributions from evolved stellar populations are made."
    },
    "0410/astro-ph0410033_arXiv.txt": {
        "abstract": "{ We present column density measurements of the HD molecule in the interstellar gas toward 17 Galactic stars.  The values for the seven most heavily reddened sightlines, with $E(B-V)$ = 0.38--0.72, are derived from observations with the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}).  The other ten values are from a reanalysis of spectra obtained with {\\it Copernicus}.  In all cases, high-resolution ground-based observations of \\ion{K}{i} and/or the CH molecule were used to constrain the gas velocity structure and to correct for saturation effects.  Comparisons of the column densities HD, CH, CN, and \\ion{K}{i} in these 17 sightlines indicate that HD is most tightly correlated with CH.  Stringent lower limits to the interstellar D/H ratio, derived from the HD/2H$_2$ ratio, range from $3.7 \\times 10^{-7}$ to $4.3 \\times 10^{-6}$.  Our results also suggest that the HD/H$_2$ ratio increases with the molecular fraction $f$(H$_2$) and that the interstellar D/H ratio might be obtained from HD by probing clouds with $f$(H$_2$) $\\sim 1$.  Finally, we note an apparent relationship between the molecular fractions of hydrogen and deuterium. ",
        "introduction": "\\begin{table*}[t] \\caption[]{Target List.} \\label{targets} \\centering \\begin{tabular}{llrcccccl} \\hline \\hline Star & $\\alpha$ (J~2000) & $\\delta$ (J~2000) & $V$ (mag) & $E_{(B-V)}$$^{\\mathrm{a}}$ & Distance (pc) & Ref & $A_{V}$ (mag)$^{\\mathrm{a}}$ & Type \\\\ \\hline HD\\,27778  & 04 23 59.79 &   +24 18 03.6 & 6.33 & 0.38&220 & 1& 1.01& B3V \\\\ HD\\,73882  & 08 39 09.53 & $-$40 25 09.3 & 7.27 & 0.72&925 & 2& 2.28&O9III \\\\ HD\\,110432 & 12 42 50.27 & $-$63 03 31.0 & 5.32 & 0.40&430 & 3& 1.32&B2IVpe \\\\ HD\\,185418 & 19 38 27.48 &   +17 15 26.1 & 7.52 & 0.51&790 & 4& 2.03 &B0.5V \\\\ HD\\,192639 & 20 14 30.43 &   +37 21 13.8 & 7.11 & 0.66&1800& 5& 1.87&O8e \\\\ HD\\,206267 & 21 38 57.62 &   +57 29 20.5 & 5.62 & 0.52&615 & 6& 1.37&O6e \\\\ HD\\,207538 & 21 47 39.79 &   +59 42 01.3 & 7.30 & 0.64&615 & 6& 1.43&B0V \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[ References. --- ]{1 = Simbad; 2 = \\citet{2000ApJ...538L..69F}; 3 = \\citet{1989ApJ...340..273V}; 4 = \\citet{2003ApJ...596..350S}; 5 =\\citet{1991A&A...252..781S}; 6 = \\citet{1999AJ....117..354D}.} \\item[$^{\\mathrm{a}}$]{ Reddening and visible extinction parameters from \\citet{2002ApJ...577..221R}.} \\end{list} \\end{table*}  It is believed that deuterium was produced in significant amounts only during the primordial Big Bang Nucleosynthesis.  Since then, deuterium has been steadily destroyed in stellar interiors.  Thus, its abundance relative to H, noted D/H, is a key measurement for studies of both cosmology and Galactic chemical evolution \\citep{2000A&A...360...15V,2004ApJ...600..544C}.  D/H in the Interstellar Medium (ISM) is characteristic of the present-day Galactic deuterium abundance.  Prior to the {\\it FUSE} mission, (D/H)$_{\\rm ISM}$ showed some dispersion \\citep[e.g.][ for a review]{1999NewA....4..231L}, likely resulting from poorly understood physical processes like astration, inefficient mixing, depletion onto grains, and perhaps some unidentified systematic errors.  The final resolution of these issues will have strong implications for our understanding of the physics of the ISM, the chemical evolution of the Galaxy, and the baryonic density of the Universe inferred from primordial D/H.  An accurate determination of (D/H)$_{\\rm ISM}$ is one of the main goals of the {\\it FUSE} mission.  So far, most of these measurements have been performed within the Local Interstellar Medium (LISM). \\citet{2002ApJS..140....3M}  reviewed those and reported (D/H)$_{\\rm LB}$ = (1.52 $\\pm$ 0.08) $\\times 10^{-5}$ within the Local Bubble (LB).  However, this value may not be representative of the true Galactic (D/H)$_{\\rm ISM}$.  \\citet{2003ApJ...599..297H} used a D/O survey to constrain the D/H variations and, assuming a constant O/H ratio, obtained a D/H ratio below $1.0 \\times 10^{-5}$ outside the LB.  Indeed, recent measurements suggest a ``canonical'' (D/H)$_{\\rm ISM}$, if it exists, likely 2 to 3 times lower \\citep{1999ApJ...520..182J, 2000ApJ...545..277S,2003ApJ...586.1094H}. Unfortunately, extended direct investigations over long Galactic disk sight lines are difficult, due to saturation and blends with neighboring \\ion{H}{i} Lyman lines.  The HD/H$_2$ ratio is an interesting alternative for D/H investigations along long sightlines.  In the diffuse ISM, the formation of HD occurs via the ion-neutral reaction: \\begin{displaymath} {\\rm H}_2 + {\\rm D}^+ \\longrightarrow {\\rm HD} + {\\rm H}^+ \\end{displaymath} while its destruction is due to photodissociation.  Because of the lower abundance of deuterium compared to hydrogen, self-shielding of HD becomes significant only at higher extinction than for H$_2$. Therefore, the transition between atomic deuterium and HD takes place deeper in a cloud than the transition between atomic and molecular hydrogen.  Whenever all the deuterium is in molecular form, the ratio $N$(HD)/2$N$(H$_2$) should be equal to the elemental abundance ratio D/H \\citep[see also][]{1973ApJ...182L..73W}.  Less than a dozen observations of deuterated molecular hydrogen have been carried out to date, from the first FUV detections with {\\it Copernicus} thirty years ago \\citep{1973ApJ...181L.116S,1973ApJ...184L.101B, 1973ApJ...182L..73W} to infrared observations with {\\it ISO} \\citep{1999ApJ...515L..29W,1999A&A...346..267B}.  Recently, the higher sensitivity of the {\\it FUSE} satellite has allowed detection of HD in extra-galactic sources \\citep{2001A&A...379...82B,2004A&A...422..483A}. In addition, more heavily reddened clouds which previously could not be investigated in the FUV are now within reach, as shown by \\citet{2000ApJ...538L..69F} toward HD\\,73882 ($A_V \\approx 2.3$ mag).   In this work, we make use of FUV data obtained with {\\it FUSE} towards seven early-type Galactic stars (see Table~\\ref{targets}).  Each sightline is ``translucent'', i.e., showing significant extinction \\citep[$1$ mag $< A_V < 5$ mag; see][]{2002ApJ...577..221R}. Because of the high sensitivity of {\\it FUSE} \\citep[10\\,000 times more sensitive than {\\it Copernicus};][]{2000ApJ...538L...1M}, the data have good $S/N$ ratios, which allow accurate measurements of equivalent widths.  However, at the large column densities needed to have most of the deuterium in molecular form, most of the HD lines available will be strongly saturated.  Prior knowledge of the sightline velocity structure is therefore crucial for the analysis. We made use of available very high resolution optical data for CH -- known to be a good tracer of H$_2$ \\citep{1984A&A...130...62D,1998ApJ...504..290M} -- in order to constrain the gas velocity structure toward our target stars.  In the next section we describe the criteria used to select the targets and note some interesting aspects of each of the seven lines of sight. In sections 3 and 4, we present the FUV observational data and our methods of reduction and analysis.  In section 5, we discuss the inferred deuterium abundances. ",
        "conclusions": "\\begin{table} \\caption{{\\it FUSE} HD Column Densities.} \\label{HD sum} \\centering \\begin{tabular}{lccc} \\hline \\hline Star & \\multicolumn{3}{c}{$N({\\rm HD})$$^{\\mathrm{a}}$} \\\\ \\cline{2-4} {} & {FIT} & {COG} & {Mean} \\\\ \\hline HD\\,27778   & $3.0^{+2.9}_{-1.7}$ &$3.4^{+3.4}_{-1.8}$ &$3.2^{+3.2}_{-1.7}$ \\\\ HD\\,73882   & $7.4^{+4.9}_{-5.2}$ &$4.2^{+2.6}_{-1.6}$ &$5.8^{+3.7}_{-3.4}$ \\\\ HD\\,110432  & $1.9^{+0.5}_{-0.4}$ &$2.0^{+1.0}_{-0.7}$ &$1.9^{+0.7}_{-0.6}$ \\\\ HD\\,185418  & $4.0^{+1.5}_{-0.8}$ &$4.6^{+2.2}_{-1.4}$ &$4.3^{+1.9}_{-1.1}$ \\\\ HD\\,192639  & $1.2^{+0.6}_{-0.5}$ &$1.7^{+1.2}_{-0.8}$ &$1.5^{+0.9}_{-0.7}$ \\\\ HD\\,206267  & $1.3^{+1.0}_{-0.6}$ &$2.9^{+2.0}_{-1.4}$ &$2.1^{+1.5}_{-1.0}$ \\\\ HD\\,207538  & $3.5^{+3.9}_{-1.8}$ &$6.6^{+6.3}_{-3.1}$ &$5.0^{+5.1}_{-2.4}$ \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$]{All column densities are times 10$^{15}$.} \\end{list} \\end{table}  The values of $N$(HD) obtained for the 7 {\\it FUSE} lines of sight are summarized in Table~\\ref{HD sum}.  There is good agreement between the column densities determined using both the COG and the PF techniques.  In these particular cases, similar values are obtained by fixing only the component velocity distribution (PF) and by fixing in addition the relative strengths and the broadening of the CH components (COG). This agreement gives an indication of the reliability of our results. For the discussion below, we adopt the unweighted mean of the column densities determined by the two techniques.  \\begin{table*} \\caption{{\\it Copernicus} HD Equivalent Widths and Column Densities.} \\label{tab_EWCop} \\centering \\begin{tabular}{lcccccccccccc} \\hline \\hline \\multicolumn{1}{l}{Star}& \\multicolumn{1}{c}{$N$(HD) $^{\\mathrm{a}}$}& \\multicolumn{2}{c}{Inferred EqWs $^{\\mathrm{a}}$}& \\multicolumn{2}{c}{Literature EqWs}& \\multicolumn{1}{c}{$N$(HD)}& \\multicolumn{1}{c}{Ref}& \\multicolumn{2}{c}{New EqWs}& \\multicolumn{1}{c}{New $N$(HD)}& \\multicolumn{1}{c}{$\\Delta N$(HD)}\\\\ \\multicolumn{2}{c}{ }& \\multicolumn{1}{c}{1054}& \\multicolumn{1}{c}{1066}& \\multicolumn{1}{c}{1054}& \\multicolumn{1}{c}{1066}& \\multicolumn{1}{c}{ }& \\multicolumn{1}{c}{ }& \\multicolumn{1}{c}{1054}& \\multicolumn{1}{c}{1066}& \\multicolumn{1}{c}{ }& \\multicolumn{1}{c}{(dex)}\\\\ \\hline HD\\,21278& 14.06 & 18.5&     &            &            &            & & &            &   14.2(0.1)& $+0.15$\\\\ $o$ Per &       &     &     &            &   31.0(2.0)&   15.77    &1& &            &   15.5(0.3)& $-0.25$\\\\ $\\zeta$ Per & 14.30 & 32.2&     &   41.4(4.9)&            &   15.60    &2& 27.7(2.6)&            &   15.6(0.2)& $+1.30$\\\\ $\\epsilon$ Per & 13.57 &  6.0&     &            &            &            & & 6.6(0.4)&            &   13.8(0.1)& $+0.25$\\\\ $\\xi$ Per  & 14.15 & 22.8& 16.4&            &            &            & & 22.9(0.6)&   17.5(1.1)&   14.4(0.1)& $+0.25$\\\\ $\\alpha$ Cam & 14.49 & 49.9& 35.8&            &            &            & & 52.0(2.0)&   48.4(2.3)&   14.9(0.1)& $+0.40$\\\\ 139 Tau & 13.84 & 11.2&     &            &            &            & & 8.5(1.0)&            &   13.9(0.1)& $+0.05$\\\\ $\\zeta$ Oph & 14.23 & 27.4& 19.7&   20.6(1.6)&   15.9(0.8)&   14.26    &3& 18.8(1.0)&   14.6(0.6)&   14.8(0.3)& $+0.55$\\\\ 59 Cyg  & 13.86 & 11.7&     &            &            &            & & 8.8(1.1)&            &   13.9(0.1)& $+0.05$\\\\ 10 Lac  & 14.41 & 41.5& 29.7&            &            &            & & 5.5(0.7)&  [3.3(1.0)]&   13.5(0.2)& $-0.90$\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[ References. --- ]{1 = \\citet{1975ApJ...201L..21S,1976ApJ...204..759S}; 2 = \\citet{1977ApJ...216..724S}; 3 = \\citet{1979ApJ...227..483W}} \\item[ Note. --- ]{Errors on new equivalent widths (EqW) are 1-sigma.} \\item[$^{\\mathrm{a}}$]{ \\citet{1973ApJ...181L.116S}; equivalent widths inferred from listed $N$(HD) and $f$-values of \\citet{1969AD......1..289A}, assuming optically thin lines (see Sect.~\\ref{sec:CHHD}).} \\end{list} \\end{table*}   \\subsection{Relation between CH and HD \\label{sec:CHHD}}   The good correlation found between the column densities of CH and H$_2$ in diffuse interstellar clouds is consistent with the production of most of the CH via a network of gas phase reactions in which H$_2$ is a major component \\citep{1984A&A...130...62D}.  This correlation recently was extended to more heavily reddened sightlines observed with {\\it FUSE} \\citep{2002ApJ...577..221R}.  The relatively tight relationship between these two species and the inferred chemical link both suggest that CH can be used as a tracer of H$_2$.  Because HD is also produced through gas phase reactions with H$_2$ and because both CH and HD are subject to destruction through photoprocesses, one could assume that CH would be a good tracer for HD as well.  To test this assumption, we plotted the column densities of H$_2$ versus HD, and H$_2$ and HD versus CH; this for all the Galactic lines of sight having estimates for $N$(HD).  The sample includes ten sightlines observed with {\\it Copernicus} \\citep{1973ApJ...181L.116S,1975ApJ...201L..21S,1976ApJ...204..759S, 1977ApJ...216..724S,1979ApJ...227..483W} (see below) and the seven {\\it FUSE} sightlines analyzed in the present work.  Because the $N$(HD) originally given by \\citet{1973ApJ...181L.116S} were estimated assuming the HD lines (at 1054 and/or 1066 \\AA) to be optically thin, those values should be taken as lower limits to the true $N$(HD).  We therefore have re-examined all the {\\it Copernicus} data and have derived new estimates for $N$(HD), using a method similar to that used for analyzing the {\\it FUSE} spectra.  The {\\it Copernicus} U1 scans encompassing the HD lines at 1054 and/or 1066 \\AA\\ were obtained from the MAST archive.  Background levels were estimated from scans of nearby saturated H$_2$ lines obtained in the same observing programs.  The equivalent widths measured from the normalized spectra (Table~\\ref{tab_EWCop}) are generally in good agreement with those inferred from the $N$(HD) listed by \\citet{1973ApJ...181L.116S} and with the values listed by \\citet{1979ApJ...227..483W}.  The most notable exception to that agreement is 10 Lac, where we suspect that the absorption near the HD line at 1054 \\AA\\ is due mostly to stellar \\ion{Fe}{iii} and/or \\ion{Cr}{iii} \\citep{1977ApJS...35...37R,1985ApJS...58..265R}; there is no strong absorption at the expected position of the HD line at 1066 \\AA.  [Spitzer et al.'s value for $N$(HD) toward 10 Lac was surprisingly high, given the modest $E_{(B-V)}$ and molecular fraction.]  Component structures derived from high-resolution spectra of CH, \\ion{K}{i}, and/or \\ion{Na}{i} were used to model the profiles and/or equivalent widths of the HD lines ($\\lambda$1054, $\\lambda$1066, and others as available).  Our ``new'' {\\it Copernicus} $N$(HD) are listed in the last column of Table~\\ref{tab_EWCop} and plotted with open squares in Fig.~\\ref{fig:cor}. In most cases, the new values are higher than those listed by \\citet{1973ApJ...181L.116S} by factors of order 2; the value for $\\zeta$ Oph is about 4 times higher.  On the other hand, the new value for 10 Lac is about a factor of 8 smaller, and the value for $o$ Per is about a factor of 2 smaller than that given by \\citet{1975ApJ...201L..21S,1976ApJ...204..759S}.  We note that these new values could still underestimate the true $N$(HD) if HD is actually concentrated in fewer components than the tracer used to model the HD equivalent widths.   \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{Correlations.eps}} \\caption{ ($top$) $N$(HD) vs $N$(H$_2$) : The correlation factor close to 1 comfirmes a good correlation between the two species, as predicted by chemical models \\citep{1984A&A...130...62D}. The slope is slightly greater than 1.0. --- ($middle$) $N$(H$_2$) and $N$(HD) vs $N$(CH) : Tight correlation between $N$(H$_2$) and $N$(CH) was already observed in wide survey \\citep{2002ApJ...577..221R}, and appears in the left panel with a slope close to 1. $N$(HD) vs $N$(CH) also shows a good relationship, but with a slope near 2.0. --- ($bottom$) $N$(HD) vs $N$(CN) and $N$(\\ion{K}{i}) : None of the species exhibit a tight relationship. --- Filled diamonds have $N$(HD) derived from {\\it FUSE} spectra; open squares have $N$(HD) derived from reanalysis of {\\it Copernicus} data. All the correlation factors and slopes are summarized in Table~\\ref{correlation}.  } \\label{fig:cor} \\end{figure}  Table~\\ref{correlation} summarizes the correlations between several different species.  As shown on Fig.~\\ref{fig:cor}, it suggests a good correlation between HD and H$_2$, and between H$_2$ and CH. Thus it is no surprise that it indicates a correlation between HD and CH column densities, but with a slope somewhat steeper than that for $N$(H$_2$) vs $N$(CH).  There are few points with $N$(HD) $\\la$ 10$^{15}$ cm$^{-2}$, however, so it would be very useful to obtain more measurements of HD (and CH) at those lower column densities.  We also investigated whether CN and \\ion{K}{i} might be useful as tracers for HD, as all three species are expected to be concentrated in the denser parts of diffuse clouds.  In addition, \\ion{K}{i} is generally more readily detected than CH and typically reveals more details of the component structure \\citep[e.g.,][]{2001ApJS..133..345W}.  Bottom panels of Fig.~\\ref{fig:cor} show HD vs CN and \\ion{K}{i}.  While CN and \\ion{K}{i} show some degree of correlation with HD, the relationships are not as tight as that with CH.  \\subsection{The D/H ratio}   \\begin{table} \\caption{Correlation Plots.} \\label{correlation} \\centering \\begin{tabular}{lrrc} \\hline \\hline \\multicolumn{1}{l}{Quantities}& \\multicolumn{1}{c}{$r$ $^{\\mathrm{a}}$}& \\multicolumn{1}{c}{$N$ $^{\\mathrm{b}}$}& \\multicolumn{1}{c}{Slope $^{\\mathrm{c}}$} \\\\ \\hline HD vs. H$_2$           & 0.923 & 16 & 1.29$\\pm$0.14  \\\\ H$_2$ vs. CH           & 0.911 & 12 & 1.18$\\pm$0.19  \\\\ HD vs. CH              & 0.807 & 13 & 1.94$\\pm$0.44  \\\\ HD vs. CN              & 0.732 &  9 & 0.66$\\pm$0.15  \\\\ HD vs. K~I             & 0.745 & 14 & 0.86$\\pm$0.20  \\\\ & \\\\ HD/2H$_2$ vs. H$_2$    & 0.291 & 16 & 0.05$\\pm$0.13  \\\\ & \\\\ $f$(HD) vs. $f$(H$_2$) -- (D/H)$_{\\rm ISM}$=0.74e-5 & 0.725 & 16 & 0.40$\\pm$.10 \\\\ $f$(HD) vs. $f$(H$_2$) -- (D/H)$_{\\rm ISM}$=1.52e-5 & 0.725 & 16 & 0.19$\\pm$.05 \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$]{ Linear correlation coefficient.} \\item[$^{\\mathrm{b}}$]{ Number of points (detections only).} \\item[$^{\\mathrm{c}}$]{ Slope and uncertainty of weighted least-squares linear fit, considering uncertainties in both quantities.} \\end{list} \\end{table}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{HDsurH2.eps}} \\caption{  $N$(HD)/2$N$(H$_2$) vs $N$(H$_2$) and $f$(H$_2$). The upper dotted line gives the LISM D/H value ($1.5 \\times 10^{-5}$) found by \\citet{2002ApJS..140....3M}; the lower dotted line gives the D/H value found toward $\\delta$ Ori ($0.74 \\times 10^{-5}$) by \\citet{1999ApJ...520..182J}, representative of the lower D/H values found in several other studies.  Except for two sigthlines, we note an increasing trend with the molecular fraction which does not occur as a function of the molecular column density. It could be an indication that $N$(HD)/2$N$(H$_2$) did not reach the D/H value in the sightlines, showing therefore the importance of examining sightlines or clouds with higher $f$(H$_2$). Filled diamonds have $N$(HD) derived from {\\it FUSE} spectra; open squares have $N$(HD) derived from reanalysis of {\\it Copernicus} data. \\label{HDsurH2}} \\end{figure}   \\begin{table*} \\caption{Column Densities.} \\label{dens sum} \\centering \\begin{tabular}{lcccccccrc} \\hline \\hline {Star} & { $E_{(B-V)}$} & { $f({\\rm H_2})$ $^{\\mathrm{a}}$}& {$N($\\ion{H}{i})} & {$N({\\rm H}_2)$} & \\multicolumn{1}{c}{$N({\\rm HD})$} & {Ref} & {$N({\\rm HD})/2N({\\rm H}_2)$} & {$N({\\rm CH})$} & {ref} \\\\ {} & {} & {}  & {[cm$^{-2}$]} & {[cm$^{-2}$]} & {[cm$^{-2}$]} & {}& {}& {[cm$^{-2}$]}\\\\ \\hline HD\\,27778    & 0.38  & 0.57  & $9.5^{+ 9.5}_{-  4.8}$ (20)&  $6.2^{+  0.9}_{-  0.8}$ (20)& $3.2^{+ 3.2}_{-  1.7}$ (15)&1&$2.6^{+  2.6}_{-  1.4}$ (-6)& $3.0\\pm0.3$ (13)& 2\\\\ HD\\,73882      & 0.72  & 0.67  & $1.3^{+ 0.5}_{-  0.4}$ (21)&  $1.3^{+  0.3}_{-  0.2}$ (21)& $5.8^{+ 3.7}_{-  3.4}$ (15)&1&$2.2^{+  1.6}_{-  1.3}$ (-6)& $3.5\\pm0.4$ (13)& 2 \\\\ HD\\,110432     & 0.40  & 0.55  & $7.1^{+ 2.9}_{-  2.1}$ (20)&  $4.4^{+  0.4}_{-  0.4}$ (20)& $1.9^{+ 0.7}_{-  0.6}$ (15)&1&$2.2^{+  0.9}_{-  0.7}$ (-6)& $1.5\\pm0.3$ (13)& 3\\\\ HD\\,185418     & 0.51  & 0.47  & $1.3^{+ 0.5}_{-  0.4}$ (21)&  $5.8^{+  0.7}_{-  0.6}$ (20)& $4.3^{+ 1.9}_{-  1.1}$ (15)&1&$3.7^{+  1.7}_{-  1.0}$ (-6)& $1.3\\pm0.3$ (13)& 2\\\\ HD\\,192639     & 0.66  & 0.32  & $2.1^{+ 0.7}_{-  0.5}$ (21)&  $4.9^{+  0.6}_{-  0.5}$ (20)& $1.5^{+ 0.9}_{-  0.7}$ (15)&1&$1.5^{+  1.0}_{-  0.7}$ (-6)& $2.8\\pm0.5$ (13)& 2\\\\ HD\\,206267     & 0.52  & 0.42  & $2.0^{+ 0.8}_{-  0.6}$ (21)&  $7.2^{+  0.7}_{-  0.6}$ (20)& $2.1^{+ 1.5}_{-  1.0}$ (15)&1&$1.4^{+  1.0}_{-  0.7}$ (-6)& $3.0\\pm0.2$ (13)& 2\\\\ HD\\,207538     & 0.64  & 0.42  & $2.2^{+ 0.7}_{-  0.5}$ (21)&  $8.1^{+  0.8}_{-  0.7}$ (20)& $5.0^{+ 5.1}_{-  2.4}$ (15)&1&$3.1^{+  3.2}_{-  1.5}$ (-6)& $4.3\\pm0.2$ (13)& 2\\\\ & \\\\ HD 21278       & 0.10 & 0.10 $^{\\mathrm{b}}$ & 5.5 (20) $^{\\mathrm{b}}$ & 3.0$^{+2.1}_{-1.2}$ (19) & 1.5$^{+0.1}_{-0.1}$ (14)    &4& 2.5$^{+1.7}_{-1.0}$ (-6)    & \\nodata     &\\nodata \\\\ $o$ Per        & 0.30 & 0.50 & 8.0$\\pm$2.4 (20) & 4.0$^{+1.6}_{-1.2}$ (20) & 3.0$^{+3.0}_{-1.5}$ (15)    &4& 3.8$^{+4.4}_{-2.0}$ (-6)    & $2.1\\pm0.2$ (13)    & 5 \\\\ $\\zeta$ Per    & 0.33 & 0.59 & 6.4$\\pm$0.6 (20) & 4.7$^{+2.4}_{-1.6}$ (20) & 4$^{+4}_{-2}$ (15)          &4& 4.3$^{+5.3}_{-2.4}$ (-6)    & $2.2\\pm0.2$ (13)    & 5 \\\\ $\\epsilon$ Per & 0.09 & 0.21 & 2.5$\\pm$0.5 (20) & 3.3$^{+2.7}_{-1.5}$ (19) & 6$^{+2}_{-1}$ (13)          &4& 9.1$^{+8.7}_{-4.9}$ (-7)    & $<$ 2.0 (12)& 6 \\\\ $\\xi$  Per     & 0.33 & 0.35 & 1.3$\\pm$0.3 (21) & 3.4$^{+1.4}_{-1.0}$ (20) & 2.5$^{+0.5}_{-0.5}$ (14)    &4& 3.7$^{+1.8}_{-1.2}$ (-7)    & $1.3\\pm0.3$ (13)    & 5 \\\\ $\\alpha$ Cam   & 0.32 & 0.35 & 8.0$\\pm$1.6 (20) & 2.2$^{+0.9}_{-0.6}$ (20) & 8$^{+2}_{-2}$ (14)          &4& 1.8$^{+0.9}_{-0.7}$ (-6)    & $6.9\\pm1.6$ (12)    & 5 \\\\ 139 Tau        & 0.15 & 0.12 & 8.0$\\pm$1.6 (20) & 5.4$^{+3.1}_{-2.0}$ (19) & 8$^{+2}_{-2}$ (13)          &4& 7.4$^{+5.0}_{-3.1}$ (-7)    & \\nodata     &\\nodata \\\\ $\\zeta$ Oph    & 0.32 & 0.63 & 5.2$\\pm$0.3 (20) & 4.4$^{+0.9}_{-0.8}$ (20) & 6$^{+6}_{-3}$ (14)          &4& 6.8$^{+7.2}_{-3.5}$ (-7)    & $2.5\\pm0.2$ (13)    & 5 \\\\ 59 Cyg         & 0.18 & 0.19 & 1.8$\\pm$0.4 (20) & 2.0$^{+1.0}_{-0.7}$ (19) & 8$^{+2}_{-2}$ (13)          &4& 2.0$^{+1.2}_{-0.8}$ (-6)    & \\nodata     &\\nodata \\\\ 10 Lac         & 0.11 & 0.06 & 5.0$\\pm$1.5 (20) & 1.7$^{+0.5}_{-0.4}$ (19) & 3$^{+2}_{-1}$ (13)          &4& 8.8$^{+6.9}_{-3.4}$ (-7)    & $2.5\\pm0.8$ (12)    & 7 \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[ References. --- ]{1 = this paper ({\\it FUSE} data); 2 = Welty et al., in prep.; 3 = \\citet{1995MNRAS.277..458C}; 4 = this paper ({\\it Copernicus} data); 5 = \\citet{2003ApJ...584..339T}, 6 = \\citet{1994ApJ...424..772F}; 7 = Thorburn, priv. comm.} \\item[ Note. --- ]{(nn) stands for 10$^{nn}$.} \\item{$^{\\mathrm{a}}$ $f({\\rm H_2})= {2 N({\\rm H}_2)}/\\left({2  N({\\rm H}_2) + N({\\rm \\ion{H}{i}})}\\right)$.} --- {$^{\\mathrm{b}}$ Inferred from $E(B-V)$.} \\end{list} \\end{table*}  Using the Meudon PDR model (Nehm\\'e et al. to be published), \\citet{2002A&A...390..369L} studied the properties of HD in a diffuse cloud with $n_{\\rm H} = 500$ cm$^{-3}$ embedded in the standard interstellar radiation field.  Under those conditions, HD becomes the reservoir of deuterium at a visual extinction of 1 magnitude, where the molecular fraction $f$(H$_2$) of hydrogen is close to 1.  Toward our sightlines, $f$(H$_2$) does not reach such a value.  However, if we assume in each case that (1) a significant part of the atomic hydrogen is not associated with the observed diffuse molecular material, and (2) all the molecular material is in one dominant cloud, then the D/H ratio in that main cloud should be equal to the $N$(HD)/2$N$(H$_2$) ratio.  Table~\\ref{dens sum} shows the $N$(HD)/2$N$(H$_2$) ratios for all the sightlines included in this paper, i.e., our {\\it FUSE} measurements and the values we have re-derived from the {\\it Copernicus} data. All these values are plotted in Fig.~\\ref{HDsurH2}, where we have added two reference values: the Local Bubble $N$(D)/$N$(H) ratio of $1.52 \\times 10^{-5}$ \\citep{2002ApJS..140....3M}, and the \\citet{1999ApJ...520..182J} value of $0.74 \\times 10^{-5}$ toward $\\delta$ Ori A \\citep[First detection by][]{1979ApJ...229..923L}. Several recent observations \\citep{2003ApJ...599..297H,2003ApJ...586.1094H,2004ApJ...609..838W} have suggested that the LB value may not be representative of the general ISM and that a ``canonical'' (D/H)$_{\\rm ISM}$ (if such exists) may be closer to the latter value.  For our sample, the average $N$(HD)/2$N$(H$_2$) yields a corresponding estimate for D/H of (2.0 $\\pm$ 1.1) $\\times 10^{-6}$, a factor 7 below the LB value and a factor 3 below the $\\delta$ Ori A value.  Such a low value is difficult to interpret in terms of enhanced stellar astration or favored deuterium depletion onto dust grains.  Indeed, given the multi-component velocity structure of the sightlines and the many special conditions that may change the HD/H$_2$ ratio from component to component, the most likely explanation is that we have not yet identified a predominantly molecular cloud.  We now turn to a more realistic view of the sightlines in which several molecule-containing clouds are present, with only some of them having all the deuterium in the molecular form.  To some extent we can expect that the molecular fraction $f$(H$_2$) will be a relevant parameter.  $f$(H$_2$) will increase with the average UV shielding and decrease with the fraction of smaller, more diffuse molecular clouds present. The models suggest that the deuterium is molecular only when hydrogen is fully in its molecular form \\citep[Fig. 1 of][]{2002A&A...390..369L}.  This can be summarized over a sightline by \\begin{equation} f({\\rm HD}) \\leq f({\\rm H_2}) \\, , \\end{equation} while the molecular fractions are linked together by \\begin{equation} \\left(\\frac{\\rm D}{\\rm H}\\right)_{\\rm ISM} = \\frac{f({\\rm H_2})}{f({\\rm HD})} \\times \\frac{N({\\rm HD})}{2 N({\\rm H_2})} \\, . \\end{equation} The average value of $N$(HD)/2$N$(H$_2$) noted above should therefore only be taken as a stringent lower limit to the (D/H)$_{\\rm ISM}$. Since the two molecular fractions reflect the average effectiveness of the self-shielding of each species, they should be highly correlated, even when summing over multiple clouds. Unlike the top panel, the bottom panel of Fig.~\\ref{HDsurH2} shows an increasing trend for $N$(HD)/2$N$(H$_2$) as a function of $f({\\rm H_2})$ -- and therefore a correlation between $f({\\rm HD})$ and $f({\\rm H_2})$ -- consistent with predictions from models of the chemistry of HD \\citep{2002A&A...390..369L}.  Toward our highly reddened targets, atomic deuterium is blended by saturated hydrogen lines, and almost impossible to measure; direct comparison between the two molecular fractions is not feasible.  We therefore used the two (D/H) ratios noted above \\citep{1999ApJ...520..182J,2002ApJS..140....3M} to examine $f({\\rm HD})$ vs. $f({\\rm H_2})$ (Fig.~\\ref{FH_FD}).  The increasing trend is obvious.  If the relation between $f({\\rm HD})$ and $f({\\rm H_2})$ can be understood, it could enable a wider survey of the D/H ratio via observations of HD.  The apparent correlation between the molecular fractions should first be confirmed by observing targets in which both D and HD are measurable.  That, however, requires targets with low overall column density but high $f$(H$_2$).  Another possibility would be to examine individual clouds, using very high resolution data in the FUV (unfortunately not available at present).  Further chemical modeling of diffuse clouds, including sums over a statistical distribution of diverse sizes and densities, could serve both to confirm the correlations and also to reveal other important parameters linked to the abundance of HD.  For example, the formation rate of HD (and therefore the molecular fraction) is directly proportional to the ionization of H and H$_2$ by cosmic rays \\citep{1996ApJ...463..181F,2002A&A...390..369L}.  It has been suggested that the observed abundances of H$_3^+$ in some diffuse lines of sight require a high flux of cosmic rays \\citep{2000MNRAS.313L...6C,2003Natur.422..500M,2004A&A...417..993L} and that the flux of cosmic rays could be very inhomogeneous in the diffuse ISM (for example due to variation in the magnetic field).  The relationship between the molecular fractions of H$_2$ and HD would provide information on the formation rates of both molecules and on the flux of cosmic rays; the variations could give some indication of the level of inhomogeneities in the cosmic ray flux.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{FH_FD.eps}} \\caption{$f({\\rm HD})$ vs $f({\\rm H_2})$. The left panel assumes that (D/H)$_{\\rm ISM} = 0.74 \\times 10^{-5}$, and the right panel assumes that (D/H)$_{\\rm ISM} = 1.52 \\times 10^{-5}$.  Both panels show a clear correlation between the molecular fractions. Understanding this correlation might provide a way to deduce (D/H)$_{\\rm ISM}$ from HD.  Filled diamonds have $N$(HD) derived from {\\it FUSE} spectra; open squares have $N$(HD) derived from reanalysis of {\\it Copernicus} data (see Sect.~\\ref{sec:CHHD}). \\label{FH_FD}} \\end{figure}  \\subsection{Summary}  We derived column densities of HD for 7 reddened sightlines using far-UV spectra obtained with the {\\it FUSE} telescope.  Most of the HD lines available in the far-UV are strongly affected by saturation (see the COG plotted in Fig.~\\ref{COG}).  We used high resolution optical CH data to determine the velocity structure in these sightlines, which allowed us to correct for the that saturation.  The analysis was done using both curve of growth and profile fitting methods.  We combined our new results for HD with re-analyzed {\\it Copernicus} measurements and compared the column densities of HD, H$_2$, and CH. A correlation between the HD and CH column densities is likely, but further measurements at low column densities are needed.  Simulations \\citep{2002A&A...390..369L} have predicted for a one-cloud model with $A_V$ $\\sim$ 1 mag a $N$(HD)/2$N$(H$_2$) ratio equal to the D/H ratio, and lower for thinner clouds.  All our sightlines have $A_V$ $>$ 1, but are unfortunately composed of multiple clouds.  We therefore obtained only stringent lower limits for (D/H)$_{\\rm ISM}$ ranging from $3.7 \\times 10^{-7}$ to $4.3 \\times 10^{-6}$.  We also noted a correlation between the molecular fraction of H$_2$ and the ratio $N$(HD)/2$N$(H$_2$) which we linked to a relationship between the self-shielding of H$_2$ and HD.  If that relationship is confirmed, it would give a mean to infer (D/H)$_{\\rm ISM}$ from observations of HD."
    },
    "0410/astro-ph0410519_arXiv.txt": {
        "abstract": "Using the NIRSPEC spectrograph at Keck II, we have obtained infrared echelle spectra covering the range $1.5-1.8~\\mu \\rm m$ for the moderately reddened bulge globular clusters NGC~6342 and NGC~6528, finding [Fe/H]=--0.60 and --0.17~dex, respectively. We measure an average $\\alpha$-enhancement of $\\approx+0.33$~dex in both clusters, consistent with previous measurements on other metal rich bulge clusters, and favoring the scenario of a rapid bulge formation and chemical enrichment. We also measure very low  $\\rm ^{12}C/^{13}C$ isotopic ratios ($\\approx$5 in NGC~6342 and $\\approx$8 in NGC~6528), suggesting that extra-mixing mechanisms due to {\\it cool bottom processing} are at work during evolution along the Red Giant Branch. ",
        "introduction": "Infrared, high resolution spectroscopy has tremendous potential as a powerful tool to study both distant and obscured stellar populations. For the heavily obscured regions of the Galactic bulge and center, infrared spectroscopy offers the best approach to measuring the composition of the old stellar populations.   Over the past few years we have commenced a high-resolution spectroscopic survey of the Galactic bulge in the near-IR using NIRSPEC, a high throughput infrared (IR) echelle spectrograph at the Keck Observatory \\citep{ml98}. H--band (1.5--1.8 $\\mu$m) spectra of bright giants in the bulge globular clusters and field population are ideal for detailed abundance analysis of Fe, C, O and other $\\alpha $-elements, using the approach of synthesizing the entire spectrum.  By observing the most luminous giants in the infrared, we have the possibility of placing the obscured stars and globular clusters toward the galactic center, and populations well studied in the optical, on the same abundance scale. The abundance distributions in the {\\it cluster} and {\\it field} populations are important in constraining the history of bulge formation and chemical enrichment \\citep{mw97}.  We have used this method to derive abundances for four bulge globular clusters: the results for NGC~6553 and Liller~1 are given in \\citet{ori02}, while abundances for Terzan~4 and Terzan~5 are reported in \\citet{ori04}. We find $\\alpha$-enhancement at a level of a factor between 2 and 3 over the whole range of metallicity spanned by the clusters in our survey, from [Fe/H]$\\approx$--1.6 (cf. Terzan~4) up to [Fe/H]$\\approx$--0.2 (cf. Terzan~ 5). In this paper we present the high resolution IR spectra and the abundance analysis of four bright giants in NGC~6342 and NGC~6528, two bulge globular clusters with relatively low reddening.  The most recent metallicity estimates based on optical \\citep{hr99} and near-IR photometry \\citep{mom03,val04} report global metallicities [M/H]$\\approx-0.4$ and $\\approx+0.0$ dex and iron abundance s [Fe/H]$\\approx-0.6$ and $\\approx-0.2$ for NGC~6342 and NGC~6528, respectively. NGC~6342 has never been studied at high spectral resolution before, while NGC~6528 has been observed with HIRES at Keck I by \\citet{car01}. They found iron abundance slightly in excess of solar ([Fe/H]=+0.07) and $\\alpha$-enhancement of Ca ([Ca/Fe]=+0.2) and Si ([Si/Fe]=+0.4) and marginal enhancement of O, Mg and Ti. Very recently, \\citet{zoc04} observed NGC~6528 giants with UVES at VLT, finding [Fe/H]=--0.15 and moderate $\\alpha$-enhancement of O, Si and Mg by $\\approx$+0.1--0.2 dex.  Our observations and data reduction follow in Sect.~2. Sect.~3 discusses our abundance analysis and in Sect.~4 the resulting metallicities and radial velocities are presented. We discuss our findings in Sect.~5. ",
        "conclusions": "Our high resolution spectroscopy in the near IR gives iron abundances for NGC~6342 and NGC~6528 which are in excellent agreement with the photometric estimates, as obtained from the RGB morphology and luminosity in the optical as well as in the near IR.  Our iron abundance of NGC~6528 is also in excellent agreement with the one estimated by \\citet{zoc04} from optical spectroscopy, while is slightly lower ($\\approx$0.2 dex) than that of  \\citet{car01}. Such a slightly higher value of iron abundance is still marginally consistent at a $\\ge 2\\sigma$ level with our IR spectra. A similar iron abundance discrepancy between IR \\citep{ori02,mel03} and optical measurements by \\citet{car01} has been found for NGC~6553.  Our O, Ca, Mg and Ti abundances of NGC~6528 are fully consistent with the study of red HB stars by \\citet{car01}, while our Si abundance is lower by $\\approx$0.3 dex. Our O, Si, and Mg abundances are also consistent (only marginally higher) with those for red giants (though less luminous than ours) by  \\citet{zoc04}. The abundances of Ti and Ca, in particular, by \\citet{zoc04} \\citep[see also][]{mr04} are significantly lower than those from our and \\citet{car01} studies.  Our composition is also consistent with that found for Galactic bulge field stars \\citep{mr94}. Photometry of NGC 6528 and a comparably metal rich bulge globular cluster, NGC 6553, indicates that these two clusters are as old as the Galactic halo with no measurable age difference \\citep{ort95}.  This has been reinforced by recent photometry in which the cluster members have been separated from the field by their relative proper motions \\citep{zoc01,felt02}. The results of our study are consistent with NGC 6342 and 6528 being related to the field population of the Galactic bulge, their composition bearing the imprint of an early era of enrichment by massive star supernovae. We find it difficult to entertain the notion that type I SNe contributed significantly to the enrichment of NGC 6528, or that the composition of this cluster is different from the other bulge clusters we have studied.  However, the significant dispersion in the composition of NGC6528, even based on analyses using high S/N spectra from 8-10m class telescopes, emphasizes that the composition of metal rich stars is far from being a solved problem.  The low $\\rm^{12}C/^{13}C$ abundance ratios measured in NGC~6342 and NGC~6528 are similar to those measured in bright giants of halo \\citep{ss91,she96,gra00,vws02,smi02,ori03} as well as other bulge \\citep{ori02,she03,ori04} globular clusters, over the entire range of metallicities between one hundredth and solar. They can be explained by additional mixing mechanisms due to {\\it cool bottom processing} in the stellar interiors during the evolution along the RGB \\citep[see e.g.][]{cha95,dw96,csb98,bs99}.    \\begin{figure} \\centering \\includegraphics[width=8.7cm]{origlia.fig3.ps} \\caption{ Probability of a random realization of our best-fitting solution and the test models with varying temperature by $\\Delta T_{eff}$ of $\\pm$200K (left panels), gravity by $\\Delta log~g$ of $\\pm$0.5 dex (middle panels), and microturbulence by $\\Delta \\xi$ of $\\mp$0.5 Km s$^{-1}$ (right panels), with respect to the best-fitting (see Sect.~\\ref{results}) for the four giants in NGC~6342. } \\label{fig3} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=8.7cm]{origlia.fig5.ps} \\caption{ Probability of a random realization of our best-fitting solution and the test models with varying temperature by $\\Delta T_{eff}$ of $\\pm$200K (left panels), gravity by $\\Delta log~g$ of $\\pm$0.5 dex (middle panels), and microturbulence by $\\Delta \\xi$ of $\\mp$0.5 Km s$^{-1}$ (right panels), with respect to the best-fitting (see Sect.~\\ref{results}) for the four giants in NGC~6528. } \\label{fig5} \\end{figure}"
    },
    "0410/astro-ph0410205_arXiv.txt": {
        "abstract": "The CELESTE atmospheric Cherenkov detector ran until June 2004. It has detected several $\\gamma$-ray flares of Mrk~421 since 1999. The new data analysis provides better background rejection. We compare our light curves with X-ray data. We significantly improved our understanding of the atmosphere, and of the optical throughput of the detector. This will allow a spectral measurement for Mrk~421 with smaller uncertainties and a more sensitive search for a signal from Mrk~501. ",
        "introduction": "CELESTE was a Cherenkov experiment using 53 heliostats of the former \\'Electricit\\'e de France solar plant in the French Pyrenees at the Th\\'emis site. Its principle is to detect the Cherenkov light of the electromagnetic shower produced in the atmosphere by the $\\gamma$-rays coming from high energy astrophysical sources. The light is reflected to secondary optics and photomultipliers installed at the top of the tower. Finally, it is sampled to be analysed (Par\\'e 2002).  We significantly improved our data analysis and have better background rejection since the 2000 status (de Naurois 2002). Changes, described elsewhere in this volume (Manseri 2004), were made in hardware (new heliostats and pointing strategy) and in software (new data analysis). The Crab sensitivity was improved from 2.2~$\\sigma/\\sqrt{\\mathrm{h}}$ to 5.8~$\\sigma/\\sqrt{\\mathrm{h}}$. We present the improved light curves for Mrk~421. We also now better understand the optical throughput of the detector with our new simulation. ",
        "conclusions": ""
    },
    "0410/hep-ph0410326_arXiv.txt": {
        "abstract": "We study the evolution of non-interacting and entangled cosmic string networks in the context of the velocity-dependent one-scale model. Such networks may be formed in several contexts, including brane inflation. We show that the frozen network solution $L\\propto a$, although generic, is only a transient one, and that the asymptotic solution is still $L\\propto t$ as in the case of ordinary (intercommuting) strings, although in the present context the universe will usually be string-dominated. Thus the behaviour of two strings when they cross does not seem to affect their scaling laws, but only their densities relative to the background. ",
        "introduction": "Cosmic strings \\cite{VS} are topological defects that might have formed at phase transitions in the early universe. Interest in their evolution and cosmological consequences has recently grown, for both observational and theoretical reasons. On one hand, it is possible that a cosmic string has been identified as the source for an otherwise unexplained gravitational lens \\cite{Sazhin1,Sazhin2}, or possibly even through a Kaiser-Stebbins effect in the cosmic microwave background. On the other hand, it has been recently realized that they can have a crucial role to play within superstring theory \\cite{Copeland}, and in particular will be inevitably produced at the end of brane inflation \\cite{Sarangi}.  Depending on the specific models, these cosmic strings may or may not have similar properties to the standard ones. An obvious difference is that they will appear in models with extra dimensions, though this alone does not guarantee that their evolution must be different. Another specific difference can potentially be more important, though. When two standard $U(1)$ strings interact, they always exchange partners (or intercommute), but in other specific models they may intercommute only with probability less than unity, pass through each other, or even entangle themselves via the formation of a bridge between the points of contact. It seems to be the case that these non-standard outcomes will be quite common in the higher-dimensional context, so it is interesting to ask what effect this will have, if any.  The evolution of standard cosmic strings has been extensively studied, both by numerical \\cite{bb,as,Moore,unified} and by analytic means \\cite{Kibble,Kibble2,Bena1,Bena2,cka,ack,ms1a,ms1b,ms2b,extend}. Here one finds that after an initial transient period (whose duration depends on the string mass scale, being shorter for heavier strings) the network will evolve in a linear scaling regime, with a characteristic length (or correlation length) being a constant fraction of the horizon, $L\\propto t$, and the RMS velocity being also constant.  Relatively little is known about the evolution of string with non-standard interactions. Relatively simple analytic arguments \\cite{Dom2} suggest that they are frozen and conformally stretched ($L\\propto a$) and end up dominating the energy density of the universe. The scenario of standard intercommuting strings dominating the universe was also discussed \\cite{Dom1}, but it can never be realized for realistic networks. On the other hand, relatively small numerical simulations \\cite{Vachaspati,Spergel,McGraw} report that linear scaling is reached except in special circumstances.  In this note we will use a suitably modified version of the velocity-dependent on-scale model of string evolution \\cite{ms1a,extend} to clarify this issue. We will restrict ourselves to the standard case of three spatial dimensions, though we shall comment on the expected differences in the case of higher dimensions. It will be shown that, just like in the standard case, an $L\\propto a$ phase may exist but it will necessarily be a transient. Due to the significant string velocities, the network will tend to evolve towards a linear scaling regime, albeit one where the strings will usually be dominating the energy density of the universe. ",
        "conclusions": "We have used the velocity-dependent one-scale model to quantitativel study the scaling behaviour of non-interacting and entangled string networks, thus clarifying some previously existing claims. We have shown that any $L\\propto a$ scaling must be a transient, simply because of the effect of the increasing velocities, though usually this regime is ended when the strings dominate the energy density of the universe. When strings dominate, linear scaling ($L\\propto t$, $v=const.$) is the attractor solution, and the evolution of the scale factor of the universe depends on the string velocity. Usually this is negligible and it will grow as $a\\propto t$, but for non-interacting strings the corrections are important---see (\\ref{noninovwithv}). We have also shown that the transient scaling laws for friction-dominated evolution still hold for non-interacting and entangled networks, albeit with correspondingly higher string densities.  Although non-interacting and entangled cosmic string networks can be formed in a number of contexts in standard $(3+1)$-dimensional scenarios (see \\cite{VS} for a review) the current wisdom seems to indicate that they should be more common  in higher dimensions. For example, two ordinarily intercommuting strings might miss each other when they cross if there are extra spatial dimensions. While it is beyond the scope of this note to analyse this case in any detail, we can nevertheless try to draw some lessons from the results above.  The main point is that $L\\propto t$ scaling is a generic attractor: cosmic string networks will eventually be straightening out as fast as is allowed by causality. The background in which they find themselves and their own properties (what happens when two cross, their mass per unit length, and so on) can only influence the density and velocity at that linear scaling regime, and also what transient regimes might exist until linear scaling is reached. We expect that this will still be the case in higher dimensions. The main difference is likely to have to do with the fact that the sizes of any extra dimensions will be quite different from those of the three ordinary ones. One therefore expects that there will be anisotropies \\cite{fossils} in the orientations and velocities of the strings, which could conceivably have observational consequences. Modelling such effects should require more than one correlation length and characteristic velocity. In that sense they would be somewhat analogous to models for wiggly \\cite{wiggly} or superconducting \\cite{supercond} strings."
    },
    "0410/astro-ph0410716.txt": {
        "abstract": "Observational evidence for anisotropy in supernovae (SN) may signal the importance of angular momentum and differential rotation in the progenitors. Free energy in differential rotation and rotation can be extracted magnetically or via turbulent dissipation. The importance that magnetohydrodyamic jets and coronae may play in driving SN motivates understanding large scale dynamos in SN progenitors.  We develop a dynamical large scale interface dynamo model in which the differential rotation and rotation deplete both through Poynting flux  and turbulent diffusion. We apply the model to a differentially rotating core surrounded by a convection zone of a SN progenitor from a initial 15$M_\\odot$ star. Unlike the Sun, the dynamo is transient because the differential rotation  is primarily due to the initial collapse. Up to $\\sim 10^{51}$erg can be drained into time-integrated Poynting flux and heat, the relative fraction of which depends on the relative amount of turbulence in the shear layer vs. convection zone and the fraction of the shear layer into which the magnetic field penetrates. % long as the turbulence in the shear layer is not too large, %the %integrated Poynting flux can tap up to $\\sim 40\\%$ %of the free-energy, but the dominant sink of the free %energy is typically turbulent dissipation. Both sinks can help facilitate explosions and could lead to different levels of anisotropy and pulsar kicks. In all cases, the  poloidal magnetic field is much weaker than the toroidal field, and the  Poynting flux is lower than previous estimates which  invoke the magnitude of the total magnetic energy. A signature of  a large scale dynamo is that the oscillation of the associated Poynting flux on $\\sim 1$ sec time scales, implying the same for the  energy delivery to a SN. %For a 15$M_\\odot$ progenitor with inner shear layer rotation rates %$ \\sim 200$rad/s, dynamo durations are $5-50$ sec and provide %up to $\\sim 5 \\ts 10^{50}$ erg from the associated Poynting flux. ",
        "introduction": "Observations suggest that many supernovae (SN) are intrinsically anisotropic.  As summarized by Wheeler (2004a,b) evidence comes from: (1) bipolar structure in supernova remnants (Dubner et al.\\ 2002); (2) optical jet and counter jet structures in Cas A (Fesen 2001); (3) X-ray jets and toroidal emission of intermediate mass elements (Hughes et al.\\ 2000; Hwang, Holt, Petre 2000; Willingale et al.\\ 2002); (4) asymmetric ejecta in SN1987A perpendicular to the major axis of the rings (%Pun et al.\\ 2001; (Wang et al.\\ 2001, 2002); and (5) Optical polarization in both Type I and Type II supernovae (see Wang 2004 for a review) with the Type Ib,c polarization being higher than that of Type II (Wang et al. 1996; Wang et al.\\ 2001), but with the latter increasing with time (Wang et al. 2001; Leonard et al. 2001). Since Type I SN represent a naked core, and the core at late times in Type II becomes exposed, a consistent interpretation is that the SN engine is a source of this asymmetry. At minimum, anisotropy likely reveals the role of rotation in SN engines and the need to incorporate non-spherically symmetric physics.  \\subsection{Magnetic Fields} The combination of rotation, highly ionized plasma, stratified turbulence, and magnetic fields has prompted various considerations of MHD outflows as a source of this asymmetry, if not the source of the SN explosion itself. Recent renewed interest in  proposals relating MHD outflows to supernovae (Leblanc \\& Wilson 1970; Meier et al.\\ 1976; Ardeljan, Bisnovatyi-Kogan, \\& Moiseenko 1998; Khokhlov et al.\\ 1999; Wheeler et al.\\ 2000; Wheeler, Meier, \\& Wilson 2002; Moiseenko, Bisnovatyi-Kogan, \\& Ardeljan 2004) is bolstered by the observational association of SN with Gamma-ray bursts (Galama et al 1998; Iwamoto et al.\\ 1998; Stanek et al. 2003; Hjorth et al. 2003; Thomsen et al. 2004; Cobb et al.\\ 2004 Gal-Yam et al. 2004; Malesani et al. 2004). %and related NS or black hole jet GRB %(e.g.  Blackman et al.\\ 1996; Wheeler et al.\\ 2000; MacFadyen \\& Woosley 1999; MacFadyen, Woosley, \\& Heger 2001). That magnetized outflows have also long been thought to be important in young stellar objects, active galactic nuclei, pulsar winds, and  microquasars (of which GRB are likely one class) suggests  that explosive bipolar outflows in astrophysics may all involve some combination of rotation and magnetic fields.  Previous work on  magnetic field amplification in pre-supernovae core conditions have employed either (i) simple shear that grows toroidal field linearly (e.g. Leblanc \\& Wilson 1970; Ardeljan, Bisnovatyi-Kogan, \\& Moiseenko 1998; Wheeler 2000), (ii) traditional kinematic convective $\\alpha-\\Omega$ dynamo models to grow ordered fields (Thompson \\& Duncan 1993), which in principle grow large-scale fields exponentially, or (iii) the exponential field growth from the magneto-rotational instability (MRI, e.g. Balbus \\& Hawley 1998) (Akiyama et al.\\ 2003; Moiseenko et al.\\ 2004), whose saturated magnetic energy was used as a rough estimate for the large-scale toroidal field. Each has its merits and its limitations. In the presence of turbulent diffusion, (i) will not sustain the field. The approach of (ii) offers exponential growth and is an important step forward, but the backreaction of the field on the driving flow, the separate generation of toroidal vs. poloidal fields, and the spatial location of the $\\alpha$ and $\\Omega$ effects still need to be determined. The recent approach of Akiyama et al.\\ in (iii) offers a useful model of the internal and rotational structure of the inner SN engine, and also makes careful estimates of the saturated energy of the magnetic field adopted from results obtained in MRI disk simulations. Whether these values also apply to pressure-supported stars remains to be studied in detail because the MRI likely transports angular momentum on spherical shells (Balbus \\& Hawley 1994), not radially; the latitudinal differential rotation may be most important. If 3-D turbulence develops, then field amplification will take place from the MRI. Moiseenko et al.\\ (2004) provide 2-D simulations of what appear to be MRI-driven SN, starting from a relatively strong, ordered poloidal field. However, in 2-D, there is not sustained dynamo action to amplify the total magnetic energy from arbitrarily small values.  An important question is whether a magnetically driven outflow requires a large-scale field in the supernova progenitor engine. By large-scale field, we mean one that maintains the sign of its flux over many dynamical times, when averaged over the size of the engine region. Large-scale fields are helpful, if not necessary, in  generating outflows in disks or stellar coronae, and perhaps by analogy, also in SN engines. If the field needs to  buoyantly rise above the dynamo region to a height at which it provides the dominant contribution to the stress, the field should be of large enough scale to avoid being shredded by turbulence. Once at the region where it dominates, the field can further relax to even larger scales. %A simple magnetic pressure gradient from %a random field might not be enough if the field is produced %in a high beta regime. Comparing the shapes of observed outflows to those from theoretical studies (e.g. Moiseenko et al.\\ 2003) may help determine the relative importance of large scale fields vs. a simple magnetic pressure gradient.  This magnetic field may drive a  ``jack-in-the-box'' type of  magnetic spring explosion. (In addition to Wheeler, Meier, \\& Wilson (2002), see also Matt et al.\\ (2004) for simulations of a magnetic explosion without specifying the field origin, and Moiseenko et al.\\ (2004) for a magnetic explosion driven SN.). %We plot values of $\\tau_b=E_{rot}/L_{mag}$ in Fig. 6, where $E_{rot}$ is %either (a) the rotational energy of the NS, $0.5 M_{ns}\\Omega_{ns}^2r_{ns}^2$, %or (b) the rotational energy of the shear layer, %${2\\pi\\over 3}[r_c^3-(r_c-L)^3]\\rho_2 v_2^2$. %The reason for considering the two choices is the uncertainty in the extent %to which the NS rotational energy couples with that for  $r\\ge r_c-L$. % The potential efficacy of magnetically-driven outflows for SN is evident from recent work by Moiseenko et al.\\ (2004). However, as in the magnetic explosion of Matt et al (2004) for planetary nebulae, Moiseenko et al. (2004) start with a significant large-scale poloidal field, comparable in strength to the saturated poloidal field whose growth we will study here from an initially small seed field. Any dynamo driven outflow must occur while the dynamo is active because afterward the field decays. %Because the large-scale toroidal and poloidal fields are out of phase, %the associated Poynting flux oscillates with period$\\sim 1$ sec. %If a SN is influenced by a large-scale dynamo outflow %then the injection would show $\\sim 1$ sec variability %during the injection duration.      Large-scale field generation does not exclude the presence of the MRI: When the MRI operates in a stratified medium, it may provide a source of helical turbulence that allows one sign of magnetic helicity %ERIC - NOTE CHANGE IN LINE ABOVE, TO SMALL-SCALE MAG. HEL. - OKAY? to migrate to large scales, producing large-scale magnetic fields (e.g. Brandenburg et al.\\ 1995; Blackman \\& Tan 2004) or a helicity flux (Blackman \\& Field 2000; Vishniac \\& Cho 2001). %but some version of large-scale %dynamo theory, which can also determine separate values of the %toroidal and poloidal fields, would be needed to describe this. In our large scale dynamo herein, the source of helical turbulence is considered to be convection driven by shock heating and neutrino deposition rather than the MRI, though the particular source is not essential for our calculations. It should also be emphasized that  studies focusing on magnetic energy amplification from the MRI (Balbus \\& Hawley 1998) typically study growth of the total magnetic energy (e.g. Akiyama et al. 2003) rather than a specifically large-scale  field in the sense we have defined above. The backreaction of the field on the shear is not usually considered in MRI studies.  Note that instead of a bulk dynamical influence of large scale fields, Ramirez-Ruiz and Socrates (2004) highlight an alternative role facilitated by large scale fields that buoyantly rise to form a corona. They  suggest that such a corona could distribute the binding energy dissipation such that the neutrino spectrum becomes non-thermal. This increases the efficiency of neutrino-matter coupling by an order of magnitude. The SN would be neutrino driven, but symbiotically dependent on a corona.  \\subsection{Turbulent Viscosity}  Magnetic fields generated in situ are amplified by extraction of the free energy in differential rotation. Thompson et al. (2001) point out that the energy lost into heat via turbulent damping of the shear supplies enough energy when combined with neutrino driving to create a SN explosion. The energy lost from the shear to heat represents a conversion of some of the binding energy into a form more efficiently coupled to the matter than neutrinos. In the presence of turbulence (which itself could be induced by the magnetic field) a competition emerges between draining the shear into Poynting flux  vs. heat. Both sinks may help drive  SN in different ways.    \\subsection{What We Do Here}   Here we study the dynamical evolution of differential rotation and large scale magnetic field growth with several key new features: 1) we explicitly consider the backreaction of the magnetic field on the shear and rotation driving the field growth and 2) allow the shear and rotation to deplete via Poynting flux and  turbulent dissipation.  The latter turns out to be important and a key result of our study is to show the relative  depletion into each channel.  We derive a dynamical generalization of the large-scale interface dynamo proposed by Parker (1993) for the Sun (see also Charbonneau \\& MacGregor 1997; Markiel \\& Thomas 1999; Zhang et al.\\ 2003). Interface dynamos differ  from the original $\\alpha-\\Omega$ dynamos in that the dominant region of shear and the dominant region of helical turbulence are not co-spatial. The shear associated with differential rotation is concentrated in a layer (the $\\Omega$-layer) lying beneath the turbulent convection zone (the $\\alpha$ layer). Such a scenario has been a leading model for the solar dynamo, and has also been applied to white dwarfs (Markiel, Thomas \\& Van Horn 1994; Thomas, Markiel \\& Van Horn 1995) and to AGB stars (Blackman et al.\\ 2001). In all of these cases, helical turbulence in the convective zone provides the dynamo $\\alpha$  effect and thus generates the poloidal field. Turbulent pumping (e.g. Tobias et al.\\ 2001) pushes the poloidal field downward into the strong shear layer where the toroidal field is amplified. The rotational and convective environment of the proto-supernova engine in core-collapse supernovae has a similar structure, with a strong shear layer surrounded by a convective envelope. However, unlike the sun where the differential rotation profile is re-established by convection, in the SN context the intial shear decays dynamically as the field grows. Whether or not convection may refuel the differential rotation needs further study, but here we assume that the dominant source of differential rotation is that from the initial collapse. This differential rotation will deplete as the field is amplified or as the shear is damped by turbulence. We explicitly include the dynamical evolution of the shear and rotation in addition to that of the magnetic field, unlike previous work.  The  dynamo becomes transient in the SN progenitor.  Dynamo research must accommodate two competing needs: the need for a theory that rigorously includes the backreaction of the growing magnetic field on the flow, and the need to model the field growth in a realistic system. Meeting these two goals is difficult both analytically and numerically, and compromises must be made on both fronts. Our aim here is to develop a  transient large scale dynamo model and apply it to the SN progenitor context as simply as possible, while still including the time-dependent nonlinear dynamical quenching of the shear and  rotation.  In section 2 we present the dynamical equations for the magnetic field, shear, and rotation. In section 3 we discuss the parameter choices and present the corresponding solutions. In section 4 we discuss the implications of the calculated Poynting flux and the fractional shear energy drained into magnetic energy vs. heat. We conclude in section 5 and identify some unresolved issues for future work. ",
        "conclusions": "The free energy in rotation and differential rotation comprise an important source of energy that can be tapped for aiding, if not driving SN explosions. This energy can be drained into  magnetic fields or into heat from turbulent dissipation and both can help power a SN: The magnetic field can mediate a SN either via direct Poynting flux driven outflows, or via a corona that leads to a non-thermal neutrino spectrum (Ramirez \\& Socrates 2005). The heat from turbulent dissipation of shear can conspire with neutrino heating to drive the  explosion (Thompson et al. 2005). Here we have investigated the relative importance of these two energy sinks by developing  a  dynamical, transient, $\\alpha-\\Omega$ interface large scale dynamo model and applying it to the proto-supernova engine structures presented in Akiyama et al.\\ (2003). Our interface dynamo, like that commonly thought to be operating in the sun, is a large-scale helical dynamo in which the shear layer providing the $\\Omega$ effect lies beneath the turbulent convection supplying the helical $\\alpha$ effect. For core-collapse proto-supernovae engines, the convection is driven by shock heating %neutrino deposition and the shear layer beneath it derives from the initial stellar core collapse. Unlike the sun however, the much stronger shear for the SN progenitor comes from the initial collapse and is not necessarily reseeded by convection. We have therefore incorporated the dynamical backreaction of the growing field on the shear, the decay of shear due to turbulent dissipation, and the extraction of rotational energy via Poynting flux. Although we also parameterize  $\\alpha$ quenching and buoyancy, the magnetic quenching and/or turbulent dissipation of the shear dominates the  dynamo quenching and makes the dynamo transient.   The transient dynamo applied to a 15$M_\\odot$ progenitor with maximal inner shear layer rotation rates of $200$rad/s lasts $\\sim 10-50$ seconds, leading to large-scale toroidal fields of strength $5 \\ts 10^{14}$G and large-scale poloidal fields of strength $\\sim 10^{13}$G at the base of the convection zone. The poloidal magnetic field has a significantly lower magnitude than the toroidal field, lowering the peak Poynting flux (which depends on the product of the two field components) compared to previous estimates that invoke the magnitude of the magnetic energy (e.g. Wheeler, Meier, \\& Wilson 2002). However extracting even 10\\% of the available rotational energy via  the integrated Poynting flux is influential in driving or making the SN anisotropic (Wheeler et al. 2000) %and/or %turbulent dissipation may be sufficient to drive the SN, %possibly in combination with  neutrinos, %and/or to drive an outflow through the stellar %envelope (Wheeler et al. 2000) that would make the SN anisotropic: because the binding energy above the core for a $15M_\\odot$ progenitor is  only $\\sim 5\\ts 10^{50}$erg (Woosley, Heger, Weaver 2002). This is less than the total rotational free energy  ($\\sim 10^{51}$erg/s) for cores rotating with $\\sim 160$ rad/sec. In addition, extracting  a modest $1\\%$  of the into pole-asymmetric outflows can supply observed pulsar kicks. Faster rotation rates than invoked here can supply even more energy, but very high rotation rates are beyond the applicability of the structure model we adopted from Akiyama et al. (2003), %and the curves of Fig. 6 start to flatten at large $\\Omega_0$. % %Because the  Poynting flux is transient, any outflow must be driven during the dynamo duration.%   A signature of the influence of a large scale magnetic dynamo is the $\\sim 1$ sec cycle periods which are also reflected in the Poynting flux energy deposition to the SN explosion (Figs. 2-5). Despite the importance of magnetic fields, it may be that only a fraction (determined by how far the toroidal field can diffuse downward in a cycle period) of the shear layer is available for magnetic energy amplification. As seen in Figs 2 and 3, the dominant energy sink for the integrated shear layer is heat. A SN induced from this heat would likely be less anisotropic than in the predominantly magnetically driven case.   %$\\tau_b$. %Observational constraints on the time scale for SN %energy injection are needed.       Our dynamo formalism differs from the $\\alpha-\\Omega$ dynamo of Duncan \\& Thompson (1993) which operates deeper inside the NS. % %large-scale toroidal field can have a  magnitude $\\sim 10^{14-15}$ G. Our results are also complementary to the MRI-based estimates of Akiyama et al.\\ (2003), who obtained magnetic energies associated with a $10^{15}$ G field, but inferred these values from turbulent magnetic energy saturation of the MRI turbulence rather than from specific values of a large-scale magnetic field.  Also unlike these previous works, we  include the dynamical evolution of the shear and rotation.    Similar dynamo and outflow studies  for faster rotating SN progenitors, and for more massive ``failed'' SN (MacFadyen \\& Woosley 1999; MacFadyen, Woosley, \\& Heger 2001) are also desired.Poynting flux produced from within such engines is a leading candidate to power GRB.  The predicted variability on time scales $\\sim 1$ s, based on the relative phase of toroidal and poloidal fields in our dynamo cycle, suggests an analogous variability in Poynting-flux dominated GRB models.   More work is also needed to understand the specific details of how the two sinks (Poynting vs. heat) for the free energy actually drive the explosion. For example, a key issue for any transient dynamo outflow model is determining the dynamical formation of a magnetic corona where the magnetic field can dominate the dynamics, and how the large scale field opens up there to mediate the outflows. Improved 3-D versions of self-consistent numerical studies that include both the amplification of a very weak seed field and the outflow generation (Matt et al.\\ 2004; Moiseenko, Bisnovatyi-Kogan, \\& Ardeljan 2004) warrant development. Future  work specifically on transient dynamo theory should also include self-consistent treatments of buoyancy losses (e.g. Cline, Brummell, \\& Cattaneo 2003) radial and latitudinal spatial dependence of the system, latitudinal differential rotation, dynamical treatments of the backreaction incorporating magnetic helicity conservation (e.g. Blackman \\& Field 2002;  Brandenburg \\& Sandin 2004). Distinguishing the relative roles of the MRI, convection driven dynamos, and possibly shear-current effect dynamos (Rogachevskii \\& Kleeorin 2003) also warrants further work. The extent to which the MRI might require latitudinal, rather than radial shear in the pressure supported interior of a star is of particular interest.  The MRI in such environments has  thus far been studied only in the linear regime (Balbus \\& Hawley 1994).  \\ni {\\bf Acknowledgments:} EGB acknowledges support from NSF grant AST-0406799, NASA grant ATP04-0000-0016, and the KITP of UCSB, where this work was supported in part by NSF Grant PHY-9907949. JTN acknowledges a Horton Fellowship from the Laboratory for Laser Energetics.  JHT was supported in part by a Senior Visiting Fellowship at the Isaac Newton Institute for Mathematical Sciences, University of Cambridge, under EPSRC Grant N09176. Thanks to the referee for inducing our inclusion of the nonlinear shear evolution, and Lifin Wang for comments."
    },
    "0410/astro-ph0410049_arXiv.txt": {
        "abstract": "In a previous paper, we described a new method for including detailed information about substructure in semi-analytic models of halo formation based on merger trees. In this paper, we compare the predictions of our model with results from self-consistent numerical simulations. We find that in general the two methods agree extremely well, particularly once numerical effects and selection effects in the choice of haloes are taken into account. As expected from the original analyses of the simulations, we see some evidence for artificial overmerging in the innermost regions of the simulated haloes, either because substructure is being disrupted artificially or because the group-finding algorithms used to identify substructure are not detecting all the bound clumps in the highest-density regions. Our analytic results suggest that greater mass and force resolution may be required before numerical overmerging becomes negligible in all current applications. We discuss the implications of this result for observational and experimental tests of halo substructure, such as the analysis of discrepant magnification ratios in strongly lensed systems, terrestrial experiments to detect dark matter particles directly, or indirect detection experiments searching for positrons, gamma-rays, neutrinos or other dark matter decay products. ",
        "introduction": "\\label{sec:1}  There is now very strong evidence from observations of the microwave background (Spergel et al.\\ 2003), galaxy redshift surveys (e.g.\\ Tegmark et al.\\ 2004), weak lensing measurements (e.g.\\ Rhodes et al.\\ 2004), and modelling of the Lyman-$\\alpha$ forest (e.g.\\ Kim et al.\\ 2004), that most of the matter in the universe is non-baryonic dark matter, and that the power spectrum of density fluctuations in this component extends to subgalactic scales, as expected in `cold' dark matter (CDM) models. The implications of the CDM power spectrum for structure formation are well established. Dark matter haloes, the dense regions that surround galaxies, groups and clusters, form from the bottom up, through the merging of progressively larger structures. This process of hierarchical merging has been studied extensively, and the overall properties of galaxy or cluster haloes formed in this way are now fairly well determined.  To learn more about dark matter, and to search for features in the power spectrum that could reveal new phases in the evolution of the very early universe, we must push the theory of structure formation to smaller scales. Most of our current understanding of the properties of dark matter on subgalactic scales comes from numerical simulations of structure formation. These simulations have been used to determine the evolution of large-scale structure and the formation of CDM haloes on scales ranging from the current horizon (Kauffmann et al.\\ 1999) down to the local neighbourhood (e.g\\ Mathis et al.\\ 2002). Furthermore, by selectively re-simulating sections of a large volume at higher resolution, recent studies have been able to `zoom in' on single objects, resolving the substructure within individual haloes in exquisite detail (e.g.\\ recent work by De Lucia et al.\\ 2004; Gill, Knebe, \\& Gibson 2004a; Gill et al.\\ 2004b; Gao et al.\\ 2004a, 2004b; Diemand et al.\\ 2004c; Weller, Ostriker \\& Bode 2004; Reed et al.\\ 2004)  There is a hard limit, however, to the dynamic range that can be achieved using  this approach of selective re-simulation. Structure formation mixes information  on many different scales as haloes form. To model the formation of a dark matter halo accurately, one needs to include the effects of very long-wavelength fluctuations as well as the smaller fluctuations that produce substructure. The minimum scale that can be included in any self-consistent  simulation of the formation of a present-day halo is determined by the requirement that the largest fluctuations in the volume studied still be  in the linear regime at the present day, and by the finite numerical  resolution available computationally. For the highest-resolution simulations  that are currently feasible, this leads to a minimum mass scale for resolved  substructure of around 10$^{-4}$--10$^{-5}$ of the mass of the main halo  considered. To study halo substructure below this mass limit requires analytic  or semi-analytic extensions to the numerical results. It is precisely this sort  of small-scale information, however, that is required in many current  applications including galaxy dynamics, strong lensing, direct or indirect  dark matter detection, or tests of dark matter physics in general.  In earlier work (Taylor \\& Babul 2001, TB01 hereafter), we developed a model for dynamical evolution of satellites orbiting in the potential of larger system. This model includes simple treatments of dynamical friction, tidal mass loss and tidal disruption.  It calculates satellite evolution over a many short timesteps, rather like a restricted $N$-body simulation, but uses only global properties of the satellite to determine its evolution, thus reducing the computational expense considerably.  More recently (Taylor \\& Babul 2004a, paper I hereafter), we have applied  this model of satellite evolution to the merging subcomponents involved in  the hierarchical formation of galaxy, group or cluster haloes, creating a  full semi-analytic model of halo formation. In a second paper  (Taylor \\& Babul 2004b, paper II hereafter), we presented the basic predictions  of this model, including distributions of subhalo mass, circular velocity,  location and merger epoch, and the correlations between these properties.  We found results similar to those of recent numerical studies, as well as for a few systematic differences.  In this paper, we compare the predictions of the semi-analytic model directly with the results of self-consistent numerical simulations of halo formation. This comparison is particularly interesting, since the only free parameters in the semi-analytic model were fixed in paper I, either by matching restricted simulations of individual subhaloes (to fix the parameters of the dynamical model), or by assuming self-similarity in the merging process (to fix the one free parameter in the pruning method). Thus we have no remaining parametric freedom when comparing our results to self-consistent simulations, making the comparison a meaningful one.  Overall, we will show that there is reasonable agreement between the  semi-analytic  and numerical results, particularly in regions where both are expected to be accurate, but also that there are systematic differences between their predictions. These could reflect inaccuracies in the semi-analytic model,  but closer examination of the numerical results suggests that at least part of the discrepancy is due to artificial numerical effects in the simulations. The quantitative estimate of the magnitude of these effects has interesting implications for the analysis of several recent observational results.  The outline of this paper is as follows. In section \\ref{sec:2}, we summarise briefly the semi-analytic model developed in paper I. In section \\ref{sec:3}, we describe the six simulated haloes used in our comparison, and analyse the general properties of their subhalo populations. In section \\ref{sec:4}, we compare the properties of individual subhaloes, as well as the cumulative distributions of subhalo mass or circular velocity, in semi-analytic model and in the numerical simulations. In particular, we examine the evidence that the central regions of the simulated haloes are subject to artificial overmerging. In section \\ref{sec:5}, we consider the implications of overmerging in two particular areas, the modelling of strongly-lensed systems, and the analysis of direct detection experiments. We summarise our conclusions in section \\ref{sec:6}. Finally, we note that as in papers I and II, in this paper we will generally consider results for the former `standard' CDM (SCDM) cosmology with $h = 0.5$ and $\\sigma_8 = 0.7$, because the simulations we compare to assumed this cosmology. In general, our main results depend only weakly on cosmology, as discussed in paper II. ",
        "conclusions": "\\label{sec:6}  In this paper, we have compared the properties of halo substructure predicted by a semi-analytic model with the substructure identified in a set of self-consistent numerical simulations of halo formation. The semi-analytic model combines merger trees, an algorithm for treating higher-order substructure, and an analytic description of satellite dynamics. While the original treatment of satellite dynamics had several free parameters (TB01), in our full model of halo formation these were fixed by comparison with the high-resolution, restricted simulations of Velazquez and White (1999) and H03, as described in paper I. Thus we have no remaining parametric freedom when comparing the predictions of the semi-analytic model to the numerical results. Our model does make a number of assumptions and approximations, however, concerning the shape and spherical symmetry of the halo density profile, for instance, as well as halo concentrations and subhalo orbits. As discussed in paper II, we estimate that modifying these assumptions would change our results at the 20--30 percent level.  Despite its uncertainties and simplifications, without any adjustment of the parameters the semi-analytic model does an excellent job of matching the numerical results in the outer regions of haloes, where the latter are the most robust. In particular, it matches the overall distribution of subhalo properties and the amplitude of the cumulative distributions of subhalo mass or peak velocity to within 10--20 percent. This is both the level of accuracy expected of our dynamical model, and is also comparable to the intrinsic scatter from one halo to the next, so we conclude that the two methods agree more or less exactly in this regime.  In contrast to this, in the central regions of haloes the semi-analytic model predicts substantially more substructure than is seen in the simulations. The excess subhaloes are predicted to be ancient, dense systems which have orbited in the central part of the halo for most of the age of the universe (10--12 Gyr, or since a redshift of $z = 2$). In the semi-analytic model, these systems survive because the overall heating and disruption rates are lower than those measured in cosmological simulations. Whether these ancient systems should survive in reality is unclear. The disagreement with semi-analytic and numerical predictions could indicate that our dynamical model for heating and mass loss is systematically less accurate. This seems unlikely, however, since the model is calibrated on the simpler, higher-resolution simulations of Velazquez and White (1999) and H03, and since many of the central subhaloes retain 20--30 percent of their original mass, and thus are at an evolutionary stage where the analytic mass-loss model matches the restricted simulations quite closely.  One physical process recently proposed to explain the higher disruption rate seen in simulations is the increased efficiency of mass loss in systems with anisotropic (internal) velocity distributions (Kazantzidis et al.\\ 2004). Our model is calibrated using simulations of isotropic systems, so in principle anisotropy could affect our results. In the example they consider, however, Kazantzidis et al.\\ find that both fairly strong anisotropy and substantial (80--90 percent) mass loss are required before the evolution of the satellite changes substantially. Further work should clarify the importance of this effect.  The other possibility is that the semi-analytic predictions are essentially correct, and that the simulations analysed in this work are affected by residual overmerging, or by problems with the group finder used to analyse their structure. If overmerging is responsible for the patterns discussed in section \\ref{sec:4}, higher-resolution numerical work will be required to establish the true level of substructure in the centres of CDM haloes definitively. In particular, based on the arguments of section \\ref{subsec:4.5}, the properties of substructure in current simulations may be unreliable within the central 10 percent of the virial radius (or 30\\,kpc for a system like the Milky Way), and it may take an increase of $\\sim$100 or more in mass resolution to get convergent results down to the equivalent of the solar radius. A similar increase in resolution would be required to increase the relaxation time in a subhalo at the nominal resolution limit of current simulations ($\\sim 10^{-5} M_{\\rm vir,m}$), and that formed at $z = 2$, until it was longer than the Hubble time. Thus while future numerical work can eventually resolve this issue definitively, achieving the required mass and force resolution will remain challenging for some time.  Whatever the final answer to the problem, we have illustrated through several examples that the survival of substructure in the innermost parts of haloes is extremely important to the analysis of many recent observational and experimental results. As it stands, our semi-analytic model provides a robust and computationally efficient basis for studying a wide range of problems related to halo structure and substructure, including the origin and evolution of galaxy morphology, tidal disruption of dwarf galaxies and globular clusters, direct detection of dark matter particles and the local phase-space density of dark matter, indirect detection of dark matter decay products such as gamma-rays and positrons, and the analysis of strong-lensing systems. We will explore these topics in future papers."
    },
    "0410/astro-ph0410563_arXiv.txt": {
        "abstract": "A new method for absolute polarimetric calibration of large telescopes is presented. The proposed method is highly accurate and is based on the calibration of a small sub-aperture, which is then extended to the full system by means of actual observations of an astronomical source. The calibration procedure is described in detail along with numerical simulations that explore its robustness and accuracy. The advantages and disadvantages of this technique with respect to other possible alternatives are discussed. ",
        "introduction": "\\label{sec:intro}  Polarimetry has become a fundamental part of solar physics and is rapidly gaining importance in many other fields of astrophysics\\cite{WS02}. One of the most challenging problems that one faces when doing astronomical polarimetry is the determination of the instrumental polarization introduced by the telescope in the observed beam. There are two approaches that have been used to deal with this problem. The first one is to place calibration optics at the telescope entrance, which allows one to feed light in a known state of polarization into the system. The measured Stokes parameters can then be used to determine the telescope response with a very high degree of accuracy. Two solar telescopes currently use this strategy: The Dunn Solar Telescope (DST), at the Sacramento Peak observatory (NM, USA) operated by the National Solar Observatories, and the German Vacuum Tower Telescope (VTT) at the Observatorio del Teide of the Instituto de Astrof\\' \\i sica de Canarias (Spain). The Swedish Solar Tower at the Observatorio del Roque de Los Muchachos (also of the Instituto de Astrof\\' \\i sica de Canarias) will soon have this capability, as well.  The second approach to determine the telescope polarimetric response is to observe relatively well-known astronomical sources of polarization. This is sometimes done in night-time astronomy by observing ``standard stars''. In solar physics, this may be accomplished with observations of a sunspot. The procedure\\cite{S04} considers Milne-Eddington fits (MEF) to sunspot umbra spectra, from which the elements of the telescope Mueller matrix are adjusted to minimize the residuals of the MEF to the observed data.  It is generally desirable to have a purely instrumental telescope calibration, so that no {\\it a priori} assumptions on the observed targets are needed. Unfortunately, the construction of calibration optics for a large aperture (larger than $\\sim$1~m) is impractical. This had not been an issue in the past because existing solar telescopes have apertures that do not exceed 1~m in diameter, whereas larger night-time telescopes did not require absolute polarimetric accuracies higher than what can be routinely achieved by means of standard stars. This situation, however, is rapidly changing due to the ever increasing accuracy demands of astronomical polarimetry and the development of new large-aperture solar telescopes, like the German telescope GREGOR\\cite{VvdLK+03} or the Advanced Technology Solar Telescope\\cite{KRH+02} \\cite{KRK+03} (ATST).  A previous paper\\cite{SN04c} introduced the ``beam-expansion'' method (BEM) for the instrumental calibration of a large telescope. This work presents a different procedure, the sub-aperture method, which has some interesting advantages over the BEM (at least for the particular case of the ATST design). We shall hereafter consider the problem of calibrating the ATST, but the proposed procedure may be easily applicable to other telescopes as well.  The sub-aperture method is based on the following basic idea. One starts by effectively turning the ATST into a small-aperture telescope. This is done by inserting a mask into the beam at any pupil image. The telescope sub-aperture is calibrated as usual by means of calibration polarization optics above the primary mirror. An astronomical polarization source (e.g., a sunspot) is observed through the sub-aperture and then again with the full aperture. The two observations are only a few seconds apart in time. The sub-aperture calibration is used to remove instrumental polarization in the first observation, which allows one to determine the actual Stokes vectors emerging from the astronomical source. Once these input vectors are known, the full aperture observation can be used to calibrate the telescope. ",
        "conclusions": "The sub-aperture method proposed in this work provides a suitable means of calibrating a large telescope. This paper explores its possibilities and limitations by means of hound-and-hares tests. Potential advantages and disadvantages with respect to two other alternatives (the BEM and MEF) are discussed.  The simulations indicate that the sub-aperture method is robust enough to meet the stringent ATST polarimetric requirements in practical situations. The amount of time required for the calibration procedure is probably inferior to $\\sim$20~minutes. Most of the time is spent observing the various configurations of the calibration optics, which is the same procedure employed at the DST (Sacramento Peak) or the VTT (Tenerife) for the polarimeter calibration. This is done currently on a routine basis about once per day during the observing runs."
    },
    "0410/astro-ph0410425_arXiv.txt": {
        "abstract": "We present Gemini integral-field unit (IFU) spectroscopy and {\\sl Hubble Space Telescope} (HST) F435W- and F814W-band images of a newly discovered four-image gravitational lens, SDSSJ140228.22+632133.3 (hereafter SDSSJ1402). The system was identified as one of 49 gravitational-lens candidates in the luminous red galaxy sample of the Sloan Digital Sky Survey, based on higher-redshift emission lines in the spectra of the lower-redshift galaxies. We are imaging the most promising lens candidates with HST as part of a Snapshot program designed to expand the sample of known gravitational lenses amenable to detailed photometric, lensing and dynamical studies; SDSSJ1402 was the first of our targets to be observed with the ACS-WFC on board HST. The lens is a smooth elliptical galaxy at a redshift of $z_{\\rm l}=0.2046 \\pm 0.0001$ with a Sloan $r$-band magnitude of $17.00 \\pm 0.05$ and a stellar velocity dispersion of 267 $\\pm$ 17 km s$^{-1}$, obtained from its SDSS spectrum. Multiple emission lines place the quadruply-imaged source at a redshift of $z_{\\rm s}=0.4814 \\pm 0.0001$. The best-fitting singular isothermal ellipsoid lens model gives an Einstein radius $b=1\\farcs35 \\pm 0\\farcs05$ (or $[4.9 \\pm 0.2]h_{65}^{-1}$ kpc), corresponding to a total mass of $(30.9 \\pm 2.3)\\times 10^{10}h_{65}^{-1}$~${\\rm M}_{\\sun}$ within the critical curve. In combination with HST photometry this gives a rest-frame $B$-band mass-to-light ratio of $(8.1 \\pm 0.7)h_{65}$ times solar within the same region. The lens model predicts a luminosity-weighted stellar dispersion within the $3 \\arcsec$-diameter SDSS aperture of $\\sigma_*\\approx 270$~ km s$^{-1}$, in good agreement with the observed value. Using the model to {\\sl de-lens} the four lensed images yields a source with a smooth, monotonically-decreasing brightness distribution. Taken in combination, the HST ACS images, Gemini IFU spectroscopy, and self-consistent mass model show SDSSJ1402 to be a genuine lens system. ",
        "introduction": "In the last decade, galaxy-size gravitational lenses have become an increasingly important tool for the study of cosmology and galaxy evolution. The number of known lenses has now reached almost 100 (see the CASTLES web page at {\\tt http://cfa-www.harvard.edu/castles/}), and subsets of lenses with suitable properties are now available for a variety of applications: the determination of the cosmological parameters from lens statistics \\citep[e.g.][]{turner_lensstat, kochanek_lensstat_96, chae_lensstat_03}, the measurement of the Hubble Constant from lens time delays \\citep[e.g.][]{kundic_h0, schechter_1115_delay, koopmans_1608_h0}, and the study of the mass distribution of E/S0 galaxies and their dark matter halos outside the local Universe \\citep[e.g.][]{kochanek_dyn_halo_94, rusin_lensstat, tk04}.  Despite the great progress, the small subsample sizes of suitable lenses is a major limitation. This is particularly true for the study of the properties of lens galaxies (typically E/S0s). Most currently known lenses have been discovered as bright quasars that show multiple images when observed at high spatial resolution. For optical quasar lenses, this biases the sample in favor of systems where the lensed source outshines the lens galaxy. Radio surveys such as the Cosmic Lens All-Sky Survey \\citep[CLASS;][]{class_myers, class_browne} are less susceptible to this bias, but redshifts in radio lens systems can be difficult to obtain, and the lensing galaxies are often faint. This situation makes it very difficult to obtain high-quality photometry, redshifts, and internal kinematics of lens galaxies, which are essential ingredients for detailed modeling. It is not a coincidence that 3/5 lenses analyzed by the Lenses Structure and Dynamics Survey (LSD; \\citealt{kt02, kt03, tk02, tk03, tk04}; hereafter collectively KT) so far have been discovered serendipitously in Hubble Space Telescope (HST) images.  To overcome this limitation we have initiated the Sloan Lens ACS (SLACS) Survey, a new survey for gravitational lenses that exploits the Sloan Digital Sky Survey (SDSS) archive and the angular resolution of HST to find and confirm new gravitational lenses. First, a large sample of SDSS early-type galaxy spectra is searched for emission lines at a higher redshift than that of the absorption features (\\citealt{bolton_speclens}; see also \\citealt{warren_0047_96}). By requiring that the foreground galaxy dominate the absorption spectrum, this procedure will tend to select bright lenses with faint background sources. The second step consists of follow-up snapshot observations of the best candidates with the Advanced Camera for Survey (ACS) to confirm the incidence of lensing and provide the lens geometry necessary to construct a mass model\\footnote{Our survey is the second HST snapshot lens survey.  For details of the first, see \\citet{maoz_snapshot}.}. Each candidate is observed for 420 seconds in each of F435W and F814W. The F435W filter is selected to optimize the detection of the lensed source galaxy (expected to be blue and starforming) while F814W is selected to optimize the signal-to-noise ratio on the (red) lens galaxy. Color information also provides important evidence in support of the lensing hypothesis. In this {\\it Letter} we present the first newly discovered lens, along with ACS images and Gemini GMOS-North integral-field spectroscopy and a lens model, and discuss its overall properties. We assume that the Hubble constant, the matter density, and the cosmological constant are H$_0=65\\,h_{65}$~km\\,s$^{-1}$\\,Mpc$^{-1}$ (with $h_{65}=1$), $\\Omega_{\\rm m}=0.3$, and $\\Omega_{\\Lambda}=0.7$, respectively. All spectroscopic wavelengths are given in vacuum values. ",
        "conclusions": "Based on the following evidence, we conclude that SDSSJ1402 is a lens system:  \\begin{enumerate} \\item The discovery of emission-lines in the SDSS spectrum at a redshift significantly larger than that of the high-$\\sigma_{*}$ LRG and within a 1.5$''$ radius from the galaxy centroid, and confirmed by follow-up GMOS-N IFU spectroscopy. \\item The discovery of four images (with two tangentially stretched) in a typical quad-lens-like configuration around the galaxy in {\\sl Hubble Space Telescope} images, with similar colors and positions between F814W and F435W filters. \\item The spatial coincidence of these images with the higher-redshift line emission, as shown by GMOS-N IFU spectroscopy. \\item The excellent goodness-of-fit ($\\chi^2/{\\rm NDF}\\approx 0.7$) for the SIE mass model (no external shear). The stellar velocity dispersion predicted by this mass model and the observed luminosity density is in excellent agreement with the SDSS value, and the galaxy and lens-model ellipticities and position angles are in significant agreement with one another. \\item A direct inversion of the lensed image---based solely on the SIE mass model determined from the image positions---leads to a very simple compact source structure in both F814W and F435W filters at the same position near the inside of the fold caustic. \\end{enumerate}  Hence, despite the relatively low signal-to-noise of the lensed images, the spectral and imaging data of SDSSJ1402 presented in this paper, in conjunction with the self-consistent lens model, convincingly show that the {\\sl first} candidate observed by the SLACS Survey is also the {\\sl first} genuine lens system from the program! The lens geometry and HST photometry yield a (largely model-independent) measurement of $(30.9 \\pm 2.3)\\times 10^{10}h_{65}^{-1}$~M$_\\odot$ and a rest-frame $B$-band mass-to-light ratio of $(8.1 \\pm 0.7)h_{65}$ times solar within the $(4.9 \\pm 0.2)h_{65}^{-1}$-kpc cylinder enclosed by the critical curve. Spatially resolved kinematic observations of this system will permit more detailed lens/dynamical modeling, which will constrain the radial mass profile and the relative fraction of luminous to dark matter in the central region of the lensing galaxy.\\\\"
    },
    "0410/astro-ph0410339_arXiv.txt": {
        "abstract": "We present a description of deep radio, optical, and near IR observations taken with the VLA, the KPNO 2m and the KPNO 4m of the region containing the rich cluster of galaxies Abell 2125. The reduction of each dataset is described. A catalog of radio sources apparently not associated with members of Abell 2125 and the associated R-band magnitudes is presented. ",
        "introduction": "Deep radio observations offer one of the important windows on the evolution of star formation and black-hole-related activity as a function of cosmological epoch. Radio, optical/NIR, FIR and X-ray data taken together have the potential eventually to give us a well-constrained picture of our history. Each window, by itself, has serious weaknesses. The optical/NIR window is the best studied and provides the most direct view. However, it suffers from uncertainties due to dust obscuration, a fundamental part of star-formation. FIR observations give us a signal from the dust but leave uncertain the origin of heating of the dust and fail to see into the heart of the AGN or star-forming regions. In X-rays and radio we can see behind the dust but we are measuring  secondary emission which is not easy to relate physically to the black hole physics or star-formation-rate (SFR). Thus we need a combination of all this information to deduce the full picture.  Below $L_{20cm} \\sim 10^{23}$ W Hz$^{-1}$, one can show that locally the radio luminosity function is dominated by synchrotron emission related to star-formation driven processes \\citep[e.g.,][]{con92}. Above this luminosity most sources are twin-jet, black-hole-driven sources. Below the  $10^{23}$ W Hz$^{-1}$ break point, the sources are mostly found in spirals rather than ellipticals and the radio and FIR luminosities correlate very well \\citep[e.g.,][]{yun01}. Through the radio-FIR correlation, one can estimate the SFR for the radio alone and can use the ratio of radio to submillimeter emission to constrain the redshift (assuming the empirical radio-FIR relation continues to hold at any redshift) \\citep{cy99}. Thus deep radio observations with limiting flux densities $<<1$ mJy, offer a way to study star formation in the distant universe.  However, there are a lot of assumptions in the previous paragraph. What one really needs is deep data in all four windows to understand the picture. Where we have such information new puzzles arise. For example, SCUBA submillimeter sources are usually associated with faint radio sources. Using the superior radio positions, one can make faint optical identifications and try to measure redshifts. Often when an optical redshift can be obtained, one finds an AGN spectrum \\citep[e.g.,][]{led02}. Even so, the redshift estimated from the Carilli/Yun relation agrees with the measured value. Perhaps this is a coincidence but it might suggest that bright black-hole-driven AGN's are part of the star-formation process in galaxies.  We are interested both in the very distant star-formation/AGN history as well as the same phenomena in intermediate-distance rich clusters($0.2 \\ls z \\ls 0.4$). A deep radio survey in the direction on such a cluster can be used for both purposes. This paper reports deep radio observations in the direction of Abell 2125. Future papers will report the combination of these data with observations in other bands to study both problems.  We began studying the Abell 2125 field with medium deep, 20cm, VLA C-array observations \\citep{dow,o99}. We compared Abell 2125 ( a richness class 4, blue cluster and $z \\sim 0.25$) with Abell 2645 (an apparently similar cluster at the same redshift but with  much redder galaxies). We found a much higher detection rate of radio galaxies (27 vs 4) in Abell 2125. The detection rate in A2645 is consistent with a normal AGN population in a lower redshift cluster. The A2125 excess population occurs entirely with luminosities below  $L_{20cm} \\sim 10^{23}$ W Hz$^{-1}$ and thus seems consistent with a star-forming population. However, only a small fraction of the excess population have optical spectra consistent with enough star-formation activity to explain the population entirely that way. Either the star-forming activity is very well hidden in the majority of the low radio luminosity sample or the some other explanation is necessary for the radio activity.  One possible clue is that Abell 2125 does not appear to be a typical very rich cluster when examined in more detail in the optical and X-ray bands. Optically the cluster consists of a central concentration together with a extension to the southwest extending at least 2 Mpc. X-ray observations \\citep{wang} show the same pattern and that the total X-ray luminosity is low for such a rich cluster. Much of the excess X-ray population is contained in the southwest extension, not the cluster core. Thus perhaps these unusual properties have something to do with the radio excess.  To study the cluster further, we have made a much deeper radio observations and higher resolution using all four VLA configurations. The deep radio data has also motivated deep optical and NIR imaging, optical spectroscopy \\citep{mil04}, deep millimeter/submillimeter observations \\citep{e03}, and a deep {\\it Chandra} exposure \\citep{w03}.In some cases, these ancillary observations were not motivated by the cluster but by the possibility of studying the background sample ({\\it e.g.} the submillimeter survey). In this paper we document the techniques used in reducing the radio and optical imaging data for this field. We also report the radio data and R-band optical identifications for these data. A related paper contains the results for the radio sources confirmed to be associated with cluster members. ",
        "conclusions": "We have presented the data reduction techniques and the catalog of radio detections of non-cluster members for the A2125 field. We also report the R-band optical identifications. We report the catalog of the detections for confirmed cluster members in a companion paper, as well as a more detailed interpretation of these results. Subsequent papers will report more detailed astrophysical results for this dataset. \\clearpage"
    },
    "0410/astro-ph0410613_arXiv.txt": {
        "abstract": "In order not to be in conflict with observations it is crucial that extra dimensions, if they exist, are stable. It is shown that in the context of homogeneous extra dimensions, this can easily be achieved during both (4D) radiation and vacuum energy dominated eras of the cosmological evolution. During matter domination, however, there is no such possibility even for a very general class of stabilization mechanisms. Even if extra dimensions could be stabilized during matter domination, it is argued that they are generically time-varying during any transition period, such as the one from radiation to matter domination. ",
        "introduction": "The idea that our world may consist of more than four space-time dimensions, with extra spatial dimensions compactified on some small scale, was first proposed by Nordstr\\\"om, Kaluza and Klein \\cite{nor}. One of the main motivations for these and subsequent works was the hope for a possible unification of all interactions (see e.g.~\\cite{bai87} for a review on Kaluza-Klein theories). Recent years have seen a great revival of interest in extra-dimensional scenarios, most notably due to the influence of string theory.  Cosmology provides an important testing ground for theories involving extra dimensions. For example, a generic prediction of these theories is that (some of) the fundamental coupling constants vary with the volume of the internal space. However, the strong cosmological constraints on the allowed variation of these 'constants' (see for example \\cite{oli}) require the extra space to be not only compactified, but also stabilized, at a time no later than big bang nucleosynthesis (BBN). One of the main tasks is therefore to find a dynamical explanation for this stabilization and to reproduce standard cosmology for late times.  Here, we will study the prospects for stabilizing homogeneous extra dimensions for a wide range of available stabilization mechanisms. We find that while the extra dimensions can easily be stabilized during both radiation and vacuum energy dominated periods of the cosmological evolution, one encounters generic problems during matter domination. For illustrative reasons, the analysis is done both for the Friedmann-like field equations that are obtained from the higher-dimensional action as well as for the dimensionally reduced and conformally transformed four-dimensional effective action. Of course, these two approaches are equivalent and lead to the same conclusions. The direct comparison, however, helps to clarify hidden subleties and leads to some insight into how one has to interpret the results. For previous work on the stabilization of homogenous extra dimensions see e.g.~\\cite{randj,frea} and references therein. ",
        "conclusions": "\\label{conc}  It is crucial for the viability of extra-dimensional scenarios that, for late times, the extra dimensions are stabilized and standard cosmology can be recovered. A wide range of possible stabilization mechanisms can be described in a phenomenological way by adding a 'radion potential' of the form (\\ref{sbg}) to the action. We have shown that the stabilization of homogeneous extra dimensions in this set-up can easily be achieved for the radiation dominated period of the cosmological evolution as well as during vacuum-energy domination. Surprisingly enough, however, there are generic problems during matter domination. By allowing for different types of stabilization mechanisms, for example of the form $W(b,\\rho)$, one might in principle get static solutions \\cite{bria} - though certainly at the cost of some serious fine-tuning (remember that is not enough to fix the size of the extra dimensions, one also has to reproduce the correct behaviour for the 4D scalefactor). As discussed in the last section, the size of the extra dimensions would in any case be expected to vary during transition periods.  Problems with the stabilization of extra dimensions even arise in certain brane world models once one allows for non-relativistic contributions to the matter content \\cite{arn}. It would be interesting to further investigate the various existing extra-dimensional models and see if and how they can avoid these problems, i.e.~whether they produce the correct late-time cosmology not only for radiation- but also for matter-domination.      \\begin{theacknowledgments} We would like to thank David Wiltshire for helpful discussions on conformal transformations and the organizers of Phi in the Sky for providing a nice and enjoyable atmosphere in Porto. \\end{theacknowledgments}"
    },
    "0410/astro-ph0410607_arXiv.txt": {
        "abstract": "After some post-natal cooling, a spinning, magnetized, canonical neutron-star (NS) has a core of superconducting protons, superfluid neutrons, and degenerate extreme relativistic electrons, all surrounded by a thin highly conducting solid crust. The quantum fluids are threaded by a dense array of quantized vortex-lines which can interact strongly with a denser and much less uniform one of quantized magnetic flux-tubes. The physics of such a core predicts the evolution of a NS's surface magnetic field and related phenomena as the star's spin changes.  Predictions include changes in NS magnetic dipole moments, anomalously small polar caps in millesecond pulsars, properties of two different families of spin-period ``glitches\", and spin-down ages much greater than true ages for some pulsars. Quantitative model-based estimates for all of these are given.  None are in conflict with observations. ",
        "introduction": "There is abundant observational evidence that the magnetic dipole-moment of a rapidly spinning neutron star (NS) evolves during a pulsar's lifetime. Most of the relevant data  comes from observed NS-periods ($P$) and spin-down rates ($\\dot P$). With simple plausible models of NS-magnetospheres the spin-down torque on a (presently) solitary NS depends only on the NS spin  $(\\Omega   = 2\\pi /P)$ and dipole moment $(\\mu   )$: $$I\\dot\\Omega \\sim - {\\mu^2\\Omega^3\\over c^3}\\eqno(1),$$ where $I$ is the NS moment of inertia. Fig. 1 shows  NS-surface-dipole magnetic fields, $B\\equiv \\mu R^{-3}$, over a large range of pulsar spin-periods. (All pulsars are assumed to have $I=10^{45}$ gcm$^2$  and radii $R=10^6$\tcm.)  Also shown is the typical evolution of $B$ with $P$ from the microphysics inside a canonical, rapidly spinning, strongly magnetized NS.  In the model considered here a NS consists mainly of a sea of superfluid neutrons together with less dense components of superconducting protons and very relativistic degenerate electrons. (Effects from possible, much smaller, central volumes of more exotic particles are ignored.) Surrounding this sea of quantum fluids is a thin solid conducting crust of thickness  $\\Delta \\sim 10^{-1}R$ in which the NS protons clump into conventional, but extremely neutron-rich, nuclei.  In such an object $B$ is expected to evolve mainly because of changes in  $\\Omega$. Fig. 1 shows evolutionary segments $a  \\rightarrow b \\rightarrow  c \\rightarrow  d \\rightarrow  e$ predicted from this canonical NS model. The slopes of $B(P)$ in the segments $a \\rightarrow  b$,  $b \\rightarrow  c$  are not significantly different from those determined directly from certain observations of $P, \\dot P$, and $\\ddot P$ (Sect. 4), and from inferred ages of radiopulsars observed in other $B-P$ regions. Fig. 1 segments $a^\\prime \\rightarrow  a$; $a^{\\prime\\prime}\\rightarrow  b-c$; $c \\rightarrow  d$; $d \\rightarrow  h\\rightarrow   g$; $d \\rightarrow  e \\rightarrow  f$;  $d\\rightarrow   h  \\rightarrow q$  give more explicit $B(P)$ predictions than what can be inferred directly from $B-P$ data, but can be compared with other kinds of observations (Sects. 3, 7). How this same model leads to  two different kinds of ``glitches\" in  pulsars is discussed in Sect. 8.  In all comparisons of observations with model-based predictions no disagreements have yet become apparent.     \\begin{figure} \\centerline{\\includegraphics[width=3.8truein]{RUDERMANFIG1.ps}} \\caption{Dipole fields (B) of pulsars inferred from observed  $P,\\dot P$, together with the $B(P)$ predicted for a typical pulsar from the model discussed in the text. The point $c^\\prime$ is about where coherent radioemission is no longer observable. Point $d$ is for $P$ and $B$ of a solitary NS after $10^{10}$  yrs.  The millisecond pulsar (MSP) population in the lower left corner is presumed to be populated by the evolution of some NSs off the $bcd$ segment from accretion-induced spin-up by companions (many of which have been evaporated by the MSPs). After accretion stops the MSPs again spin-down like solitary radiopulsars. Open circles are pulsars in binaries which are not presently interacting with their companions. (The $P$, $\\dot P$ of many more radiopulsars are shown in reference [28].) The diagonal dashed line is the canonical accretion-driven spin-up line.}  \\end{figure}   In almost all biographies, certainly including those of NS $B$-fields, authors and readers have good reason to be much less confident about descriptions of earliest life (conception, infancy, and childhood) than of adolescence, maturity, and old age.  This is even more the case here where there is not yet a consensus about what is happening in any one of these stages. We shall begin our account of the development of a NS magnetic field with this {\\it caveat} very much in mind. ",
        "conclusions": ""
    },
    "0410/astro-ph0410431_arXiv.txt": {
        "abstract": "The combination of the first-year Wilkinson Microwave Anisotropy Probe (WMAP) data with other finer scale cosmic microwave background (CMB) experiments (CBI and ACBAR) and two structure formation measurements (2dFGRS and Lyman $\\alpha$ forest) suggest a $\\Lambda$CDM cosmological model with a running spectral power index of primordial density fluctuations. Motivated by this new result on the index of primordial power spectrum, we present the first study on the predicted lensing probabilities of image separation in a spatially flat $\\Lambda$CDM model with a running spectral index (RSI-$\\Lambda$CDM model). It is shown that the RSI-$\\Lambda$CDM model suppress the predicted lensing probabilities on small splitting angles of less than about 4$^{''}$ compared with that of standard power-law $\\Lambda$CDM (PL-$\\Lambda$CDM) model. ",
        "introduction": "Mapping the mass distribution of matter in the universe has been a major challenge for modern observational cosmology. One of the direct procedures to weigh matter in the universe is measuring its deflection of light by gravity. The statistics of gravitational lensing can provide us with a very powerful probe of the mass distribution of the Universe. It is well known that the Jodrell-Bank VLA Astrometric Survey (JVAS) and the Cosmic Lens All-Sky Survey (CLASS) \\citep{2000IAUS..201E..47B,2003MNRAS.341....1M, 2003MNRAS.341...13B} have provided us with observations of strong lensing probabilities for small image separations ranging from $0.3''$ to $3''$. By comparing predicted lensing probabilities with observations, we can examine the mass distributions of dark matter halos, in particular, their inner density slopes \\citep{2004ApJ...600L...7H,2003MNRAS.344.1237R}.  Based on the Cold Dark Matter (CDM) model, which has become the standard theory of cosmic structure formation, the lensing probabilities strongly depend on the density profiles of CDM halos. The lensing model is usually described by a singular isothermal sphere (SIS), the Navarro-Frenk-White (NFW) model \\citep{1996ApJ...462..563N,1997ApJ...490..493N}, or generalized NFW (GNFW) density profiles of dark halos \\citep{1996MNRAS.278..488Z}. \\cite{2002ApJ...566..652L} employed a semi-analytical approach to analyze the gravitational lensing of remote quasars by foreground dark halos and checked the plausibility of various lensing models. They found that no model can completely explain the current observations: the SIS models predict too many lenses with large splitting angles, while the NFW models predict too few small splitting angles. They therefore further developed a two-population halo model for lensing: small mass halos with a steep inner density slope and large mass halos with a shallow inner density slope, concluding that a combination of SIS and NFW halo models can reproduce the current observations reasonably well. Unlike previous work that directly models the density profiles of dark matter halos semi-analytically, \\cite{2004ApJ...602L...5Z} generalized the density profiles of dark matter halos from high-resolution N-body simulations by means of generalized Navarro-Frenk-White (GNFW) models of three populations with slopes, $\\alpha$, of about -1.5, -1.3 and -1.1 for galaxies, groups and clusters, respectively. He presented the calculations of lensing probabilities using these GNFW profiles for three populations in various spatially flat cosmological models with a cosmological constant $\\Lambda$. He showed that the compound model of density profiles does not match well with the lensing probabilities derived from the combined data of JVAS/CLASS. Recently, \\cite{2004astro.ph..7298M} compared predictions on small scale structure for the $\\Lambda$CDM model by numerical simulations with observed flux ratios, and found that the disagreements between monochromatic flux ratios and simple lens models can be explained without any substructure in the dark matter halos of primary lenses. However, spectroscopic lensing observations of Q2237+0305 require more small mass dark halos than that expected in the $\\Lambda$CDM model.  In particular, a power spectrum of primordial fluctuation, $P_p(k)$, should be assumed in advance in the calculation of lensing probabilities. Inflationary models predict a approximately scale-invariant power spectra for primordial density (scalar metric) fluctuation, $P_p(k)\\propto k^n$ with index $n=1$ \\citep{1982PhRvL..49.1110G,1983PhRvD..28..679B}. The combination of the first-year Wilkinson Microwave Anisotropy Probe (WMAP) data with other finer scale cosmic background (CMB) experiments (Cosmic Background Imager [CBI], Arcminute Cosmology Bolometer Array Receiver [ACBAR]) and two observations of large-scale structure (the Anglo-Australian Telescope Two-Degree Field Galaxy Redshift Survey [2dFGRS] and Lyman $\\alpha$ forest) favor a $\\Lambda$CDM cosmological model with a running index of the primordial power spectrum (RSI-$\\Lambda$CDM), while the WMAP data alone still suggest a best-fit standard power-law $\\Lambda$CDM model with the spectral index of $n\\approx 1$ (PL-$\\Lambda$CDM) \\citep{2003ApJS..148..175S,2003ApJS..148..213P}. However, there still exist the intriguing discrepancies between theoretical predictions and observations on both the largest and smallest scales. While the emergence of a running spectral index may improve problems on small scales, there remain a possible discrepancy on the largest angular scales. It is particularly noted that the running spectral index model suppress significantly the power amplitude of fluctuations on small scales \\citep{2003ApJS..148..175S,2003ApJ...598...73Y}. This imply a reduction of the amount of substructure within galactic halos \\citep{2002PhRvD..66d3003Z}. \\cite{2003ApJ...598...73Y} studied early structure formation in a RSI-$\\Lambda$CDM universe using high-resolution cosmological N-body/hydrodynamic simulations. They showed that the reduced small-scale power in the RSI-$\\Lambda$CDM model causes a considerable delay in the formation epoch of low-mass minihalos ($\\sim10^6 M_{\\sun}$) compared with the PL-$\\Lambda$CDM model, although early structure still forms hierarchically in the RSI-$\\Lambda$CDM model. Thus the running index probably affect the strong lensing process of distant sources by intervening dark matter halos because the lensing probabilities strongly depend on the abundance of dark halos formed in the evolution of the universe. In this letter, we will present the first calculation of lensing probabilities in a RSI-$\\Lambda$CDM model to explore the effect of running spectral index of primordial fluctuation on strong lensing probabilities. We adopt a spatially flat $\\Lambda$CDM cosmological model through this letter.  The reminder of this paper is organized as follows. We describe strong lensing probabilities and mass function of dark halos in Section 2 and 3 respectively. The cross-sections for producing multiple images are described in Section 4. Our results is shown in Section 5, while a conclusion and discussion are given in Section 6. ",
        "conclusions": "Motivated by the new result on the index of primordial power spectrum from a combination of WMAP data with other finer scale CMB experiments and other large-scale structure observations, we present the first study on the predicted lensing probability in a $\\Lambda$CDM model with a running spectral index. In a popular cosmological model from the new fit results mentioned above: $h=0.71$, $\\Omega_{\\mathrm m}=0.27$, $\\sigma_8=0.84$, we calculate lensing probabilities with image separations greater than $\\Delta\\theta$ according to the compound model of SIS and NFW halo profile in both PL-$\\Lambda$CDM and RSI-$\\Lambda$CDM models. From the analysis above, we can see that the running spectral index mainly affect the predicted lensing probability on small image separation. It is well known that structures in the universe forms hierarchically in standard CDM models. \\cite{2003ApJ...598...73Y} found that although this hierarchical formation mechanism do not work well in RSI-$\\Lambda$CDM model compared with that in PL-$\\Lambda$CDM model and it also is not clear that the PS theory can be used in RSI-$\\Lambda$CDM model, the mass function measured by high-resolution cosmological N-body/hydrodynamic simulations overall match the PS mass function for both RSI-$\\Lambda$CDM and PL-$\\Lambda$CDM model.  In addition, because the running spectral index model predicts a significant lower power of density fluctuation on small scales than the standard PL-$\\Lambda$CDM model \\citep{2003ApJS..148..175S,2003ApJ...598...73Y}, it should also attract considerable attention in studies on weak lensing by large-scale structure \\citep{2003astro.ph..8446I}, especially on skewness  \\citep{2003ApJ...592..664P,2003ApJ...598..818Z} which characterizes the non-Gaussian property of $\\kappa$ field in the nonlinear regime.  \\begin{figure} \\epsscale{0.8} \\plotone{f2.eps} \\caption{The lensing probabilities for image separations greater than $\\Delta\\theta$. The solid-line histogram represents observed lensing probabilities from JVAS/CLASS. The solid line is the lensing probabilities for running spectral index $\\Lambda$CDM model, while the dashed line is that for power law $\\Lambda$CDM model.} \\label{f2eps} \\end{figure}  We are very grateful to the anonymous referee for constructive suggestion. T.J.Zhang would like to thank Ue-Li Pen, Peng-Jie Zhang, Xiang-Ping Wu, Bo Qin and CITA for their hospitality during his visits to the Canadian Institute for Theoretical Astrophysics(CITA), University of Toronto and the cosmology groups of the National Astronomical Observatories of P.R.China. This work was supported by the National Science Foundation of China (grants No.10473002 and 10273003)."
    },
    "0410/astro-ph0410741_arXiv.txt": {
        "abstract": "{We consider various classes of persistent extragalactic astrophysical sources which have been suggested in literature as possible emitters of ultra-high-energy cosmic rays (UHECR). We compare the strength of the claimed correlations by a uniform procedure for all classes of sources by making use of the AGASA, Yakutsk and HiRes stereo data. BL Lac type objects correlate with the cosmic rays detected by all three independent experiments and are more probably, compared to other astrophysical sources, related to the UHECR origin. With the account of the Galactic magnetic field (not possible for the HiRes data at the moment), apart of BL Lac type objects, unidentified gamma-ray sources may be correlated with AGASA and Yakutsk cosmic rays. } ",
        "introduction": "Thousands of cosmic rays with energies higher than $10^{19}$~eV have been detected by various experiments. Still, the origin of these energetic particles remains unknown. Deflections of charged particles by cosmic magnetic fields are relatively small (though significant) at these energies; moreover, presence of a fraction of neutral particles, which propagate rectilinearly, is not excluded. This opens a possibility for direct searches of astrophysical sources of the high-energy cosmic rays by positional correlations.  Several difficulties limit the application of this approach. Firstly, at least a large fraction of air showers with $E\\gtrsim 10^{19}$~eV is believed to be caused by protons, whose trajectories are bend by cosmic magnetic fields by several degrees. Lack of knowledge about these fields makes it difficult to account for the deflections while at highest energies, where the deflections are believed to be small, low statistics of the cosmic rays prevents one from making definite conclusions.  Secondly, the angular resolution of the cosmic-ray experiments is very poor by the astronomical standards. The most precise ground array, AGASA, has an average error of $1.8^\\circ$ in determination of the arrival directions at the highest observed energies, and $2.5^\\circ$ at $E\\sim 10^{19}$~eV~\\cite{AGASA:list}. As a result, an impressive number of astrophysical objects fall in the error box of each particular cosmic ray. The situation improved with the availability of the first set of the HiRes stereo data~\\cite{HiRes} with claimed angular resolution of $0.6^\\circ$. Still, this resolution is quite poor for direct identification: dozens of optical, radio and X-ray sources are contained in a given circle of $0.6^\\circ$ radius on the Celestial sphere.  With these complications, statistical methods give the only possible clue to search for positional correlations between the cosmic rays and their potential sources. An excess of cosmic rays at small angular distances (in particular, at those compared to the experiment's angular resolution) from objects of a certain class may suggest that this class of objects contains more probably the sources of these cosmic rays than other classes, which do not exhibit such an excess.\\footnote{The background of random coincidences is always present, and one cannot claim that one particular object is a source of a given cosmic ray based on these statistical arguments only.}  Many astrophysical sites have been suggested which can host an accelerator powerful enough to produce particles with $E\\sim 10^{19}$~eV and higher. Most of them are extragalactic (which is consistent with the lack of global galactic anisotropy of the arrival directions of high-energy cosmic rays). Various cosmic accelerators may operate by similar mechanisms, so the most efficient way to distinguish between them is based on direct positional correlations. Up to now, the correlations with different classes of objects were studied by different methods, with different data sets, or were not studied at all. This work aims at filling this gap and calculating the chance probabilities of the claimed correlations by a uniform method, which allows to compare the strength of correlations and to figure out more probable sources among those suggested previously.  The rest of the paper is organized as follows.  In Sec.~2, we discuss the logic of the statistical analysis of correlations, describe the methods we use and emphasize several subtle points. Sec.~3 lists potential astrophysical sources, briefly recalls motivations for these candidates and previous correlation studies. For each class of sources, we give here references for the catalogs we use in the current study. In Sec.~4, we present the results of the correlation studies for several representative samples of cosmic rays. The conjectures about positional correlations were formulated on the basis of the cosmic rays observed by AGASA, and sometimes by the Yakutsk experiment. We use the samples of cosmic rays with energies $4\\cdot 10^{19}<E<10^{20}$~eV observed by AGASA (49 rays) and by AGASA and Yakutsk (59 rays) to treat the suggested sources uniformly and to reformulate the conjectures. {We consider independently neutral and proton primaries and in the latter case, we account for proton deflection in the Galactic magnetic field. } Then, we test the same conjectures with recently published HiRes stereoscopic data (271 cosmic ray with $E>10^{19}$~eV), assuming only neutral primaries, since energies of the HiRes events are unpublished. Finally, we look for correlations of suggested sources with the cosmic rays of the highest ($E>10^{20}$~eV) energies. Sec.~5 contains conclusions and discussion. ",
        "conclusions": "We performed a comparative study of correlations between the arrival directions of the cosmic rays observed by AGASA, Yakutsk and HiRes experiments with 21 catalogs of astronomical objects suggested by different authors as possible UHECR sources. No a posteriori cuts were imposed on the catalogs. The results are described in Sec.~4.2 but should be interpreted with great care because the angular resolution of the cosmic-ray experiments is poor, cosmic magnetic fields are known not quite well and the statistics is not overwhelming.  The analysis of the AGASA and Yakutsk samples demonstrates the presence of correlations with unidentified EGRET sources, BL Lac's and nearby Seyfert galaxies. The corrections for the Galactic magnetic field are very important for these effects.  Only BL Lac type objects exhibit correlations with the HiRes stereo sample as well. With the uncut catalog of 491 confirmed BL Lac's and the set of 271 cosmic rays with energies $E>10^{19}$~eV, we obtained the probability of $0.04$ for these correlations to result from a random coincidence. For the updated catalog of 178 bright ($m<18$) BL Lac's, the probability is as low as $5\\cdot 10^{-4}$. For the AGASA sample (with the GMF correction), ${\\cal P}=4\\cdot 10^{-4}$ for the same catalog; for AGASA and Yakutsk ${\\cal P}=7\\cdot 10^{-4}$. The probability that one of 21 catalogs exhibits, by chance, ${\\cal P}\\le 7\\cdot 10^{-4}$ both for the HiRes sample and one of four other samples is about $3\\cdot 10^{-5}$. The correlations are present at the angular distances of order HiRes angular resolution and much less than the expected deflections of charged particles in cosmic magnetic fields. This suggests that a fraction of the cosmic-ray flux at energies $E>10^{19}$~eV is carried by neutral primaries~\\cite{BL-HiRes}.  Clearly, more detailed information about the cosmic rays observed by HiRes would be very important for understanding the effects discussed here.  Knowledge of the energies of particles would allow to reconstruct (to some accuracy) their deflections in the Galactic magnetic field. A study of shower profiles might help to understand what kind of neutral particles may be present in the cosmic rays and contribute to correlations. Finally, if, for example, a neutral primary is a photon emitted by a blazar, the arrival time of the corresponding air shower may coincide with a temporary change in the state of the source.  This would give much stronger support to conjectures about particular cosmic-ray sources than the statistical studies of the arrival directions.  Further detailed study will specify which particular physically distinguished subclasses of BL Lac's and other objects correlate strongly with the cosmic rays and will give quantitative predictions for the future experiments, in particular for the Pierre Auger Observatory."
    },
    "0410/astro-ph0410094_arXiv.txt": {
        "abstract": "Galaxies undergoing formation and evolution can now be directly observed over a time baseline of some 12 Gyr.  An inherent difficulty with high--redshift observations is that the objects are very faint and the best resolution (HST) is only $\\sim0.5$ kpc.  Such studies thereby combine in a highly synergistic way with the great detail that can be obtained for nearby galaxies through \"archaeological\" studies. Remarkable advances are being made in many areas, due to the power of our observatories on the ground and in space, particularly the unique capabilities of the HST ACS.  Three new developments are highlighted. First is the derivation of stellar masses for galaxies from spectral energy distributions (SEDs) using HST and now Spitzer data, and dynamical masses from both sub--mm observations of CO lines and near--IR observations of optical nebular lines like H$\\alpha$.  A major step has been taken with evidence that points to the $z\\sim 2-3$ LBGs having masses that are a few $\\times 10^{10}$ $M_{\\odot}$. Second is the discovery of a new population of red, evolved galaxies, again at redshifts $z\\sim 2-3$ which appear to be the progenitors of the more massive early--type galaxies of today, with dynamical masses around a few $\\times 10^{11}$ $M_{\\odot}$.  Third are the remarkable advances that have occurred in characterizing drop--out galaxies (LBGs) to $z\\sim6$ and beyond, less than 1 Gyr from recombination. The HST ACS has played a key role here, with the dropout technique being applied to $i$ and $z$ images in several deep ACS fields, yielding large samples of these objects. This has allowed a detailed determination of their properties (e.g., size, color), and meaningful comparisons against lower--redshift dropout samples.  The use of cloning techniques has overcome many of the strong selection biases that affect the study of high redshift populations.  A clear trend of size with redshift has been identified, and its impact on the luminosity density and star formation rate estimated.  There is a significant, though modest, decrease in the star formation rate from redshifts $z\\sim 2.5$ out through $z\\sim 6$.  The latest data also allow for the first robust determination of the luminosity function at $z\\sim 6$.  Last, but not least, the latest UDF ACS (optical) and NICMOS (near--IR) data has resulted in the detection of some galaxies at $z\\sim7-8$. ",
        "introduction": "Direct observation of galaxies in their formative stages is proving to be an extremely powerful approach for understanding how galaxies form and grow.  The key physical processes are, in principle, directly observable.  However, the primary challenge for direct observations is that high redshift galaxies are faint and small, and the rest frame observations are often in the UV, a spectral region that has been poorly characterized locally -- though current observations by GALEX are beginning to provide the needed $z\\sim0$ benchmark.  A key consideration, of particular relevance for this workshop, is that studies of high redshift galaxies are complementary to what can be learned from nearby galaxies -- although strongly synergistic may be a more appropriate characterization.  The ability to do both extraordinarily detailed ``archaeological'' studies of nearby galaxies at one epoch (now) contrasts with the much more superficial characterization that can be carried out for distant galaxies.  The resolution at high redshift, $6-8$ kpc arcsec$^{-1}$ from $z\\sim 0.5$ to $z\\sim 6+$, is dramatically worse than can be achieved at $z\\sim 0$, and presents a serious challenge, even with HST.  The discussion that has taken place at this workshop about bars in $z\\sim1+$ galaxies is an example of how small changes in resolution (from WFPC2 to ACS) can lead to significant differences in conclusions about the properties of high redshift galaxies.  Nonetheless the opportunity to measure the properties of galaxies directly, even at the global level, to within $<1$ Gyr from the Big Bang, is exciting and valuable, even if the details are mostly lacking.  In the talk by Ken Freeman (this volume), a further example was given of how observations  at high and low redshift combine to challenge our view of galaxy evolution.  The thin disk of our galaxy appears to have been undisturbed for some 10 Gyr, corresponding to a relatively quiescent life since $z\\sim 2$.  Yet large disks appear to become rare around $z>1.2$. Are we missing many such disks in our observations?  Or were we just ``lucky'' in the Milky Way....?  A very interesting constraint on the nature of high redshift galaxies comes from the cosmic baryon budget as discussed by Fukugita, Hogan and Peebles (1998).  They compiled a census of where the baryons are at $z\\sim0$.  The vast majority are in gas or plasma (83\\%). Of the 17\\% locally that are in stars, 73\\% are in spheroids (bulges/ellipticals), 25\\% are in disks and 2\\% are in late-type galaxies.  Thus a key issue at high-redshift (at $z>1-2$) is identifying and characterizing spheroid (bulge and elliptical) buildup. Most of the star formation at high redshift must be part of the buildup of spheroids, with different classes of objects (LBGs, evolved $J-K_s$ objects, SCUBA sources etc) exemplifying different phases of this process.  This is particularly true if the 25\\% in disks was assembled relatively late --  maybe after $z\\sim 1.5$?. ",
        "conclusions": "\\begin{figure} \\begin{center} \\includegraphics*[width=.95\\textwidth]{gdi_fig08.ps} \\end{center} \\caption[]{Images ($3'' \\times 3''$) in the $V_{606}, i_{775}, z_{850}, J_{110}, H_{160}$ bands (AB mags) of a $z\\sim 7-8$ object (a $z_{850}$--dropout), plus a very red $(z_{850}-J_{110})_{AB}=1.1$ object (bottom) which nearly met our selection criteria and could be a reddened starburst at $z\\sim6.5$ (or, improbably, a reddened early--type at $z\\sim1.6$) (see Bouwens et al.\\ 2004e). The $J_{110}+H_{160}$ image for each object is included, along with its position in color--color space, in the two rightmost panels.  The sources are in the UDF--IR field (Thompson et al.\\ 2004).  } \\end{figure}   There has been remarkable progress on characterizing galaxies at high redshift, particularly since the ACS came on line after the HST servicing mission SM3B in 2002.  Significant numbers of evolved $z\\sim2-3$ galaxies have been detected. These $J-K_s$ galaxies, or DRGs, appear to be the progenitors of ellipticals and massive bulges (at one stage of their development).  Substantial progress has been made on determining mass scales from kinematics and from stellar populations (SEDs).  This is challenging but the future holds great promise with observations from HST, Spitzer, ALMA, and ground-based 8--10 m IR spectrographs.  Large samples of LBG (strongly star-forming) galaxies at $z>2$ to $z\\sim6$ have been detected. The ACS GOODS + HDF $U,B,V,i$ dropout samples extend to $\\sim$27 AB mag ($\\sim0.2-0.5 L_{*,z=3}$) and total some 1900 galaxies, of which $\\sim$200 are $z\\sim2.5$ $U$--dropouts, $\\sim$1230 are $z\\sim 4$ $B$--dropouts, $\\sim$410 are $z\\sim 5$ $V$--dropouts, and $\\sim$60 are $z\\sim 6$ $i$--dropouts.  The UDF, UDF--Ps and other fields increase the $i$--dropout sample to over 200.  Our cloning analysis of several datasets including the GOODS CDF--S \\& HDF--N, and the UDF plus the UDF--Ps fields has shown that there is a systematic decrease in the size of dropout galaxies as a function of redshift ($2<z<6$).  The best fit is $(1+z)^{-1}$, but the precise scaling may depend on luminosity, ranging from  $(1+z)^{-1}$ to $(1+z)^{-1.5}$.  This implies a $\\sim2-3\\times$ decrease in galaxy size from $z\\sim2.5$ to $z\\sim6$.  Our recent studies on the UDF--Ps + GOODS fields indicated that there is a $\\sim2.5\\times$ increase in the rest--frame continuum UV luminosity density from $z\\sim6$ to $z\\sim3.8$ (within a period of $\\sim$1 Gyr).  The uncorrected star formation rate density at $z\\sim6$ was just $0.38\\pm 0.08\\times$ the star formation rate density at $z\\sim3.8$.  New data from the GOODS, the UDF--Ps, and UDF fields are permitting an $i$--dropout luminosity function to be constructed down to $<0.1 L_{*,z=3}$ ($\\sim29$ AB mag).  Studies indicate that these galaxies could provide the needed reionizing flux at $z\\sim6$ and earlier.  Remarkably, HST ACS and NICMOS data in the UDF have led to the likely detection of (a few) $z\\sim7-8$ $z$--dropout galaxies."
    },
    "0410/astro-ph0410577_arXiv.txt": {
        "abstract": "Full-polarization MERLIN and VLBA observations of the 1720-MHz maser emission from the OH molecule in the supernova remnant W28 are presented. Zeeman splitting ($|\\mathbf{B}| \\approx 0.75$~mG) has been directly resolved between right- and left-circularly polarized spectra for the first time. Linear-polarization position angles and circular-polarization Zeeman splitting observed in the maser emission permit interpretation of a comprehensive picture of the magnetic field at the supernova/molecular cloud interface marked by the masers. We find the post-SNR-shock magnetic field to be well-ordered over the $\\sim 1$~pc covered by the maser region (compared with the approximately 30-pc diameter of the entire SNR) and well-aligned with the shock front that is traced by synchrotron radiation and molecular emission. Based on MERLIN data having a resolution of 200~mas and VLBA data with a resolution of 15~mas, the masers are measured to have deconvolved angular sizes of 90 to 350~mas (225 to 875~AU) with compact cores 20~mas (50~AU) in size, consistent with theoretical expectation and previous observations. ",
        "introduction": "Spectral-line emission from the OH (1720 MHz) satellite transition $({^2}{\\Pi}_{3/2},J={3\\over2},F=2{\\rightarrow}1)$ in supernova remnants (SNR's) was first observed by Goss (1968) toward W28. The maser nature of the emission was noted shortly thereafter along with the recognition that the 1720-MHz line emission, when observed in association with absorption in the other ground-state OH lines, requires a pump mechanism different than that used to explain \\ion{H}{2} region OH masers (Goss \\& Robinson 1968; Ball \\& Staelin 1968; Turner 1969; Robinson, Goss, \\& Manchester 1970). Much later, observations of W28 by Frail, Goss, and Slysh (1994) renewed the interest in the SNR masers while also recognizing (1) the astrophysical diagnostic value of the masers' association with molecular clouds and (2) the viability of the collisional pump suggested by Elitzur (1976) to explain the level population inversion. Subsequent studies and surveys (Yusef-Zadeh et al.\\ 2000; Yusef-Zadeh et al.\\ 1999; Koralesky et al.\\ 1998; Green et al.\\ 1997; Frail et al.\\ 1996; Yusef-Zadeh et al.\\ 1996) have since observed approximately 200 Galactic supernova remnants and found 22 remnants with associated OH (1720~MHz) maser emission.  Observations support the premise that OH (1720~MHz) SNR masers occur where the shock front from a supernova explosion encounters a molecular cloud ({\\it e.g}.\\ W28:  Frail, Goss, and Slysh 1994; Kes~78:  Koralesky et al.\\ 1998; 3C391:  Frail et al.\\ 1996; W44: Wootten 1977). In such interactions, a $C$-type (non-dissociative) shock can produce the relatively rare conditions ($n \\approx 10^5\\,{\\rm cm}^{-3}$, $T \\approx 90$~K, $N_{\\rm OH} \\sim 10^{16-17}\\,{\\rm cm}^{-2}$) needed for the collisionally-pumped OH (1720 MHz) masers (Elitzur 1976; Wardle 1999; Lockett, Gauthier, \\& Elitzur 1999; see also Draine \\& McKee 1993). The suggested temperature, density, and OH column of the OH (1720~MHz) collisional pump are consistent with those determined from observation ({\\it e.g}.\\ Claussen et al.\\ 1997, hereafter C97; Hoffman et al.\\ 2003a, hereafter H03). Furthermore, OH (1720~MHz) SNR masers are observed at the SNR/molecular cloud interface where the shock front is edge-on, moving transversely across the plane of the sky ({\\it e.g}.\\ Reynoso \\& Mangum 2000). Frail \\& Mitchell (1998) and Arikawa et al.\\ (1999) have examined the post-shock gas in W28 using observations of shocked CO ($J = 3 \\rightarrow 2$) using the James Clerk Maxwell Telescope. Figure \\ref{c97+CO} shows the good positional agreement between the OH masers, the shocked gas, and the limb of the synchrotron emission from the SNR (this agreement is discussed in \\S\\ref{lin_disc}).  The intrinsic, transverse sizes of the masers have been the subject of several recent studies. Indeed, the narrow-band maser emission does not permit simple multi-wavelength fitting of the scattering properties to the $\\lambda^2$ wavelength dependence expected of scattering due to the ionized interstellar medium at radio frequencies. Claussen et al.\\ (2002) suggest that the angular scatter-broadening of the maser images in W28 is not significant based on studies of the angular scattering of an adjacent extragalactic source and the temporal broadening of a background pulsar. Hoffman et al.\\ (H03) observed the OH (1720~MHz) SNR masers in IC~443 near the Galactic anticenter where interstellar scatter-broadening is expected to be negligible. Conversely, Yusef-Zadeh et al.\\ (1999) find the OH (1720~MHz) SNR masers in the direction of the Galactic center to be appreciably scatter-broadened. Both Claussen et al.\\ (2002) and H03 find that the observed sizes of the masers are consistent with those predicted from collisional pump theory (${\\sim}10^{15}$~cm, Lockett et al.\\ 1999). However, OH (1720~MHz) SNR maser observations using the VLBA and MERLIN (H03; Brogan et al.\\ 2003; Claussen et al.\\ 1999, hereafter C99) are not sensitive to the very extended (60\\arcsec), ``face-on'' maser emission observed in W28 using the VLA (Yusef-Zadeh, Wardle, \\& Roberts 2003, $60\\arcsec \\approx 10^{18}$~cm projected linear size for a distance $d=2.5$~kpc to the SNR, Vel\\'{a}zquez et al.\\ 2002). The current study will examine only the approximately 40 individual maser ``spots'' observed by C97 (using the VLA) at the location where the shock front of W28 is viewed edge-on (moving transversely across the sky). The maser region observed for this paper is about 1~pc in diameter compared with the approximately 30-pc diameter of the W28 SNR (Fig.~\\ref{c97+CO}).  The intrinsic angular size of OH (1720~MHz) SNR maser spots (on the order of 100~mas) is a relatively problematic angular scale for observation with most existing instruments. Radio interferometers with baselines on the order of 10~km in length do not resolve the maser emission and, furthermore, typically convolve several maser spots together within a large beam (on the order of 1\\arcsec). This ``spatial blending'' confuses both the angular and spectral character of the individual masers, preventing a determination of intrinsic emission properties. Past high-angular-resolution studies have employed VLBI-baseline lengths on the order of 1000~km (fringe spacings of approximately 30~mas). Although OH (1720~MHz) SNR masers are known to have a compact core component (C99; H03), VLBI experiments typically ``resolve out'' most maser spots ({\\it i.e}.\\ the observations are not sensitive to emission on 100-mas angular scales due to low correlated fringe visibility). However, the angular sensitivity of the MERLIN array (baselines on the order of 100~km in length) is well-matched to the angular scale of the masers. The minimum angular resolution of observations using MERLIN (approximately 100~mas) is sufficient to alleviate spatial blending. The largest angular scale sampled in observations using MERLIN (approximately 3\\arcsec) is sufficient to recover all of the flux density of the masers.  Since the OH radical is paramagnetic, the polarization state of OH emission is strongly affected by the magnetic field in the maser gas (compared with non-paramagnetic molecules such as H$_2$CO). Indeed, circular-polarization studies of the Zeeman splitting of OH spectral lines is a significant source of information about the interstellar magnetic field ({\\it e.g}.\\ Caswell 2004). Zeeman observations of the OH (1720~MHz) maser emission in supernova remnants have provided estimates of the magnitude of the magnetic fields in supernova/molecular cloud interactions ($|\\mathbf{B}| \\approx 0.5$~mG: C99; Brogan et al.\\ 2000, hereafter BFGT). Also, observations of the linear polarization of the maser radiation can yield the orientation of the post-SNR-shock magnetic field. However, the suggested theoretical conversion from observational parameters to the three-dimensional magnetic field geometry responsible for the observed maser polarization is not straightforward ({\\it e.g}.\\ Elitzur 1998; Watson \\& Wyld 2001, hereafter WW01; Gray 2003). Nevertheless, measurement of the following three observational polarization parameters is a necessary step toward understanding the physics of the maser emission: (1) the magnitude of the level splitting due to the Zeeman effect, (2) the position angle $\\chi$ of the linear polarization, and (3) the fractional linear polarization $q \\equiv \\sqrt{Q^2+U^2}/I$ of the maser radiation (where $Q$, $U$, and $I$ are Stokes parameters).  To date, all of the full-polarization studies of OH (1720~MHz) SNR masers were completed by C97 and BFGT using the VLA. Since much of the prevailing theory pertinent to these masers was published after the C97 paper (Elitzur 1998; Lockett, Gauthier, \\& Elitzur 1999; Wardle 1999), BFGT contains a thorough discussion of the theories and their implications for both the C97 and BFGT data. The interpretation of many of the results of the current paper also rely on BFGT. However, due to the insufficient spectral and angular resolution of the VLA observations presented in C97 and BFGT, the theoretical suggestions presented by BFGT could not be confidently constrained using the data available at the time. Indeed, the observational hurdle currently confronting OH (1720~MHz) SNR maser studies is the acquisition of full linear- and circular-polarimetry of resolved maser sources. The purpose of the current study is to observe a large sample of OH (1720~MHz) SNR masers with spectral and angular resolution sufficient to apply more definitively the wealth of maser polarization theory discussed by BFGT. The W28 SNR is chosen for this purpose because of the large number of masers (approximately 40).  This paper presents full polarization observations of the OH (1720~MHz) masers in W28 made using the Multi-element Radio Linked Interferometry Network\\footnote{MERLIN is operated as a National Facility by the University of Manchester, Jodrell Bank Observatory, on behalf of Particle Physics and Astronomy Research Council (PPARC).} (MERLIN) and the Very Long Baseline Array (VLBA) of the NRAO\\footnote{The National Radio Astronomy Observatory (NRAO) is a facility of the National Science Foundation operated under a cooperative agreement by Associated Universities, Inc.}. This paper will use the source numbering convention of Claussen et al.\\ (C97 Table~1). ",
        "conclusions": "Using MERLIN, we have imaged all but two of the OH (1720~MHz) SNR masers in W28 with intensity greater than 1~\\jb\\ as observed using the VLA by Claussen et al.\\ (1997). Of the 28 MERLIN detections, we have detected 4 masers (approximately 15\\%) using the VLBA+Y1, consistent with other VLBI detection rates of OH (1720~MHz) SNR masers. Based on MERLIN data with a resolution of 200~mas and VLBA data with a resolution of 15~mas, the masers have deconvolved sizes of 90 to 350~mas (225 to 875~AU) with compact cores 20~mas (50~AU), consistent with theoretical expectation and previous observations.  Based on a number of constraints, the masers appear to be partially saturated, allowing quantitative theoretical examination of the polarization state of the emission. These data contain the first direct detection of resolved Zeeman splitting in OH (1720~MHz) SNR masers. Analysis of the circular-polarization profiles indicates magnetic fields in the maser gas $|\\mathbf{B}| \\approx 0.75$~mG. The linear polarization state of the masers is in good agreement with (1) theoretical expectations of the shock geometry, (2) theoretical suggestions concerning the radiative transfer in the masers, and (3) the magnetic field determined using observations of the synchrotron emission from W28, indicating a projected position angle of $-10$\\arcdeg\\ for the magnetic field in the limb of the SNR near the masers. Also, the position angle of the magnetic field measured using the linear polarization of the masers is in good agreement with both the orientation of the shock observed using thermal molecular emission ({\\it e.g}.\\ Frail \\& Mitchell 1998) and with the magnetic field orientation expected from the SNR expansion ({\\it e.g}.\\ Balsara et al.\\ 2001).  The polarization properties of the masers are observed to change as a function of velocity within the profile, indicating small (approximately 10\\arcdeg) changes in the magnetic field orientation within the maser gas. Using theoretical models, the angle between the line-of-sight and the magnetic field vector is inferred to be $\\theta \\approx 60\\arcdeg$."
    },
    "0410/hep-th0410230_arXiv.txt": {
        "abstract": "In this work we show that the gravitational Chern-Simons term, aside from being a key ingredient in inflationary baryogenesis, modifies super-horizon gravitational waves produced during inflation. We compute the super-Hubble gravitational power spectrum in the slow-roll approximation and show that its overall amplitude is modified while its spectral index remains unchanged (at leading order in the slow-roll parameters). Then, we calculate the correction to the tensor to scalar ratio, $T/S$.  We find a correction of $T/S$ which is dependent on $\\cal{N}$ (more precisely quadratic in ${\\cal N}$), the parameter characterizing the amplitude of the Chern-Simons terms. In a stringy embedding of the leptogenesis mechanism, $\\cal{N}$ is the ratio between the Planck scale and the fundamental string scale. Thus, in principle, we provide a direct probe of leptogenesis due to stringy dynamics in the Cosmic Microwave Background (CMB). However, we demonstrate that the corresponding correction of $T/S$ is in fact very small and not observable in the regime where our calculations are valid. To obtain a sizable effect, we argue that a non-linear calculation is necessary. ",
        "introduction": "Cosmic baryogenesis stands as one of the unresolved problems of particle cosmology. Most models address baryogenesis after the inflationary epoch. Recently the authors of Ref.~\\cite{APS} demonstrated that the baryon asymmetry can be generated during inflation from gravity waves. In this model the lepton number was generated by a quantum expectation value of the Chern-Simons density from Ultra-Violet (UV), birefringent gravitational waves during the inflationary epoch. In a subsequent paper the authors showed that this model can be embedded in string theory, in a model independent manner, through the Green-Schwarz mechanism~\\cite{Gates}. In the stringy embedding there was a huge enhancement of the lepton asymmetry due to a hierarchy in the fundamental string scale and the four dimensional Planck scale.  \\par  Other investigators have searched for parity violation in the Cosmic Microwave Background (CMB) which ultimately leads to a birefringence in gravitational waves ~\\cite{Lue,Pogosian,Balaji}. In this note, we shall study the super-horizon power spectrum and tensor to scalar ratio of scalar and tensor birefringent perturbations produced during inflation. Specifically, we study the spectrum of super-horizon gravity waves whose UV counterparts were responsible for leptogenesis. Is it possible to see a signature of the leptogenesis mechanism in the super-horizon power spectrum?  To address this question we shall derive the tensor to scalar ratio and show that it contains a direct signature of the leptogenesis mechanism which occurred in the UV. Furthermore we show that the scalar to tensor ratio contains the string scale in a model independent way and is in an observable window to this physics.  \\par  The paper is organized as follows. In section~II we derive the equations for the gravitational waves in the presence of the Chern-Simons term. In section~III we provide the exact solutions at various scales and derive the power spectrum as well as the corrected tensor to scalar ratio. In section~IV we relate this modified tensor to scalar ratio to the stringy embedding of gravitational leptogenesis and we conclude with some open issues concerning a consistent quantization and further directions. ",
        "conclusions": "We have evaluated the super-Hubble power spectrum and the tensor to scalar ratio for birefringent gravitational waves produced during inflation. The power spectrum exhibits two interesting regimes, linear and non-linear.  The non-linear regime occurs when $k\\eta \\sim \\Theta ^{-1}$ because the effective potential controlling the evolution of the linear perturbations blows up. At this point, the linear theory of cosmological perturbations that we used is no longer valid. This divergence occurs for all modes (\\ie for all comoving wavenumber $k$) but at different times.  \\par  In this present investigation we only considered the linear regime since at the present moment we were not able to perform a rigorous analysis of the non-linear phenonema. We found corrections which survive to second order in $\\Theta$.  Therefore, in this regime the tensor to scalar ratio gets corrected by $\\Theta$ but this effect is very small.  \\par  If $\\Theta \\lsim 10^{-5}$, one can push the non-linear regime (\\ie the divergence in the effective potential) into the trans-Planckian region where, anyway, other effects (for instance, non perturbative stringy effects) are likely to become important. Somehow, this corresponds to the standard situation where the evolution from the Planck scale to the super-horizon scales is under control and where the perturbations are assumed to emerge from the trans-Planckian regime in the vacuum state, thus ignoring the modifications of the initial conditions that the trans-Planckian physics could cause (in the very same way that we have ignored the effect of the divergence in the potential, provided it is in the trans-Planckian region). However, it is important to keep in mind that this is mostly a technical trick which allows us to work with the linear theory. At a deeper level, the trans-Planckian effects are not expected to play a more important role than in the standard situation. In particular, if the divergence is not in the trans-Planckian regime, only the non-linear theory of cosmological perturbations is necessary in order to calculate the modified $T/S$ irrespectively of any trans-Planckian effects.  \\par  It is interesting to note that the linear regime (where $\\Theta \\sim 10^{-5}$) is compatible with the Stringy Embedding of inflationary Baryogenesis (SEB)~\\cite{Gates}. In this context, the value of $\\Theta$ enhances and gives the resonant frequency associated with the observed baryon asymmetry. As already mentioned before, this value is completely fixed by the string scale and coupling in a model independent fashion. Explicitly, the value of the number ${\\cal N}$ which appears in the definition of $\\Theta $, see Eq.~(\\ref{deftheta}), is given by \\begin{equation} {\\cal N}= \\pi ^2\\sqrt{\\frac{g_{\\rm s}}{2}} \\left(\\frac{M_{_{\\rm Pl}}}{M_{_{\\rm 10}}}\\right)^2\\, , \\end{equation} where $M_{_{\\rm 10}}$ is the ten-dimensional fundamental scale and $g_{\\rm s}$ is the string coupling. Therefore, we established a direct link between stringy quantities and CMB anisotropies. Explicitly, Eq.~(\\ref{cccheck}) can be re-written as \\begin{equation} \\frac{\\left(T/S\\right)_{\\Theta \\neq 0}}{\\left(T/S\\right)_{\\Theta =0}} \\simeq 1+\\frac{0.022}{4}\\left(\\frac{H_{_{\\rm inf}}}{M_{\\rm 10}}\\right)^4 g_{\\rm s}\\epsilon \\, . \\end{equation} In a recent paper the authors of Ref.~\\cite{Gates} found that for reasonable values of string coupling (\\ie weak) and the string scale set to $10^{16}{\\rm GeV}$, both $\\Theta $ can be as small as $10^{-2}$ and the observed baryon asymmetry can be generated. Of course the stringy embedding admits much larger values of $\\Theta$ putting our analysis into the non-linear regime.  \\par  If $\\Theta \\gsim 10^{-5}$, a non-linear calculation is mandatory and one can hope to obtain a significative modification of the ratio $T/S$, maybe observable by future high accuracy CMB experiments. Clearly, the present situation is not very satisfactory since the regime for which sizable effects are expected turns out to be very complicated from the technical point of view. Furthermore, we suspect that the scale associated to the divergence of the effective potential corresponds to a resonant production of lepton number. We wish to report on this issue in a future paper.  \\vspace{0.5cm} \\centerline{\\bf Acknowledgments} \\vspace{0.2cm}  We wish to thank R.~Brandenberger, M.~Peskin and P.~Peter for many enlightening comments and discussions."
    },
    "0410/nucl-th0410043_arXiv.txt": {
        "abstract": "A possibility of the crystalline state of nuclear matter is discussed in a medium--modified Skyrme model.  The interaction energy per nucleon in nuclear matter is evaluated by taking into account the medium influence on single nucleon--skyrmion properties and the tensor part of the nucleon--nucleon potential, and by using a variational method of Hartree--Fock type including zero--point quantum fluctuations.  It is shown that in this approach the ground state of nuclear matter has no crystalline structure due to quantum fluctuations as well as medium modifications of hadron properties. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410211_arXiv.txt": {
        "abstract": "Milliarcsecond resolution Very Long Baseline Array (VLBA) observations of the archetype WR+O star colliding-wind binary (CWB) system WR\\thinspace140 have been obtained at 23 epochs between orbital phases 0.74 to 0.97. The emission in the wind-collision region (WCR) is resolved as a bow-shaped arc which rotates as the orbital phase progresses. This rotation provides for the first time the inclination of the orbit ($122\\deg\\pm5\\deg$), the longitude of the ascending node ($353\\deg\\pm3\\deg$), and the orbit semi-major axis ($9.0\\pm0.5$~mas). The implied distance is $1.85\\pm0.16$~kpc, which requires the O star to be a supergiant, and leads to a wind-momentum ratio of 0.22. Quasi-simultaneous Very Large Array (VLA) observations show the synchrotron spectra evolve dramatically through the orbital phases observed, exhibiting both optically thin and optically thick emission.  The optically thin emission maintains a spectral index of $-0.5$, as expected from diffusive shock acceleration. ",
        "introduction": "Observations of WR\\thinspace140 were obtained at 23 epochs using the VLBA. The campaign started on Jan 4, 1999 near the peak of radio emission around orbital phase 0.75 and was completed Nov 18, 2000 when the radio emission had declined to its low level.  The resulting phase-referenced 8.4-GHz images at three of the observed epochs are shown in Fig.~\\ref{fig:xband_vlba}. Additionally, we also obtained closely concurrent observations with the VLA at five frequencies between 1.4 and 22 GHz.  \\begin{figure*} \\vspace{3.8cm} \\special{psfile=doughertyfig1c.eps angle=0 hoffset=225 voffset=-21 hscale=28 vscale=28} \\special{psfile=doughertyfig1b.eps angle=0 hoffset=105 voffset=-21 hscale=28 vscale=28} \\special{psfile=doughertyfig1a.eps angle=0 hoffset=-15 voffset=-21 hscale=28 vscale=28} \\caption{\\footnotesize 8.4-GHz VLBA observations of WR140 at orbital phases 0.74, 0.86 and 0.93. The contour levels are -1, 1, 1.6, 2.6, 4.1, 6.6, 10.5$\\rho$ where $\\rho=220~\\mu$Jy~beam$^{-1}$. The contour levels and greyscale are identical in each image. The beam size is shown in the lower left corner of the images, and is typically $2.0\\times1.5$~mas$^2$, which at a distance of 1.5 kpc gives a linear resolution of 3 AU. Rotation and proper motion of the WCR are clear. \\label{fig:xband_vlba}} \\end{figure*}  The 8.4-GHz emission detected by the VLBA is clearly resolved.  We identify this emission as arising from the WCR in WR\\thinspace140 since this is the only emission in the system with sufficient brightness temperature ($>10^{5}$~K) to be detected by the VLBA.  The stellar winds, with brightness of $\\sim10^4$~K are undetected.  A bow-shaped ridge of emission is observed at most epochs, as anticipated for the WCR from model calculations \\citep[e.g. see][]{Eichler:1993, Canto:1996, Dougherty:2003}, with the bow shock wrapping around the star with the lower wind momentum - typically the O star. Between orbit phase 0.74 and 0.95, the WCR exhibits rotation from ``pointing'' NW to W, in addition to an east to west proper motion of $\\sim10$~mas. ",
        "conclusions": ""
    },
    "0410/astro-ph0410027_arXiv.txt": {
        "abstract": "Arcsecond-resolution spectral observations toward the protobinary system IRAS 16293-2422 at 344 and 354 GHz were conducted using the Submillimeter Array. Complex organic molecules such as CH$_3$OH and HCOOCH$_3$ were detected. Together with the rich organic inventory revealed, it clearly indicates the existence of two, rather than one, compact {\\it hot molecular cores} ($\\lesssim$ 400 AU in radius) associated with each of the protobinary components identified by their dust continuum emission in the inner star-forming core. ",
        "introduction": "The low-mass protostellar source IRAS 16293-2422 (hereafter I16293) is located in the $\\rho$~Ophiuchus cloud complex at a distance 160 pc from the Sun.  High angular-resolution observations revealed that I16293 is also a protobinary system of two components with a projected separation of $\\sim$5.2\\arcsec~($\\sim$840 AU), and 3-mm continuum observations \\citep{loo00} suggest that the northwest component (0.61 M$_{\\sun}$, hereafter I16293B) is slightly more massive than the southeast component (0.49 M$_{\\sun}$, hereafter I16293A). Single-dish submillimeter line-surveys led \\citet{van95} to conclude that, within a 20\\arcsec~region, I16293 consists of a cold (T$_{\\rm kin} \\simeq$ 10$-$20 K) outer molecular envelope, a warmer (T$_{\\rm kin} \\simeq$ 40 K) circumbinary envelope of 10\\arcsec$-$15\\arcsec~in size, and a hot (T$_{\\rm kin} \\gtrsim$ 80 K) region of dense gas only 3\\arcsec$-$10\\arcsec~in size. This innermost gas was found to be rich in organic molecules and it was suggested that their presence was due to outflows \\citep{hir01} interacting with the circumbinary envelope. However, \\citet{cec99,cea00} showed that the emitting region is warm (100 K), dense $n_{\\rm H_2} \\gtrsim 10^7~{\\rm cm}^{-3}$, and very compact ($\\sim$2\\arcsec$-$3\\arcsec). This led to the suggestion that I16293 contained a {\\it hot molecular core} about 150 AU in size where high dust temperatures (T$_{\\rm dust} \\gtrsim$ 100 K) can facilitate the evaporation of icy grain mantles \\citep{cea00}.  This conclusion is supported by recent observations which show that I16293 {\\it inner} hot core has a molecular inventory rich in complex molecules, similar to the hot molecular cores (HMCs) associated with massive protostars \\citep{caz03}.   Nevertheless, several questions remain that cannot easily be answered by single-dish observations. One major concern is that its inferred size of $\\sim$150 AU ($\\lesssim$~1\\arcsec) \\citep{cea00,sch02} is much smaller than the binary separation in I16293, and comparable to the dimensions of protostellar disks.  Furthermore, it is unknown whether high resolution observations of low-mass hot cores will exhibit clear chemical differentiation on small spatial scales, similar to that seen in massive HMCs \\citep{wyr97}.     Submillimeter Array\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica} observations with arcsecond resolution (1\\arcsec=160 AU) were carried out on 2003 March 14 (compact configuration) and July 12 (extended configuration) with 5 antennae. The phase tracking center of the observations was $\\alpha$(J2000) = 16$^h32^m22^s$.91, $\\delta$(J2000) = -24\\arcdeg28\\arcmin35\\arcsec.52. The digital correlator at the time was configured with eight overlapping segments (\"chunks\") each of 104 MHz bandwidth and 128 channels, except one chunk which had 512 channels, giving a frequency resolution of 0.812 and 0.203 MHz, respectively. The total bandwidth of each sideband was $\\sim$670 MHz covering two frequency ranges: 343.555$-$344.225 GHz (LSB) and 354.211$-$354.881 GHz (USB). Quasars 1743-038 and NRAO530 were observed for phase and amplitude calibration. The flux scale is estimated to be accurate to 25\\%. The synthesized beam sizes are $\\sim$1\\arcsec.3$\\times$2\\arcsec.7 at 344 GHz and $\\sim$1\\arcsec.1$\\times$2\\arcsec.5 at 354 GHz with natural weighting; these beams correspond to a linear scale of $\\sim$200$\\times$400 AU at a distance of 160 pc. The data were calibrated using the MIR software, and imaging was done with MIRIAD software. ",
        "conclusions": "To make positive identifications of the spectral features detected, in addition to the basic requirement of frequency coincidence, an iterative process with the following stringent selection criteria was applied to each candidate species: 1). Expected line intensity - unfavorably weak transitions were rejected. 2). Looking for transitions with similar or stronger line intensities in the bands - if missing, the particular candidate line was excluded. 3). Looking for slightly less favored transitions of the same molecule for a possible sign of low-level ($<$ 3-$\\sigma$) \"detections\" if no other stronger transitions predicted were available in the bands. This would be an indirect verification of the particular line identified. 4). Literature check against lists of known molecules in I16293 in the literature - not mandatory for a detection as SMA is sensitive to dense hot cores. 5). Making the requirement that the molecular candidate has to be a known hot core molecule.   Table~\\ref{tbl-1} lists some of the representative molecular transitions detected. The LSR velocity of each molecular line measured is also listed in Table~\\ref{tbl-1}. Single-dish observations reveal that the LSR velocity of I16293 is centered at 3.9 to 4.2 km s$^{-1}$ and is often found within the range of 3.3$-$4.7 km s$^{-1}$. Asymmetric double-peak line profiles with the blueshifted peak at 2.5$-$3.5 km s$^{-1}$ and the red one at 5.0$-$6.0 km s$^{-1}$ are readily seen in single-dish observations.  On the other hand, an interferometer like the SMA can easily filter out the diffuse ambient cloud, such as the warm gas envelope of I16293 at V$_{\\rm LSR}$ $\\sim$3.9$-$4.2 km s$^{-1}$, and its small beam will tend to pick up the fine kinematic details of dense hot regions. The SMA beam is small enough to resolve the two protobinary components in I16293, thus SMA is able to trace individual velocity component in the rotating disks and/or hot cores. As shown in Fig.~\\ref{fig1}$a$, the two protobinary components are clearly resolved in the continuum emission. Sample images of five organic molecules are shown in Figs.~\\ref{fig1}{\\it b-f}. Fig.~\\ref{fig2} shows the sample spectra of four large organic molecules plus HC$^{15}$N in I16293A and I16293B. By examining HC$^{15}$N emission toward I16293A, two major velocity components, one at 1.7 km s$^{-1}$ and the other at 4.7 km s$^{-1}$ were seen (Fig. 2$a$); in I16293B, one component at 2.5 and the other at 6.0 km s$^{-1}$ were also exhibited (Fig. 2$b$). Furthermore, the small beam of an array suffers less beam dilution, thus is sensitive to high-velocity components within compact sources. In the case of HCN 4$-$3 and HC$^{15}$N, for example, linewidths $\\gtrsim$8 km s$^{-1}$ are seen toward I16293A. The HCOOCH$_3$ line in I16293A is also fairly wide, with a linewidth of $\\sim$8 km s$^{-1}$ (Fig.~\\ref{fig2}$e$). The large linewidth (5$-$8 km s$^{-1}$) is common for simple molecules like HCN, or strong molecular emission like HC$_3$N and CH$_3$OH. For weaker emission, typical linewidths are of 2$-$4 km s$^{-1}$ in I16293A; however, toward I16293B, the linewidths of most molecular lines are notably narrower, i.e., between 2 and 4 km s$^{-1}$.   Because of array characteristics and sensitivity, it is not unusual that an array can preferentially pick up only the stronger emission peak of a double-peak profile as seen by a single-dish. The CH$_3$OH line shows V$_{\\rm LSR}$ $\\sim$ 2.5 km s$^{-1}$ toward I16293A, with two velocity components visible near 4.6 and 1.8 km s$^{-1}$ (Fig.~\\ref{fig2}$c$). The kinematic structure traced is thus similar to what seen by HC$^{15}$N; a similar line profile of $^{13}$CH$_3$OH was also observed toward I16293A. Toward I16293B, the CH$_3$OH emission is weaker and is mainly from the blueshifted component near 2.5 km s$^{-1}$ (Fig.~\\ref{fig2}$d$), which again agrees well with what was seen by HC$^{15}$N. Since the optically thin HC$^{15}$N emission basically traced the densest regions such as the disks, it suggests methanol may also reside in the same dense regions. However, it is also plausible that the two velocity components seen in the methanol line profile are due to self absorption at $\\sim$3.9 km s$^{-1}$. Toward I16293B, the LSR velocities of HCOOCH$_3$ lines were found consistent with each other at $\\sim$3.9 km s$^{-1}$, within the velocity uncertainty of $\\sim$0.7 km s$^{-1}$. However, a non-trivial velocity scattering with V$_{\\rm LSR}$ between 3.9 and 6.0 km s$^{-1}$ was observed toward I16293A. A close examination of the integrated intensity maps of HCOOCH$_3$ show their peak emission positions are situated at $\\sim$0\\arcsec.5 east (V$_{\\rm LSR}$ = 3.9 km s$^{-1}$) and south (V$_{\\rm LSR}$ = 5.9 and 6.0 km s$^{-1}$) of the I16293A center position where V$_{\\rm LSR}$ $\\sim$ 4.6 km s$^{-1}$. Again, such a velocity field variation traces the general velocity gradient seen by the HC$^{15}$N emission nicely.     All spectral emission appears to originate from two compact regions of a beam-convolved size of $\\sim$200$\\times$400 AU in radius, concentrated toward the two protobinary components I16293A and I16293B. An exception is the 4$-$3 HCN transition which shows complicated kinematic structure (Takakuwa \\etal~2004, in preparation).   The beam-averaged line-of-sight column densities and fractional abundances of all observed molecules are listed in Table~\\ref{tbl-2}, where a gas column density $N_{\\rm H_2} = 1.6 \\times 10^{24}$ cm$^{-2}$ \\citep{sch02} was adopted for both protostellar cores.  Column densities were measured at the peak emission positions of each integrated spectral line and were derived assuming thermalized level populations and optically thin lines. At these high densities, gas and dust are expected to be thermally well-coupled and so an excitation temperature T$_{\\rm ex}$ $\\simeq$~T$_{\\rm dust}$ = 100 K was adopted, except for molecular transitions with energy levels $>$ 500 cm$^{-1}$ ($\\gtrsim$700 K) where $T_{\\rm ex}$ = 300 K was used \\citep{sch02}. For HCOOCH$_3$, with multi-line detections, the column densities were derived from rotation diagrams (Fig.~\\ref{fig3}); we found T$_{\\rm ex}$ = 105 K for I16293A, and 116 K for I16293B. The assumption of T$_{\\rm ex}$ = 100 K applied for transitions of lower energy levels is hence warranted.   The derived fractional abundances and relative column densities are in good agreement with those derived from submillimeter observations of the Orion KL and Sgr B2 massive HMCs \\citep{sut91,sut95}. When molecular emission is mainly from the compact cores, abundances derived using the smaller SMA beam can be much higher (e.g. $c$-C$_3$H$_2$) than found in previous single-dish observations \\citep{van95,caz03}. This explains why CH$_3$OH abundances derived from modeling higher energy transitions \\citep{sch02} in the {\\it inner} I16293 core show good agreement with our values, and suggests that the lower array abundances for HCN and HCOOCH$_3$ is indicative of less compact emission where the extended component is filtered out by the interferometer.   The high abundances of organic molecules, particularly methanol, indicates that icy mantles have recently been evaporated \\citep{cha92}. For many molecules detected in both sources, the measured abundances only differ by factors of about 2$-$3; as expected if both cores collapsed from the same cold core material.  Nevertheless, the apparent absence of $c$-C$_3$H$_2$ and CH$_2$CO in I16293A, if confirmed, raises the prospect that future observations could discover pronounced chemical differentiation between the two cores.        We have resolved the hot cores surrounding both protostellar sources in I16293 and measured the chemical composition in each.  These preliminary observations demonstrate that the SMA will be an important tool to further explore the connection between the volatile organic chemical composition of the hot cores associated with massive protostars, those of Solar mass, and the composition of comets."
    },
    "0410/astro-ph0410357_arXiv.txt": {
        "abstract": "{ Spectroscopic and imaging data of the peculiar galaxy Mkn 298 are presented in this paper. Narrow-band \\ion{H}{$\\alpha$} and broad-band $R$ images are used to study  the star formation rate in the galaxy and its morphology, which is typical of a merging system. Long-slit and integral field spectra are used to assess the kinematics of gas and stars, and the nature of the ionizing source at different distances from the nucleus. In particular, the nucleus of Mkn 298 is characterized by peculiar line ratios: \\ion{[N}{ii]}$\\lambda6583$/\\ion{H}{$\\alpha$} is typical of \\ion{H}{ii}-like regions, while \\ion{[O}{i]}$\\lambda6300$/\\ion{H}{$\\alpha$} could indicate the presence of an active galactic nucleus. We show that models where a shock component is  added to  photoionization from a starburst allow to reproduce the observed line ratios. Mkn 298 is thus most likely a star forming galaxy, rather than a galaxy hosting an active nucleus. ",
        "introduction": "\\object{Mkn 298} (IC 1182) is known as a morphologically peculiar system (Figure~\\ref{fig:bim}) located in the Hercules supercluster. It shows a chain (c, d, e) of small compact blue emitting regions \\citep{stock} aligned on its eastern side up to $\\sim 80$\\arcsec\\ from the main body of the galaxy, which is classified in NED (NASA/IPAC Extragalactic Database) as a SA0+ galaxy. These regions are aligned with a tail (b) resembling a spiral arm located on the eastern side of the galaxy (a). On the western side a very faint trace of spiral arm is also visible. A faint emission is visible in the c knot, behind a foreground star \\citep{stock,met:pron}. The d and e knots are  characterized by an emission line spectrum. They were identified by \\citet{braine} as a tidal dwarf galaxy (TDG hereafter) forming from material ejected from the disk of a galaxy after a collision.   The classification of Mkn 298 as a galaxy hosting an active nucleus is controversial in literature. \\citet{koski} and \\citet{vo} classified it as a Seyfert 2 galaxy with an extremely low \\ion{[N}{ii]}$\\lambda$6583/\\ion{H}{$\\alpha$} line ratio. Conversely \\citet{vieg:gruen} enclosed Mkn 298 in a sample of LINERs. Variability was claimed by \\citet{met:pron} who observed  changes of the relative intensity of \\ion{[O}{iii]}$\\lambda\\lambda$4959,5007 and \\ion{H}{$\\beta$} in the period 1969--70, 1977. The detection  of variable soft X-ray emission \\citep{rafanelli} with a luminosity typical of Seyfert 2 galaxies,  seems to confirm the presence of an active nucleus. However, as noted by \\citet{moles} it is not possible to exclude that the X-ray emission is produced in star-forming regions \\citep{zezas}. Finally, \\citet{moles} detected the presence of several knots in the nucleus; they interpreted Mkn 298 as a very luminous starburst (SFR $\\sim$ 90 $M_{\\sun}$/yr) with very low metallicity ($Z/Z_{\\sun} \\sim 0.06 - 0.1$).   An extended \\ion{H}{i} distribution peaked on the TDG was detected  using VLA data \\citep{braine,iglesias}. \\citet{braine} report the detection of CO(2-1) and CO(1-0) emission in Mkn 298, however the emission is peaked on the main galaxy and is below detection in the  TDG. They computed a total amount of H$_2$ mass in the system $M_{\\rm mol} > 5 \\times 10^9$ $M_{\\sun}$, and an upper limit $M_{\\rm mol} <6 \\times 10^7$ $M_{\\sun}$ in the TDG. This would imply that  only a small fraction of HI was transformed into H$_2$ in the TDG, while in the nucleus the conversion already took place. This is consistent with Mkn 298 being in the late stage of a merging: as outlined by \\citet{hibbard}, in advanced mergers the atomic gas is mostly relegated to the outer tidal features and is almost absent in the central regions, whereas the molecular gas is concentrated in the remnant body.   The aim of this paper is to study in detail the physical conditions in the different components (galaxy body and tail) of the Mkn 298 system by the analysis of their emission-line spectra. Observations and data reduction steps are described in Sec.~\\ref{sec:observations}. The morphology, star formation rates and kinematics are analyzed and discussed in  Sec.~\\ref{sec:morphology},  Sec.~\\ref{sec:sfr} and Sec.~\\ref{sec:kinematics} respectively. In Section~\\ref{sec:linerat} emission line ratios are used to derive chemical abundances. Observed ratios are then compared with the results from both pure  photoionization models and photoionization+shock models, using the {\\sc Cloudy} and {\\sc Suma} codes respectively. Conclusions are finally drawn in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "}  In this paper we have presented new imaging and integral-field spectroscopic data analysis of Mkn 298, aimed to investigate its controversial nature. By means of ground-based broad-band R and HST band V images we have performed a detailed analysis of the morphology of this galaxy, obtaining that Mkn 298 is an early-type spiral, likely a S0 galaxy (B/T $\\sim 0.78$), with a de Vaucouleurs bulge and without any evident nuclear point-like source. In addition, we have pointed out the presence of weak stellar shells around the galaxy. These shells, together with the bright and extended eastern tail, and the nuclear dust lane, contribute to strongly support the idea that Mkn 298 is the result of a major merger event. Ionized gas in Mkn 298 is distributed in two or maybe three main knots, visible in the H$\\alpha$-continuum subtracted image, and whose spectra show mostly the typical emission lines produced by thermal ionization. This gas is extended and oriented as the dust lane, suggesting that these two features are likely to be connected. The kinematic investigation of H$\\alpha$ and \\ion{[O}{iii]} velocity fields show clearly that the gas is rotating with a maximum deprojected velocity of $+ 250$ \\kms with respect to the systemic velocity of the galaxy. Therefore, the integral-field data confirm the previous hypothesis by \\citet{moles} that a disk dominated galaxy is the secondary component of the merger event.  The analysis of the emission line ratios shows that both AGN photoionization models and starburst+shock models with sub-solar metallicity ($Z/Z_{\\sun}$ $\\le$ 0.6) are able to fit the observed values. In particular the high \\ion{[O}{i]}/H$\\alpha$, which could indicate the presence of an AGN, is also well reproduced by starburst+shock models, where photoionization is given by a young ($t < $ 2.5 Myr) stellar population  and shocks are produced in low velocity$-$density clouds ($V_{\\rm s}$ = 80 km/s, $n_{\\rm 0}$ = 60 cm$^{-3}$). AGN+shock models are ruled by the \\ion{[O}{iii]}$\\lambda$4363/\\ion{H}{$\\beta$} ratio which is too high if compared with observed values and by the flux from the AGN which is unrealistically low.   Considering that emission line ratios are constant over all the regions where they are detected, with no well defined spatial trend, we conclude that the starburst+shock models provide a more likely possibility than an AGN: merging would induce both collisional processes and star formation in the gas, thus producing the observed line ratios. No evidence was found for the very low metallicity  claimed by \\citet{moles}, Z/Z$_{\\sun}$ $\\le$ 0.1. We do not confirm the high SFR ($\\sim$ 90 $M_{\\sun}$/yr) found by \\citet{moles}; we find instead that  SFR/\\ion{H}{$\\alpha$} $\\le$ 10 $M_{\\sun}$/yr. A comparable SFR was found taking the IRAS infrared luminosity at the position of Mkn 298, even if this value is very uncertain.  As it concerns the knots in the tidal tail (d,e), the presence of shocks is not so compelling as both pure photoionization from stars and photoionization+shock models provide a good fit of the observed ratios. A sub-solar metallicity ($Z/Z_{\\sun}$ $\\le$ 0.4) is found in these regions as well.    \\vskip 50pt"
    },
    "0410/astro-ph0410482_arXiv.txt": {
        "abstract": "We present two classes of stars with yet unknown evolutionary phase: the B[e] supergiants and the so-called unclassified B[e] stars. While the B[e] supergiants are luminous post-main sequence stars with high mass progenitors, not much is known of the unclassified stars. We discuss how it might be possible to determine their evolutionary phases in the area of new large telescopes. ",
        "introduction": "Stars showing the B[e] phenomenon have several characteristics of their optical spectrum in common: they are of spectral type B, have strong Balmer emission lines, and many permitted and forbidden emission lines of predominantly low ionization metals like e.g. Fe{\\sc ii} and O{\\sc i}. In addition, these stars are known to possess a strong near- or mid-infrared excess due to hot circumstellar dust. Since most of the characteristics are based purely on the optical spectrum, these stars can be of quite different evolutionary phase, and \\citet{l1998} grouped all stars with the B[e] phenomenon accordingly. The different classes they found are: B[e] supergiants, Herbig Ae/B[e] stars, and compact planetary nebula B[e] stars. The biggest group, however, is formed by stars of unknown evolutionary phase, the so-called unclassified B[e] stars. In this paper we want to concentrate on two classes of B[e] stars: the supergiants and the unclassified stars. Both classes have in common, that their evolutionary phase is still unknown and we want to strengthen the need of large telescopes for the determination of the evolutionary phases of these stars. ",
        "conclusions": ""
    },
    "0410/astro-ph0410161_arXiv.txt": {
        "abstract": "The extra-planar H\\,I in the inner parts of the Milky Way has been discovered to contain numerous cloud-like structures when observed in the 21cm line with the Green Bank Telescope. These halo clouds have motions consistent with Galactic rotation and do not seem to be related to the classic high-velocity clouds. They are found to distances $>$1 kpc from the plane and can contain hundreds of $M_{\\sun}$ of H\\,I. Spectra of many of the halo clouds show evidence of coexisting cool and warm H\\,I phases. A preliminary high-resolution study of one of the clouds suggests that it consists of a diffuse envelope and a few dense cores, with a peak $N_{H\\,I}$ reaching $4 \\times 10^{20} $ cm$^{-2}$. The clouds are often organized into larger structures, one example of which was discovered near $\\ell=35\\deg$ rising higher than 2 kpc above the Galactic plane. New observations should answer some fundamental questions about the nature of these clouds. ",
        "introduction": "The Milky Way provides an opportunity to study extra-planar gas in great detail, yet the view is often confused and uncertain: at high latitudes where spectra can be simple, there is rarely any way to estimate a distance to the gas being measured,  while at low latitudes, where Galactic rotation can  separate gas kinematically, spectra are blended and confused. Neutral gas in significant amounts has been found far from the Galactic plane, and measurements in optical, radio, and UV absorption lines have provided detailed information on its physical properties, but the data do not form a consistent picture, do not reveal the structure of the material, and leave us with questions about how to generalize from one observation to the next \\citep{munchzirin,savagedeboer,l84,lhs,savage90, spitzfitz93,albert,sembsav,spitzfitz97,kalberla98,howk_sf}.  Our understanding of the lower halo changed dramatically upon commissioning of the new Green Bank Telescope (GBT) \\citep{fjlspie98, jewellspie} whose $9\\arcmin$ angular resolution and superb  sensitivity to low surface brightness emission revealed structures in Galactic H\\,I which were only hinted at in earlier 21 cm data. We now know that much of the extra-planar H\\,I in the inner parts of the Milky Way is organized into discrete objects, some of which look like archetypical diffuse interstellar clouds \\citep{fjl2002, fjl2004}.  The distance to many of the clouds can be estimated with some accuracy, allowing us to derive quantitative information on their temperature, density, and mass. We are  now in a new era in the study of the Galactic halo: it may come to be the model for the disk-halo interface in other systems as well. ",
        "conclusions": ""
    },
    "0410/astro-ph0410410_arXiv.txt": {
        "abstract": "Galaxy interactions/mergers, gravitational instabilities and density waves, such as bars, are frequently invoked to trigger starbursts. These mechanisms have been explored through numerical simulations, with the help of various star formation recipes.  Gravitational instabilities are necessary to initiate star formation, but the main trigger might be the gas flows, to provide sufficient fuel in a short time-scale. Gas accretion is also acting on the dynamics, in favoring bars/spirals, which will drive the gas inwards. Large amounts of external gas accretion are required to explain the bar frequency, and this accretion rate can be provided by the cosmic filaments, as supported by cosmological simulations. Subsequent interactions can then trigger starbursts by driving this accreted gas inwards. ",
        "introduction": "It is now widely accepted that strong galaxy interactions can produce violent starbursts, and ultra-luminous infra-red galaxies (ULIRGs) are all mergers (e.g. Sanders \\& Mirabel 1996). But interacting galaxies in general do not show intense starbursts (Bergvall et al 2003), or only in their centers. The tidal parameter, depending on mass and distance of companions, is not correlated to any star formation tracer. Starbursts are rare in the local universe, and much rarer than galaxy interactions. The only perceptible effect is that interacting galaxies have more mass - gas, dust \\& young stars, and are more concentrated to the center (Bergvall et al 2003). One of the main conclusions that can be drawn for giant galaxy starbursts is that galaxy interactions is a necessary condition to trigger them, but not sufficient. Certainly, other parameters have to be taken into account, such as the gas content, the geometry/velocity of the encounter, and also the phases of the interaction, the starburst occuring only during about 10\\% of the interaction time, according to simulations (e.g. Mihos \\& Hernquist 1996).  For dwarf galaxies, interactions are even not a necessary condition for triggering: starbursting dwarfs have no more companions (Brosch et al 2004), although it might be different for Blue Compact Dwarfs (BCD, Hunter \\& Elmegreen 2004). Dwarf galaxies with starbursts present asymmetries, and a possible trigger could just come from gas accretion: the gas infalling in dark-matter haloes may experience sloshing and oscillations favoring compression and instabilities. This mechanism might be relevant for primordial galaxies (Brosch et al 2004).  One of the best evidence of triggering can be obtained by tracing the fossil records of the star formation history in nearby galaxies. In the Small Magellanic Cloud, the number of stars at a given age has been determined by Zaritsky \\& Harris (2004). They noticed that the derived SMC star formation history reveals some bursts corresponding to pericenters of its orbit around the Milky Way. A model taking the tidal triggering into account fits better the observations, although the quantitative influence of the tide is uncertain, between 10 and 70\\% of the star formation could be tidal in origin, cf figure \\ref{zar-f11}. One of the problem in the fit comes from the too large increase in star formation about 3-4 Gyrs ago, which can only be explained by external gas accretion.  A significant gas infall, by about 50\\% of the total gas mass, improves the fit, and accounts also for the observed age-metallicity relation. The metallicity is not monotonously increasing with decreasing age, but remains constant for old stars until 3 Gyrs, indicating that dilution is required from deficient gas infall.   \\begin{figure} \\includegraphics[width=.6\\textwidth]{combes-bh-f1.ps} \\caption{Star formation history in the SMC derived from a tidally triggered model (dash-line), or a tidally triggered model with infall (solid line), compared to the observed number of stars at each age, from Zaritsky \\& Harris (2004). The insert is a zoom of the recent 2.5 Gyr.} \\label{zar-f11} \\end{figure} ",
        "conclusions": "Star formation depends essentially on the gas supply and is only slightly influenced by the large-scale dynamics. Galaxy interactions act to drive the accreted gas radially inwards and trigger central starbursts. The accretion of gas not only regulates the star forming history in galaxies but also their dynamical state  (bars, spirals, warps, m=1...), and both phenomema interact in a feedback loop. For galaxies in the field, accretion is dominant, and explain bars and spirals, and the constant star formation rate for intermediate types. In rich environments, galaxies experience quicker evolution, mergers are much more important, and secular evolution of galaxies through gas accretion is halted at z$\\sim 1$, since galaxies are stripped from their gas reservoirs."
    },
    "0410/astro-ph0410626_arXiv.txt": {
        "abstract": "We have used the Giant Meterwave Radio Telescope (GMRT) to measure the Galactic HI 21-cm line absorption towards 102 extragalactic radio continuum sources, located at high ($|b|~>~15^{\\circ}$) Galactic latitudes. The Declination coverage of the present survey is $\\delta~>~\\sim~-45^{\\circ}$. With a mean rms optical depth of $\\sim0.003$, this is the most sensitive Galactic HI 21-cm line absorption survey to date. To supplement the absorption data, we have extracted the HI 21-cm line emission profiles towards these 102 lines of sight from the Leiden Dwingeloo Survey of Galactic neutral hydrogen. We have carried out a Gaussian fitting analysis to identify the discrete absorption and emission components in these profiles. In this paper, we present the spectra and the components. A subsequent paper will discuss the interpretation of these results. ",
        "introduction": "\\label{sec:intro} The distribution atomic hydrogen in the Galaxy and its physical properties have been extensively studied. Soon after the discovery of the HI 21-cm line, a number of single dish HI surveys were conducted (See Burton, \\nocite{wbb88}1988 for a useful compilation of the early HI surveys). The single dish HI absorption spectra are limited by the errors due to the variation of HI emission intensity over the angular scales smaller than the telescope beam. Interferometric surveys are a better alternative for HI absorption studies. An interferometer rejects the low spatial frequencies where HI emission is dominant, resulting in reliable absorption spectra. There have been a number of interferometric HI absorption studies, often supplemented by single dish observations to measure the HI emission.  The early HI 21-cm line absorption and emission studies led to the emergence of a global picture of the interstellar medium (Clark, Radhakrishnan \\& Wilson, \\nocite{crw62}1962; Clark, \\nocite{barry65}1965). These and later studies paved way for the models of the interstellar medium of the Galaxy. Field, Goldsmith \\& Habbing  \\nocite{fgh69}(1969) modeled the ISM as cool dense concentrations of gas, often referred to as ``interstellar clouds'' (the Cold Neutral Medium or CNM) in pressure equilibrium with a warmer intercloud medium (the Warm Neutral Medium or WNM). While this initial model of the ISM has been refined considerably by later studies (Wolfire et al, \\nocite{wolfire95}1995), the basic picture of the ISM with cold diffuse clouds and a warmer intercloud medium has survived. The spin temperature of the cold clouds which constitute the CNM was estimated to be $\\sim$80 K and that of the WNM to be $\\sim$8000 K. The cold clouds manifest as narrow Gaussian features with typical velocity dispersions of a few km s$^{-1}$ in both HI emission and absorption profiles. The intercloud medium, on the other hand, is identified with broad Gaussian features with typical velocity dispersions $\\gtsim$10 km s$^{-1}$. The intercloud medium is usually detected in emission, since the HI absorption in the warm gas is very weak. We summarize the results from some of the important surveys below.  Radhakrishnan et al (\\nocite{rad72a}1972a, \\nocite{rad72b}1972b) used the Parkes Interferometer to study the HI absorption towards 35 extragalactic radio sources, in the Galactic latitude range 6$^{\\circ}$ $\\le$ $|b|$ $\\le$ 74$^{\\circ}$ and 53 Galactic radio sources in the lower Galactic latitudes, mostly located at $|b|$ $<$ 2$^{\\circ}$. In addition, the Parkes 64 m telescope was used to obtain the HI emission towards all these directions. The velocity resolution of these observations was 2.1 km s$^{-1}$. The rms noise in the HI optical depth profiles varied from $\\sim$0.01 to $>$0.1, depending on the flux density of the background source. They concluded that the HI absorption features are arising in discrete concentrations of gas with a spin temperature in the range 60 -- 80 K. They also derived the number density of these features to be $\\sim$2.5 per Kpc (Radhakrishnan \\& Goss, \\nocite{rg72}1972). The number of such concentrations of gas for a given optical depth $\\tau$ was $\\propto$ $e^{-\\tau}$. Lack of absorption in the intercloud medium, which was identified as the wide ``shoulders'' in HI emission profiles, enabled them to put a limit of $>$750 K for the spin temperature of this gas.  Dickey \\& Benson \\nocite{dickey82}(1982) used the NRAO 300 ft and 140 ft telescopes as an interferometer to detect absorption in the 21-cm line towards 64 radio continuum sources. The HI emission profiles were obtained using the 300 ft telescope alone. These sources were both Galactic and extragalactic and spread over a range of latitudes 0$^{\\circ}$ $\\ltsim|b|$ $\\ltsim70^{\\circ}$. They have produced HI absorption profiles with an rms optical depth in the range $\\sim$0.007 -- $\\sim$0.16, with a velocity resolution ranging from 1.3 $-$ 5.3 km s$^{-1}$. Among the main results of this study was the realization that for lower Galactic latitudes ($|b|$ $\\ltsim$15$^{\\circ}$), HI emission surveys using single dish telescopes would miss a significant amount of gas ($\\sim$40\\%) due to HI self absorption, wherein the cold HI gas in the foreground absorbs the HI line emission from the background gas. They also found more HI gas with lower spin temperature (100 $\\ltsim$T$_{S}$ $\\ltsim$150 K) at lower latitudes ($|b|$ $<$ 2$^{\\circ}$) as compared to the previous HI line surveys. They concluded that such a behaviour is the result of velocity blending.  Mebold et al \\nocite{egs81}(1981) carried out an HI absorption survey towards 69 sources in the range 0$^{\\circ}$ $\\ltsim|b|$ $\\ltsim$80$^{\\circ}$ using the NRAO 3-element interferometer. They obtained the corresponding HI emission using the Effelsberg 100 m telescope, the 91 m Green Bank telescope or the 64 m Parkes telescope. The velocity resolution of these profiles were in the range 0.42 -- 3.3 km s$^{-1}$. These profiles, on an average, had an rms in HI optical depth $\\gtsim$0.05. They found most of the HI absorption features to have a spin temperature in the range 20 to 140 K. For the HI absorption data at $|b|$ $>$ 15$^{\\circ}$, they found indications for a bimodal distribution in the radial velocity distribution of absorbing features. However, since the number of components at higher radial velocities were small, they were unable to study its significance.  Till recently, the Very Large Array (VLA) used to be the only instrument with a collecting area comparable with large single dish telescopes. The survey by Dickey et al. \\nocite{dickey83}(1983) using the VLA is limited to lower Galactic latitudes ($|b|$ $<$ 10$^{\\circ}$). Moreover, the optical depth detection limit for this survey is $\\sim$0.1 (3$\\sigma$). From the various HI absorption surveys carried out so far (Radhakrishnan et al., \\nocite{rad72a}1972a, \\nocite{rad72b}1972b; Mebold et al., \\nocite{egs81}1981; Dickey \\& Benson, 1982 \\nocite{dickey82}; Dickey et al., \\nocite{dickey83}1983), more than 600 absorption spectra are available, but the optical depth detection limits of more than 75\\% of these are above 0.1. From the available results, the cloud population observed in HI absorption seem to have a Gaussian random velocity distribution. The dispersion in the random velocities of HI absorption features is $\\sim$7 km s$^{-1}$ (Dickey \\& Lockman, \\nocite{araa90} 1990, and references therein).  The low optical depth regime of Galactic HI is largely unexplored except for the HI absorption studies using the Arecibo reflector by Dickey et al \\nocite{dst78}(1978) and more recently by Heiles \\& Troland \\nocite{ht03a}(2003a, \\nocite{ht03b}b). Dickey et al \\nocite{dst78} measured HI absorption and emission towards 27 extragalactic radio continuum sources located at high and intermediate Galactic latitudes ($|b|$ $>$ 5$^{\\circ}$). The rms optical depth in their spectra were typically $\\sim$0.005. In many of the profiles the systematics in the band dominate the noise in the spectrum. As we noted earlier, these observations are not impervious to HI  emission fluctuations introducing errors in the absorption profile (Dickey \\& Lockman, \\nocite{araa90}1990). Radhakrishnan \\& Goss \\nocite{rg72}(1972) found the number of HI absorption features for a given optical depth, $\\tau$ to be proportional to $e^{-\\tau}$. However, the data by Dickey et al \\nocite{dst78}(1978) as well as the later survey using the Green Bank Interferometer (Mebold et al, \\nocite{egs82}1982) indicated that for $|b|$ $>$ 15$^{\\circ}$ this dependence is steeper than $e^{-\\tau}$. This trend was explained as due to increase of low optical depth features with increasing angular and velocity resolution (Mebold et al, \\nocite{egs82}1982). Dickey et al \\nocite{dst78}(1978) also noted that the velocity distribution of HI absorption features is dependent on the optical depth. For the optically thin clouds ($\\tau$ $<$ 0.1), the velocity dispersion was $\\sim$11 km s$^{-1}$, whereas for the optically thick clouds ($\\tau$ $>$ 0.1) this value is $\\sim$6 km s$^{-1}$. Heiles \\& Troland (\\nocite{ht03a}2003a, \\nocite{ht03b}b) analysed HI absorption and emission profiles toward 79 lines of sight. A good fraction of these directions (66 out of 79) were at Galactic latitudes $|b|$ $>$ 10$^{\\circ}$. They found evidence for an excess of low column density ($N_{\\scriptscriptstyle{\\rm{HI}}}$ $<$ 5 $\\times$ 10$^{19}$ cm$^{-2}$) CNM components. ",
        "conclusions": ""
    },
    "0410/astro-ph0410556_arXiv.txt": {
        "abstract": "We investigate the dependence of observational properties of black hole X-ray binaries on the inclination angle $i$ of their orbits.  We find the following: (1) Transient black hole binaries show no trend in their quiescent X-ray luminosities as a function of $i$, suggesting that the radiation is not significantly beamed.  This is consistent with emission from an accretion disk.  If the X-rays are from a jet, then the Lorentz factor $\\gamma$ of the jet is $<1.24$ at the 90\\% confidence level.  (2) The X-ray binary 4U1543--47 with $i\\sim21\\degr$ has a surprisingly strong fluorescent iron line in the high soft state.  Quantifying an earlier argument by Park et al. (2004), we conclude that if the continuum X-ray emission in this source is from a jet, then $\\gamma<1.04$.  (3) None of the known binaries has $\\cos i<0.25$ or $i>75\\degr$.  This fact, plus the lack of eclipses among the 20 black hole binaries in our sample, strongly suggests at the 99.5\\% confidence level that systems with large inclination angles are hidden from view.  The obscuration could be the result of disk flaring, as suggested by Milgrom (1978) for neutron star X-ray binaries.  (4) Transient black hole binaries with $i\\sim70-75\\degr$ have significantly more complex X-ray light curves than systems with $i\\lesssim65\\degr$.  This may be the result of variable obscuration and/or variable height above the disk of the radiating gas. ",
        "introduction": "An accretion disk is intrinsically non-spherical since the angular momentum vector of the orbiting gas specifies a unique direction. Depending on the direction from which one views the disk --- face-on or edge-on or something in between --- one expects differences in what one observes.  Apart from the obvious $\\cos i$ factor which determines the projected area of the disk, $i$ being the angle between the disk normal and the direction to the observer, there could be additional beaming effects if the radiating gas has relativistic motion.  The orbital motion of the gas in the disk leads to blue- and redshifts for observers who are close to edge-on ($i\\to 90\\degr$), while the radiation from a jet is relativistically boosted for observers who are more face-on ($i\\to 0\\degr$).  There may also be effects due to obscuration.  For instance, the radiation emitted by one part of the disk may be blocked from the observer's view by other parts of the disk; alternatively, the secondary star may do the obscuring, in which case one expects eclipses in the light curve.  In order to study these inclination effects, we need a sample of disk systems for which we have independent estimates of $i$.  X-ray binaries are particularly useful in this regard.  Because the gas accreting onto the compact star is supplied by the secondary, the angular momentum of the material is aligned with the orbital angular momentum of the binary system.  Thus, the accretion disk is expected to be coplanar with the binary orbit.  The inclination of the binary orbit with respect to the observer's line-of-sight can be estimated from the light curve of the secondary (e.g., Orosz \\& Bailyn 1997). This inclination directly gives the inclination angle $i$ of the disk.  We focus in this paper on black hole X-ray binaries, limiting our study to the 20 black hole systems with dynamical mass information (McClintock \\& Remillard 2004; Orosz et al. 2004; Casares et al. 2004).  Of these 20 sources, 17 are transient black hole X-ray novae (BHXN) and the remaining 3 are persistent sources.  We collect from the literature estimates of the inclination angles of these binaries and investigate their luminosities, spectra and light curves. Our aim is to identify any patterns or systematic variations with respect to inclination in the data.  In \\S 2, we discuss relativistic beaming effects and look for evidence of jets in certain spectral states of black hole binaries.  Then, in \\S 3 and \\S 4 we discuss obscuration effects.  We conclude in \\S 5 with a summary. ",
        "conclusions": "Starting with the available information on the orbital inclinations of black hole X-ray binaries, we have reached several conclusions in this paper.  \\noindent 1. We find that the quiescent X-ray luminosities of BHXN show no trend with inclination $i$ (Table 1, Fig. 1).  Comparing the observations with the expected trend if the X-ray emission is from a jet oriented perpendicular to the binary orbital plane, and considering in particular the low-inclination source 4U1543--47, we conclude that the bulk Lorentz factor $\\gamma$ of the jet is $\\leq1.24$ at the 90\\% confidence level and $\\leq1.41$ at the 95\\% confidence level.  This result for the quiescent state agrees with a similar conclusion reached by Gallo et al. (2003) for the radio-emitting gas and Maccarone (2003) for the X-ray-emitting gas in the low hard state.  \\noindent 2. The source 4U1543--47 has only an upper limit on its quiescent X-ray luminosity.  Because of its low inclination angle ($i=20.7\\pm1.0\\degr$) this source is most sensitive to relativistic beaming from a jet.  We urge greater efforts to measure the quiescent luminosity of the source.  A detection at a level well below the present upper limit could be problematic for any jet model that invokes a reasonable velocity for the radiating gas.  \\noindent 3. 4U1543--47 showed a strong fluorescent iron line with an equivalent width of 300 eV in outburst in the high soft state.  Quantifying an argument by Park et al. (2004), we show that the X-ray continuum must have originated in gas with an outward bulk Lorentz factor $<1.04$ (Fig. 2).  If this gas was in a jet, then the jet must have been highly sub-relativistic.  Alternatively, the X-rays could have come from a slowly expanding corona (Beloborodov 1999; Malzac et al. 2001) or from the disk.  \\noindent 4. Of the 11 sources with estimates of $i$ (Table 1 plus the source XTE J1550--564), we show that there are none with $i>75\\degr$ or $\\cos i <0.25$.  The probability that this would happen by chance is only 4\\%, suggesting that there is a selection effect that prevents edge-on systems from being seen.  \\noindent 5. None of the 20 binaries in our complete sample is an eclipsing source.  The probability that it could happen by chance is only 0.5\\% (calculated based on the data given in Table 2).  This completely independent argument, which moreover does not require estimates of $i$, confirms that high inclination systems are hard to detect.  \\noindent 6. In our view, the most plausible explanation for the absence of eclipsing sources is that the accretion disks in black hole X-ray binaries are flared (or corrugated or warped) by about $\\pm15\\degr$. The disks thus permanently occult the X-ray emission for observers located within $\\pm15\\degr$ of the binary orbital plane.  Milgrom (1978) suggested the same explanation for the rarity of eclipses in neutron star X-ray binaries.  \\noindent 7. We show that there is a surprising dichotomy between the X-ray light curves of BHXN with $i \\lesssim 65\\degr$ and those with $i\\sim70-75\\degr$ (Fig. 3).  The former sources have smooth well-behaved outburst light curves with exponential decays, whereas the latter have noisy light curves, often with a deep minimum around the mid-point of the outburst.  There are similar indications also in the optical light curves.  \\noindent 8. We tentatively suggest that the systems with $i\\sim70-75\\degr$ have noisy light curves because they are close to being occulted by the disk.  Fluctuations in the flaring angle of the disk and/or the height of the X-ray emitting gas above the disk could cause the observed variations in the flux.  \\noindent 9. The source SAX J1819.3--2525 has the largest inclination angle among all our sources: $i=75\\pm2\\degr$.  It also has the most extreme X-ray light curve.  In this BHXN, we suggest that the primary X-ray source was occulted almost throughout its outburst, with only a small amount of radiation being observed from the outer regions of a large corona.  However, for a period of just two hours, the source became very bright in X-rays.  We suggest that during this brief episode there was a powerful jet ejection; X-rays from the jet at a relatively large distance from the black hole were observed directly without being obscured."
    },
    "0410/astro-ph0410283_arXiv.txt": {
        "abstract": "A substantial fraction of barred galaxies host additional nuclear bars which tumble with pattern speeds exceeding those of the large-scale (primary) stellar bars. We have investigated the mechanism of formation and dynamical decoupling in such nested bars which include gaseous (secondary) nuclear bars within the full size galactic disks, hosting a double inner Lindblad resonance. Becoming increasingly massive and self-gravitating, the nuclear bars lose internal (circulation) angular momentum to the primary bars and increase their strength. Developing chaos within these bars triggers a rapid gas collapse --- bar contraction. During this time period, the secondary bar pattern speed $\\Omega_{\\rm s} \\sim a^{-1}$, where $a$ stands for the bar size. As a result, $\\Omega_{\\rm s}$ increases dramatically until a new equilibrium is reached (if at all), while the gas specific angular momentum decreases --- demonstrating the dynamical decoupling of nested bars. Viscosity, and therefore the gas presence, appears to be a necessary condition for the prograde decoupling of nested bars. This process maintains an inflow rate of $\\sim 1~{\\rm M_\\odot~yr^{-1}}$ over $\\sim 10^8$~yrs across the central 200~pc and has important implications for fueling the nuclear starbursts and AGN. ",
        "introduction": "A large fraction, $\\sim 1/3$, of barred galaxies host a secondary (nuclear) bar in the central regions, in addition to the primary, large-scale stellar bar (Laine et al. 2002; Erwin \\& Sparke 2002).   Ironically, such nested bar systems have been predicted theoretically, including their most intriguing property --- the state of dynamical decoupling (Shlosman, Frank \\& Begelman 1989; Shlosman, Begelman \\& Frank 1990). Few examples of double barred galaxies have been observed already in the 1970s, but have been classified as having triaxial (elliptical) bulges rather than nuclear bars (de Vaucouleurs 1974; Sandage \\& Brucato 1979; Kormendy 1982). Both stellar and mixed type bars (e.g., Buta \\& Crocker 1993;  Shaw et  al. 1995;  Friedli et al. 1996; Jungwiert, Combes \\& Axon 1997; Jogee, Kenney \\& Smith 1998; Knapen, Shlosman \\& Peletier 2000), as well as gaseous nuclear bars (e.g., Ishizuki et al. 1990; Devereux, Kenney \\& Young 1992; Forbes et al. 1994; Mirabel et al. 1999; Kotilainen et al. 2000; Maiolino et al. 2000) have been detected. Random mutual orientation of the observed nested bars confirms that they spend a substantial fraction of their lifetime in the decoupled state --- tumbling with different pattern speeds.  Numerical simulations have shown indeed that pattern speeds of the mixed nuclear bars are substantially larger than those of the primary stellar bars (e.g., Friedli \\& Martinet 1993: Combes 1994; Heller \\& Shlosman 1994). However, purely stellar nuclear bars form only under special initial conditions (e.g., Friedli \\& Martinet 1993), while gaseous bars have been claimed not to decouple (Wada \\& Habe 1992; Combes 1994). A number of issues have never been resolved --- what triggers the dynamical decoupling of nested bars? Do purely gaseous self-gravitating bars form, and can they decouple from the primary bars? In this Letter we use numerical simulations of self-gravitating gas in galactic disks to show that gaseous bars form and decouple as a result of stellar bar-driven inflow into the central kpc. Furthermore, we study the physical reasons for this dynamical runaway in nested bar systems (see also Shlosman 2003; and in preparation). ",
        "conclusions": ""
    },
    "0410/hep-ph0410313_arXiv.txt": {
        "abstract": "We show that atmospheric electric fields may affect the cosmic ray observations in several ways and may lead to an overestimation of the cosmic ray energies.  The electric field in thunderclouds can be as high as a few kV/cm.  This field can accelerate the shower electrons and can feed some additional energy into the shower.  Therefore, ground array observations in certain weather conditions may overestimate the energy of ultrahigh-energy cosmic rays if they don't take this effect into account. In addition, the electric field can bend the muon trajectories and affect the direction and energy reconstruction of inclined showers.  Finally, there is a possibility of an avalanche multiplication of the shower electrons due to a runaway breakdown, which may lead to a significant miscalculation of the cosmic ray energy. ",
        "introduction": "Ultrahigh-energy cosmic rays (UHECR) produce extensive air showers in the atmosphere, which are observed by two different techniques: (1) optical detection of fluorescent light, and (2) surface array detection of charged particles in the shower.  The first technique, employed, for example, by HiRes detector \\cite{hires}, can be used exclusively in good weather, while the second technique, used, for example, by the AGASA experiment\\cite{AGASA}, is usually believed to afford accurate observations regardless of the atmospheric conditions.  The new Pierre Auger experiment is using a combination of the two techniques \\cite{AUGER}.   On a clear day the atmosphere is permeated by the electric fields of order a few volts per cm~\\cite{intro}.  However, in thunderclouds the electric field reaches much higher magnitudes, up to a few kV/cm~\\cite{el_book}. Balloon measurements~\\cite{Marshall} show the electric fields ${\\cal E}> 1$~kV/cm, at altitudes 0-12 km, which may switch polarity at several altitudes.  A shower developing in the atmosphere and going through a thundercloud may, therefore, pass between layers with either polarity at some angle.  The electric field along the path of a random shower may have a magnitude of a few kV/cm and an arbitrary direction, or alternating direction.  As we discuss below, the electron energy gain in the electric field of a few kV/cm can be comparable to the energy losses due to ionization.  It, therefore, important to examine what effect this may have on detection of UHECR.  Several experimental studies confirm the influence of electric fields on secondary particles in a shower.  First, observations of low-energy cosmic ray electrons show short-term variations during thunderstorms \\cite{thunder_low_energy}.  A similar effect has been observed for muons with a higher energy, $E>100$~MeV~\\cite{thunder_low_energy,thunder_low_energy2}, but it is not nearly as strong as the effect on the low-energy electrons.  Second, it has been established that giant electron-gamma bursts are triggered in thunderclouds by the passage of an extensive air shower (EAS) from a $10^{16}$ cosmic ray \\cite{electr_thunder}.  EAS were measured at the Tien Shan Mountain in coincidence with detection of a radio signal from the giant electron-gamma bursts. Theory of these bursts, based on the phenomenon of runaway breakdown, has been developed in recent years~\\cite{uspekhi}.  It has was also suggested that acceleration of secondary electrons from EAS in thunderclouds can be used for radio detection of cosmic rays with energies $E>10^{17-19}$~eV~\\cite{radio_high_energy}.   We will show that the atmospheric electricity associated with thunderclouds may affect the ground array observations of $E>10^{19-20}$~eV cosmic rays in several ways.  First, acceleration of the shower electrons can alter the low-energy spectrum and the number of particles in EAS at ground level. This could lead to errors in energy determination by a ground array. Second, the muon trajectories may be deflected by the atmospheric electric fields, which could lead to errors in reconstruction of energy and direction of inclined showers. Finally, a runaway breakdown and discharge triggered by an UHECR EAS may produce a dramatic increase in the number of electrons detected by the surface array. ",
        "conclusions": "We have shown that atmospheric electric fields may affect the energy measurements by the ground arrays in certain weather conditions. Our simplified calculations show that, in the presence of thunderclouds, one may overestimate the shower energy by as much as 20\\%. In addition, the deflection of muons may affect the directional reconstruction of inclined showers, which relies on muon arrival times. Finally, if an UHECR shower is accompanied by a runaway breakdown, the number of shower electrons  can increase dramatically, which may lead to a significant miscalculation of the cosmic ray energy."
    },
    "0410/astro-ph0410006_arXiv.txt": {
        "abstract": "We review several aspects of the Sunyaev-Zel'dovich (SZ) effect associated with the large scale baryon distribution and its characteristic signatures in the statistics of cosmic microwave background (CMB) anisotropies. We discuss (1) the contributions to the angular power spectrum from the thermal and kinetic SZ effects, (2) the effect of SZ non-Gaussianities on cosmological parameter estimation, and (3) the SZ-induced CMB polarization towards galaxy clusters. We also discuss the extent to which recent measurements of small angular scale CMB anisotropy can be accounted for by SZ clusters and the potential use of SZ polarization towards large samples of galaxy clusters as a probe of dark energy parameters. ",
        "introduction": "The Sunyaev-Zel'dovich effect (SZ; \\cite{SunZel80}) is the upscattering of cosmic microwave background (CMB) photons as they pass through the hot gas of a galaxy cluster.  This upscattering introduces frequency-dependent secondary temperature anisotropies in the CMB that are proportional to the integral of the electron pressure $\\Pi_e = n_e k_B T_e$ along  a given line of sight through the Universe: \\begin{equation} \\frac{\\Delta T^{\\rm SZ}}{T_{\\rm CMB}}\\equiv g_{\\nu}(x) y=g_{\\nu}(x)\\int \\, d\\eta a(\\eta) \\frac{\\sigma_{\\rm T}}{m_{\\rm e}}\\Pi_e(\\eta)  \\, . \\label{eqn:sz} \\end{equation} where we have defined $y$, the Compton $y$-parameter, $\\sigma_{\\rm T}$ is the Thomson cross-section, and $m_{\\rm e}$ is  the electron rest-mass \\footnote{Note that here and in what follows we adopt units where the speed of light $c$ is set to unity.}. The frequency dependence of the SZ anisotropies is accounted for by the function $g_{\\nu}(x)=x{\\rm coth}(\\frac{x}{2})-4$ with $x={h\\nu}{(k_B T_{\\rm CMB})^{-1}}$, which, in the low frequency Rayleigh-Jeans part of the spectrum has the limit $g_{\\nu} \\rightarrow -2$.  This form for the frequency spectrum leads to an apparent temperature decrement, relative to the primordial CMB black body spectrum, at frequencies below the null at $x_{\\rm cross}$, and an increment above  $x_{\\rm cross}$.  This SZ null occurs at approximately $217~{\\rm GHz}$ with a slight dependence on the electron temperature $T_{e}$ due to higher-order relativistic corrections to the inverse-Compton scattering process neglected in the Thomson approximation.  In Eq.~(\\ref{eqn:sz}), $\\eta \\equiv \\eta(z)$ is the conformal distance from the observer at redshift $z=0$ to the cluster at redshift $z$, given by \\begin{equation} \\eta(z) = \\int_0^z {dz' \\over H(z')} \\, , \\end{equation} where the expansion rate of cold dark matter (CDM) cosmological models with a cosmological constant $\\Lambda$ is \\begin{equation} H^2 = H_0^2 \\left[ \\Omega_m(1+z)^3 + \\Omega_K (1+z)^2 +\\Omega_\\Lambda \\right]\\, . \\end{equation} The Hubble distance today is $H_0^{-1} = 2997.9 \\, h^{-1} $Mpc with $h =0.72 \\pm 0.08$~\\cite{Freedman:2000cf, Spergel:2003cb}. $\\Omega_i$ denotes the contribution of component $i$ in units of the critical density $\\rho_{\\rm crit} \\equiv 3H_0^2/8\\pi G$, with $i=c$ for the CDM, $b$ for the baryons, and $\\Lambda$ for the cosmological constant. We also define the auxiliary quantities $\\Omega_m \\equiv \\Omega_c+\\Omega_b$ and $\\Omega_K \\equiv 1-\\sum_i \\Omega_i$, which represent the total matter density and the contribution of spatial curvature to the expansion rate, respectively.\\\\  The thermal SZ effect has been imaged in the CMB towards massive galaxy clusters whose presence is a priori known from optical data \\cite{Caretal96,Jonetal93}. The temperature of the scattering medium in certain clusters can reach up to 10 keV producing temperature changes in the CMB of order 1 mK at Rayleigh-Jeans (RJ) wavelengths. The individual galaxy cluster SZ images have a variety of astrophysical and cosmological applications including a direct measurement of the angular diameter distance to the cluster through a combined analysis with X-ray data and a measurement of the gas mass, and, thus the baryon fraction of the universe. We do not discuss these applications of the SZ effect in the present review, but rather will concentrate on the statistical study of wide-field CMB data where SZ effects lead to anisotropies in the temperature distribution both due to resolved and unresolved galaxy clusters. In fact, the thermal SZ contribution is the dominant signal beyond the damping tail of the primary anisotropy power spectrum.\\\\  In addition to the thermal SZ effect, the bulk flow of electrons in galaxy clusters and other virialized halos that scatter the CMB photons leads to a second contribution to temperature fluctuations, the kinetic SZ effect. This effect arises from the well known Doppler effect \\cite{Kai84,CooHu00}: \\begin{equation}\\label{equ:dop} \\frac{ \\Delta T^\\dop}{T_{\\rm CMB}}(\\bn) = \\int_0^{\\eta_0} d\\eta g(\\eta) \\bn \\cdot {\\bf v}_g(\\eta,\\bn \\eta)\\, , \\end{equation} where ${\\bf v}_g$ is the baryon velocity and $g(\\eta)$ is the visibility function (also interpreted as the probability of scattering within $d\\eta$ of $\\eta$) \\begin{equation}\\label{equ:visibility} g =  \\dot \\tau e^{-\\tau} = X_e(z) H_0 \\tau_H (1+z)^2 e^{-\\tau}\\,. \\end{equation} Here $\\tau(\\eta) = \\int_0^{\\eta} d\\eta' \\dot\\tau(\\eta')$ is the optical depth out to distance $\\eta$, $X_e(z)$ is the ionization fraction as a function of redshift, and \\begin{equation} \\tau_H = 0.0691 (1-Y_p)\\Omega_b h\\,, \\end{equation} is the optical depth due to Thomson scattering to the Hubble distance today, assuming full hydrogen ionization with primordial helium fraction of $Y_p (=0.24)$. From equation (\\ref{equ:dop}) follows the secondary temperature anisotropy due to the kinetic SZ effect. The kinetic SZ fluctuation is written as an integral over the product of the radial velocity, $\\bn \\cdot {\\bf v}_g$ , and the density fluctuation associated with the cluster, $\\delta_g$: \\begin{equation}\\label{equ:kSZ} \\frac{\\Delta T^\\ksz}{T_{\\rm CMB}}(\\bn) = \\int_0^{\\eta_0} d\\eta g(\\eta) \\bn \\cdot {\\bf v}_g(\\eta,\\bn \\eta) \\cdot \\delta_g(\\eta, \\hat{\\bf n} \\eta)\\, . \\end{equation}   The angular power spectrum of anisotropies generated by the Doppler effect peaks at angular scales corresponding to the horizon at the time of scattering. The effect cancels out significantly at scales smaller than the horizon since photons scatter against the crests and troughs of the perturbation. The kinetic SZ effect arises from the second-order modulation of the Doppler effect by non-linear density fluctuations associated with virialized objects such as galaxy clusters, and avoids the strong cancellation associated with the linear Doppler effect. The kinetic SZ effect is also known as the Ostriker-Vishniac effect \\cite{OstVis86} when the density fluctuations are associated with the linear density field of the large-scale structure. In addition to modulation of the velocity field due to density perturbations, the velocity field is also modulated by fluctuations in the ionization fraction of electrons in a partly reionized universe during the reionization epoch. This latter contribution is generally referred to as patchy-reionization \\cite{Aghanim:1996ib,Gruzinov:1998un,Santos:2003jb}. Due to the density weighting, the kinetic SZ effect peaks at small angular scales (sub arcminutes for $\\Lambda$CDM). For a fully ionized universe, contributions are broadly distributed in redshift so that the power spectra are moderately dependent on the optical depth $\\tau$.\\\\  In this review we focus on the statistical characterization of secondary CMB temperature fluctuations in the form of the angular power spectrum of SZ-induced anisotropies. SZ maps have important additional applications including number counts of halos where a comparison to analytical predictions can be used to derive cosmological parameters. These studies will not be the subject of this short review, but are discussed elsewhere in these proceedings.\\\\  The outline of this article is as follows. In \\S2 we derive the power spectrum of the thermal SZ effect using the halo model of large scale structure. \\S3 compares these theoretical predictions to the current CMB data. In \\S4 we derive the power spectrum of the kinetic SZ effect. Finally, we discuss secondary CMB polarization due to free electrons in galaxy clusters and its cosmological use as a powerful probe of dark energy in \\S5. ",
        "conclusions": ""
    },
    "0410/astro-ph0410230_arXiv.txt": {
        "abstract": "Short period sdB binaries with cool companions are crucial to understand pre-CV evolution, because they will evolve into cataclysmic variables, when the sdB will have left the extended horizontal branch. Recently we discovered the sixth such system, HS~2333+3927, consisting of an sdB star and an M dwarf (period: 0.172~d) with a very strong reflection effect, but no eclipses. The reflection is stronger than in any of the other similar systems which renders a quantitative spectral analysis very difficult because the Balmer line profiles may be disturbed by the reflected light. A spectroscopic analysis results in $ {T_{\\rm eff}}$ = 36\\,500\\,K, $\\log{g}$ = 5.70, and $\\log (n_{\\rm He}/n_{\\rm H}) = -2.15$. Mass-radius relations were derived from the results of the analysis of light and radial-velocity curves. Comparison with the mass-radius relation derived from the surface gravity of the sdB star favours a rather low mass of 0.38~\\mbox{$\\rm M_{\\odot}$}\\ for the primary. The mass of the companion is 0.29~\\mbox{$\\rm M_{\\odot}$}. HS~2333+3927 is the only known sdB+dM system with a period above the CV period gap. ",
        "introduction": "SdB binaries are important to clarify the evolutionary origin of sdB stars because the analysis of light and radial velocity curves can constrain their dimensions and masses. However, only five suitable systems are known up to now.  There is general consensus that the sdB stars can be identified with models for Extreme Horizontal Branch (EHB) stars (Heber 1986). Like all HB stars they are core helium burning objects. However, their internal structure differs from typical HB stars, because their hydrogen envelope is very thin ($<$1\\% by mass) and therefore inert. As a consequence EHB stars evolve directly to the white dwarf cooling sequence, thus avoiding a second red giant phase. How they evolve to the EHB configuration is controversial. The problem is how the mass loss mechanism of the progenitor manages to remove all but a tiny fraction of the hydrogen envelope at {\\em precisely} the same time as the He core has attained the mass ($\\sim$0.5~\\mbox{$\\rm M_{\\odot}$}) required for the He flash.  Considerable evidence is accumulating that a significant fraction of the sdB stars reside in close binaries (Maxted et al.~2001; Napiwotzki et al., 2004, see also Karl et al, these proceedings). Therefore mass transfer should have played an important role in the evolution of such binary systems. Detailed investigations of sdB binaries, in particular eclipsing systems, are crucial to determine their masses. However, only three such eclipsing binaries, HW~Vir, PG~1336-013, and HS~0705+6700 (see Drechsel et al. 2001) and two non-eclipsing ones are known up to now, which consist of an sdB star and an M dwarf companion revealed by reprocessed light from the primary (reflection effect). Recently, Heber et al. (2004) discovered another related system, HS~2333+3927, which is, however, not eclipsing, but otherwise possesses very similar system parameters and configuration. ",
        "conclusions": "\\subsection{Results}  From the analysis presented above we conclude that a model for a low mass ($\\approx$0.38~\\mbox{$\\rm M_{\\odot}$}) core helium burning sdB star and a 0.29~\\mbox{$\\rm M_{\\odot}$}\\ M-dwarf fits the observations best. However, the results have to be taken with a grain of salt. The spectral analysis of the Balmer lines, is plagued by reflected light, as is evident from the line profile variations observed in H$\\alpha$, but less obvious in other Balmer lines. An improved measurement of the gravity, therefore, is urgently needed.  Another important observational constraint would be a precise measurement of the projected rotational velocity. The rotation of the sdB star in HS~2333+3927  is very likely tidally locked to the orbital motion. Drechsel et al. (2001) showed this to be true for HS~0705+6700. Determining the projected rotation velocity would therefore allow an independent estimate of the inclination of HS~2333+3927.  The spectral lines of the Lyman series are sensitive gravity indicators and plenty of metal lines can be used to measure $v\\,\\sin{i}$. These measurements will allow us to constrain mass and radius much better  \\subsection{SdB stars and pre-CV Evolution}  Short period sdB binaries with main sequence companions, like HS~2333+3927, are important not only to understand the formation and evolution of sdB stars. When the sdB star will have left the EHB, it will evolve into a cataclysmic variable. Therefore, these objects are also crucial to understand pre-CV evolution. Schenker (these proceedings) has suggested that cataclysmic variables with short periods, i.e. below the CV period gap evolve from sdB + dM binaries. While four other known systems have periods between 1.7~h to 2.8~h, i.e. below the CV period gap, the period (4.13~h) of HS~2333+3927 is a bit above the CV period gap. Its orbital separation is so small that the secondary is appreciably distorted. The shrinkage of the binary orbit by gravitational wave radiation will initiate mass transfer turning the system into a cataclysmic variable. It is likely that the HS~2333+3927 system will then still have a period larger than the CV period gap. Hence sdB + dM binaries form an evolutionary channel to the population of longer period cataclysmic variables as well as to the short-period population."
    },
    "0410/astro-ph0410189_arXiv.txt": {
        "abstract": "Intergalactic Ly$\\alpha$ opacity suggests that H~I was reionized at $z \\approx 6$, while He~II reionization was delayed to $z \\approx 3$. Both epochs are in conflict with inferences from CMB optical depth (WMAP) which suggest early reionization at $z = 10-20$. One of the major contributions of \\FUSE\\ to cosmological studies has been the detection of He~II Ly$\\alpha$ (``Gunn-Peterson\") absorption in the spectra of AGN at redshifts $z \\geq 2.03$.  Spectra of He~II absorbers, taken in concert with corresponding H~I (Ly$\\alpha$) lines, allow us to fix the epoch of helium reionization at $z_{\\rm HeII} \\approx 2.8 \\pm 0.2$.  Here, I review \\FUSE\\ observations of He~II absorption, together with their implications for the sources and transport of ionizing radiation over 30-50 Mpc distances through the IGM.  \\FUSE\\ observations of He~II absorption toward HE~2347-4342, combined with {\\it Keck} and {\\it VLT} observations of H I, are consistent with photoionization by QSOs, with a wide range of intrinsic spectral indices, and modified by filtering and reprocessing in the IGM.  The He~II/H~I ratio ($\\eta$) exhibits variations over 1 Mpc distance scales ($\\Delta z \\approx 10^{-3}$).  Intriguingly, this $\\eta$-ratio is also correlated with Ly$\\alpha$ filaments and voids. The ionizing radiation field appears to be softer (higher He~II/H~I) in the voids.  These void regions may be ionized by local soft sources (dwarf starburst galaxies), or the QSO radiation may softened by escape from AGN cores and transport through denser regions in the ``cosmic web\". The differences in ionizing spectra may explain the 1.4 Gyr lag between H~I and He~II epochs. ",
        "introduction": "Before beginning a discussion of the \\FUSE\\ contributions to the study of He~II reionization, I provide an overview of the cosmological significance of ``reionization\". This term refers to the sudden increase in ionization fraction of hydrogen gas in the intergalactic medium (IGM), initiated at redshifts $z \\geq 6$ by the first sources of ionizing radiation.  After the epoch of recombination at $z_{\\rm rec} \\approx 1300$, the IGM entered a long period of ``dark ages\" when the IGM was mostly neutral.  This period ended sometime before $z \\approx 6$, perhaps as early as $z = 10-20$.  The ionization state of the IGM affects the transmission of radiation through the Ly$\\alpha$ absorbers (H~I and He~II), as well as the electron-scattering optical depth to the cosmic microwave background (CMB) radiation.  As individual ionizing sources (stars, galaxies, quasars) began emitting ionizing radiation, regions of ionized gas were produced, preceded by ionization fronts (I-fronts) in the IGM.  These isolated ionized zones (H~II regions, He~III regions) eventually overlapped (Gnedin 2000), after which the IGM became transparent to Ly$\\alpha$ radiation.   Residual absorption continued as a ``picket fence\" of discrete, redshifted Ly$\\alpha$ absorbers, often termed the ``Ly$\\alpha$ forest\" -- see spectra in Songaila (1998).  The rapid increase in IGM ionization occurred as a topological phase change in the volume filling factor of ionized zones. The longer the first ionizing sources were on, the greater the volume of the zones. In order to effectively absorb all the Ly$\\alpha$ radiation at $z = 6$ requires a neutral fraction of just $x_{\\rm HI} = 2 \\times 10^{-4}$, even in gas that is underdense by a factor of 10 relative to the mean.  Simulations of this reionization process (Gnedin 2000, 2004) show that the transition from neutral to ionized is quite extended in time, between $z = 5-10$.  The first stage (pre-overlap) involves the development and expansion of the first isolated ionizing sources. The second stage marks the overlap of the I-fronts and the disappearance of the last vestiges of low-density neutral gas.  Finally, in the post-overlap stage, the remaining high-density gas is photoionized from the outside.  The final drop in neutral fraction and increase in photon mean free path occur quite rapidly, over $\\Delta z \\approx 0.3$ (see Fig.\\ 1 of Gnedin 2004).  Reionization is tied closely to the onset of galaxy formation beginning at $z \\geq 20$.  For the ionization of He~II, and somewhat for H~I, reionization depends on the rise in the numbers and luminosities of active galactic nuclei from $z = 5 \\rightarrow 2$. According to numerical simulations (e.g., Ricotti, Gnedin, \\& Shull 2002, herafter denoted RGS), the earliest epochs of galaxy formation began at $z \\approx 20-30$ triggered by H~I and H$_2$ cooling of protogalactic halos with virial temperatures $T_{\\rm vir} = 10^3$~K to $10^{4}$~K.  RGS found that the average star formation rate density, $\\rho_{\\rm SFR}$, increased monotonically from $10^{-3}$ to $10^{-1}$ $M_{\\odot}~{\\rm yr}^{-1}~{\\rm Mpc}^{-3}$ between $z = 20$ and 10, according to the approximate formula: \\begin{equation} {\\rm SFR} = (0.003~M_{\\odot}~{\\rm yr}^{-1}~{\\rm Mpc}^{-3}) 10^{(20-z)/5} \\; . \\end{equation}  The effects of reionization can be seen indirectly in the electron-scattering optical depth of the cosmic microwave background (CMB), as we discuss in the next section. However, the most direct observations of the IGM ionization state come from their influence on the Ly$\\alpha$ line-blanketing opacity of the IGM at wavelengths shortward of the H~I and He~II Ly$\\alpha$ lines at $(1215.670~{\\rm \\AA})(1+z)$ and $(303.782~{\\rm \\AA})(1+z)$, respectively. As the amounts of H~I or He~II increase, these line-blanketed regions grow into black (``Gunn-Peterson\") absorption troughs, whose presence can be used to identify high-redshift galaxies (Steidel et al.\\ 2002).  The best probes of the epoch of reionization have come from analyzing the weak transmission of AGN continuum flux in these troughs, as the high-redshift IGM transforms from a neutral to ionized medium.  The He~II absorption becomes detectable in the far ultraviolet at redshifts $z \\geq 2.03$ with \\FUSE\\ and at $z \\geq 2.79$ with HST.  From the latest Ly$\\alpha$ data, the reionization epochs occur at $z \\approx 6.1\\pm0.3$ for H~I (Becker \\etal\\ 2001; Fan \\etal\\ 2002; Gnedin 2004) and at $z \\approx 2.8\\pm0.2$ for He~II (Kriss \\etal\\ 2001; Smette \\etal\\ 2002; Shull \\etal\\ 2004; Zheng \\etal\\ 2004). ",
        "conclusions": ""
    },
    "0410/astro-ph0410140_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec::intro}  The ``Big Bang'' model is misnamed, as one might expect when a term is coined as an insult. Cosmologist Fred Hoyle first used the name in a BBC radio interview to denigrate the theory, which opposed the then-dominant Steady State model. The name Big Bang stuck, as did its evocation of a mighty explosion in space. But the image of an \\emph{explosion} is highly misleading. What the model actually posits is that the Universe began hot, dense, and expanding.  Within the first second, roughly 13.7 billion years ago, the Universe achieved temperatures on the order of 1 trillion Kelvin (K, degrees above absolute zero; Schwarz, 2003). The density during that second was high enough to stop neutrinos, which interact so weakly with matter that they can pass unmolested through a quadrillion kilometers of lead. What ties this hot, dense beginning to the Universe we see today is expansion. A useful metaphor for the expanding universe is the surface of an inflating balloon. As the balloon inflates, space-time itself is stretched; every point moves away from every other point. Density falls as the universe expands. If you picture a wave oscillating over the surface of the balloon, the wavelength increases. Increasing the wavelength of light corresponds to reducing its temperature. The Universe thus cools as it expands.  Within the first three minutes, the Universe's temperature was over one billion K. The energy density in space was so high that atoms could not form. Space was filled with a stew of photons, baryons (e.g., protons and neutrons), electrons, neutrinos, and other matter. As the temperature cooled below 1 billion K, light-element nuclei (deuterium, helium, some lithium) formed as well, in proportions that fit well with observations. During this period, photons (radiation) were the dominant form of energy in the Universe. Any fluctuations in density caused by gravity (which affects light and matter) were quickly smoothed out and so could not grow.  When the temperature of the primordial photons had fallen below approximately 12,000 K, photons were no longer dominating the interactions among all particles. Photons and baryons became coupled in a mathematically perfect fluid, while exotic kinds of matter began to clump under the influence of gravity. The interaction between this photon-baryon fluid and such gravitational overdensities are of critical importance and will be described below.  When the temperature reached about 3000 K, roughly 380,000 years after the Big Bang, electrons and protons could combine to form atoms. This decoupled the photon-baryon fluid, and the photons flew free through space. This period is named \\emph{recombination} and happened, in cosmic terms, very quickly. After another 200 million years, hydrogen formed after recombination had clumped enough for the first stars to form, which began the synthesis of heavy elements and the formation of galaxies that we see today.  Most of the photons released at recombination have travelled through space for billions of years without interacting with matter. The temperature of these primordial photons has now cooled to about 2.7K, barely above absolute zero, which puts them in the microwave part of the electromagnetic spectrum. This primordial radiation field, which still pervades the Universe, is called the Cosmic Microwave Background (CMB). The CMB thus provides a snapshot of the moment of recombination, and fluctuations in the temperature across the sky contain information about the physics of the early universe.  \\subsection{The Cosmic Microwave Background Radiation}   As we will explain in the remainder of this section, the temperature fluctuations in the CMB give a snapshot of the physics in the early Universe and provide critical tests of cosmological models. In 1992, the Cosmic Microwave Background Explorer (COBE) satellite discovered fluctuations in the blackbody temperature of the CMB (Smoot et al. 1992). These fluctuations are small: approximately one thousandth of the mean temperature over the sky. Indeed, almost thirty years of experiments since the CMB's discovery could not detect any deviation from uniformity. During the ten years following COBE, many more refined measurements were taken; notable experiments include MAXIMA, DASI, BOOMERANG (Lee et al. 2002, Halverson et al. 2002, Netterfield et al. 2002). In 2003, the Wilkinson Microwave Anisotropy Project (WMAP) considerably refined the picture, increasing spatial resolution by a factor of 33 and sensitivity by a factor of 45 over COBE (Bennett et al. 2003). In Figure \\ref{fig::cobeWMAP}, we compare the COBE and WMAP temperature sky maps after removing the mean temperature $T=2.726$ Kelvin and adjusting for the motion of our galaxy through the Universe. The fluctuations' magnitudes are just right to explain the large-scale structure in the Universe we see today. For example, if they had been much smaller, there would not be enough local concentration of mass to seed the formation of galaxies, galaxy clusters, et cetera.  \\begin{figure}[t] \\vbox { \\psfig{file=cobe_wmap.epsf,height=5.5in,angle=0} } \\caption{(top) The CMB as seen by the COBE satellite. The angular resolution of the satellite is about $10^{\\circ}$ and the various shades correspond to hot and cool spots with respect to the CMB blackbody temperature. (bottom) The CMB from the Wilkinson Microwave Anisotropy Probe. Notice the high angular resolution. Also notice that the large-scale structures are apparent in both the COBE and the WMAP data. Image courtesy of the WMAP Science Team and available at the WMAP Mission website: {\\tt http://map.gsfc.nasa.gov}.} \\label{fig::cobeWMAP} \\end{figure}  Perhaps the most important summary of the temperature measurements used by cosmologists is the power spectrum, which gives the temperature variance as a function of spatial frequency. The spectrum's shape, and in particular the location and height of its peaks, relates directly to the parameters in cosmological models. (See Appendix 2 for a description of these parameters.) Thus, a critical statistical question is how accurately can these peaks be estimated. Of particular interest are the height and location of the first peak and the relative heights of the successive peaks.  Figure \\ref{fig::cosmospec} displays an estimated spectrum commonly used by cosmologists and highlights the peaks of interest. We will give a more precise definition of the spectrum in Section \\ref{sec::spectrum}, but here we want to explain how the spectrum's shape relates to the physics in the time up to recombination.  \\begin{figure}[t] \\vbox { \\psfig{file=cosmospec-gray.epsf,height=5.5in,angle=0} } \\caption{Estimated CMB spectrum showing the three peaks of interest. The underlayed boxes give the data ranges and uncertainties from a variety of older CMB experiments, not including WMAP. From Hu (2001).} \\label{fig::cosmospec} \\end{figure}   A key to understanding the physics before recombination is, as mentioned earlier, that photons and baryons became coupled into a (perfect) fluid. Mathematical techniques for studying fluid dynamics apply well in this scenario and have been investigated by many authors (see for instance Hu and Sugiyama 1995; Hu 1999, 2001, 2003; Hu and Dodelson 2002). The properties of the fluid are determined by the relative density of photons and baryons in the fluid. Photons provided pressure, and the baryons provided inertia. As the fluid falls into a gravitational potential well around a clump of higher density, the pressure from the photons resists compression and the inertia of the baryons increases it. (Large, isolated potential wells were likely rare in the early universe; instead, there were random density fluctuations at many scales.) The result is an oscillation that produces pressure waves -- sound -- in the photon-baryon fluid. These \\emph{acoustic oscillations} account for much of the interesting structure in the spectrum, particularly the size and arrangement of peaks. The imprint of those waves remains in the CMB as a pattern of hot and cold spots.  To understand the peaks in the power spectrum, it is helpful to decompose the acoustic oscillations into their basic components, or modes. The first peak of the spectrum represents the fundamental tone of the oscillations, and the other peaks in the spectrum represent harmonics of this tone.  The fundamental corresponds to the mode for which one compression occurs between the Big Bang and recombination. Each successive harmonic corresponds to an additional half-cycle, compressions followed by rarefication (decompression). Thus, the second peak represents modes that had time to compress and then rarefy before the photons were released from the photon-baryon fluid. The third peak represents compression-rarefication-compression, and so on.  The height of the first peak is determined by the total energy density.  Roughly, with more matter, the gravitational attraction requires more force to counteract, deepening the compression and thus increasing the amplitude of oscillation.  Now suppose we increase the density of baryons in the photon-baryon fluid.  This increases the inertia of the fluid, deepening each compression phase without changing the rarefication. The oscillations become asymmetric. What this means is the odd-numbered peaks, whose modes end on a compression, are enhanced relative to the even-numbered peaks, whose modes end on a rarefication.  Thus as the baryon fraction increases we should (over some range) see a differential effect on the odd and even numbered peaks.  The third peak in the spectrum provides the clearest support for the existence of ``dark matter''  -- a substance of unknown composition that interacts at most weakly with baryons (e.g., neutrons, protons) or with photons (that's why its dark). To see why, it is illuminating to compare the oscillations in two example cases. In the ``radiation-dominated era,'' when photons were the dominant form of interaction in the universe, density fluctuations were short-lived and unstable.  A compressed region of photon-baryon fluid would rarefy as described earlier, but as it did so, the overdensity that caused the original gravitational well would disappear.  Thus, in this case, at most one cycle of oscillation would occur between the Big Bang and recombination.  We would see only one small peak in the temperature power spectrum corresponding to the mode (component of oscillation) that reaches maximum compression at the time of recombination.  When the matter fraction is low, the peak would be small, increasing with the baryon fraction (inertia).  In the ``matter-dominated era,'' however, most of the energy density was in the form of dark matter.  The rarefication phase of the oscillation would not eliminate the local overdensity, allowing multiple cycles of oscillation.  The result is a spectrum with multiple harmonics and thus multiple peaks. The existence and contribution of dark matter is only distinguishable from that of baryons alone with three or more peaks. Moreover, the magnitude of the third peak constrains the time of transition between a radiation and matter dominated universe. In particular, a finding that the second and third peaks were comparable in magnitude would suggest that dark matter dominated before recombination, which is a fundamental prediction of Big Bang cosmology. The magnitude of the third peak is also of interest for estimating the fraction of dark matter in the Universe. Astronomers have several methods for inferring the dark matter fraction (e.g., studying the rotation of galactic disks in the recent Universe), and it is vital to determine if these estimates are comparable to those produced by the physics of the early Universe.  Finally, the pattern of CMB hot and cold spots we see on the sky corresponds to those photons just reaching us from the moment of recombination. (Recombination was relatively quick but not instantaneous, so there is some blurring of high spatial frequencies from the scatter of photons during that finite period.) The contribution to this pattern from each acoustic mode maps to a spherical mode of fluctuations on the sky. The analysis then proceeds by decomposing the observed fluctuations into spherical modes and using the contributions of these modes to understand the acoustic oscillations. We discuss this in the next subsection.   \\subsection{The CMB Temperature Power Spectrum} \\label{sec::spectrum}  Our focus in this paper is inference about the CMB temperature power spectrum and in particular the peaks in the spectrum. In this section, we describe the spectrum and some of the issues that arise in estimating it. Marinucci (2004, this issue) gives a more complete derivation upon which ours is based.  Let $T(\\theta,\\vartheta)$ denote the temperature field as a function of colatitude (zero at the zenith) $0\\le \\theta\\le \\pi$ and longitude $0\\le\\vartheta < 2\\pi$. Let $\\overline{T}$ denote the average of the temperature field over the sphere.  Define the temperature fluctuation field by $$ Z(\\theta,\\vartheta) = \\frac{T(\\theta,\\vartheta) - \\overline{T}}{\\overline{T}}. $$ Note that $Z$ is a random field with mean zero and is assumed to have finite second moment. We can expand $Z$ in terms of a orthonormal basis on the sphere. The usual choice of basis is the set of spherical harmonics $\\{Y_{\\ell,m}(\\theta,\\vartheta)\\}$, for positive integers $\\ell=1,2, \\ldots$ and integers $-\\ell \\leq m\\leq  \\ell$. (Here $\\ell$ is called the multipole index, or loosely ``multipole moment.'') These are defined as follows: $$ Y_{\\ell,m}(\\theta,\\vartheta) = \\sqrt{\\left(\\frac{ 2\\ell + 1}{4\\pi}\\right)\\frac{ (\\ell-|m|)!}{ (\\ell+|m|)!}}\\; P_\\ell^{|m|}(\\cos \\theta)\\, e^{i m \\vartheta}, $$ where the $P_{\\ell,m}$ $\\ell = 1, 2, \\ldots$ and $m = 0,\\ldots,\\ell$ are the associated Legendre functions defined by \\begin{eqnarray*} P_\\ell^m(x) &=& (-1)^m (1-x^2)^{m/2} \\frac{d^m}{d x^m} P_\\ell(x)\\\\ \\noalign{\\noindent with Legendre polynomials} P_{\\ell}(x) &=& \\frac{1}{2^\\ell \\ell!} \\frac{d^\\ell}{d x^\\ell}(x^2-1)^\\ell. \\end{eqnarray*}  We can now write \\begin{equation} Z(\\theta,\\vartheta) = \\sum_{\\ell=1}^\\infty \\sum_{m=-\\ell}^{\\ell} a_{\\ell,m}Y_{\\ell,m}(\\theta,\\vartheta) , \\end{equation} where, \\begin{equation} a_{\\ell,m} = \\int_{0}^{2\\pi}\\int_{0}^\\pi Z(\\theta,\\vartheta) Y_{\\ell,m}(\\theta,\\vartheta) \\sin \\theta d \\theta d\\vartheta. \\end{equation} Since $Z$ is a mean zero random field, the coefficients $a_{\\ell,m}$ are random variables. They have mean 0, variance $$ C_\\ell \\equiv \\E |a_{\\ell,m}|^2, $$ and are uncorrelated. The \\emph{power spectrum} is defined to be $C_\\ell$ as a function of $\\ell$.  Usually, it is assumed that $Z$ is a Gaussian field (but see the paper by Marinucci in this issue) which implies that the $a_{\\ell,m}$ have a Gaussian distribution. If we were to observe $Z$ without measurement error, we could estimate $C_\\ell$ by, say, \\begin{equation}\\label{eq::Cellhat} \\tilde{C}_\\ell = \\frac{1}{2\\ell +1}\\sum_{m=-\\ell}^{\\ell} a_{\\ell,m}^2, \\end{equation} and thus for large $\\ell$ we have $\\tilde C_\\ell \\approx C_\\ell$ because we are averaging a large number of $a_{\\ell,m}^2$. We call $\\tilde{C}_\\ell$ the realized spectrum. Another important implication of equation (\\ref{eq::Cellhat}) is that even with perfect observations, we would not know the true power spectrum. Because our Universe is viewed as one realization of a stochastic process, $\\tilde C_\\ell$ will in general differ from $C_\\ell$, especially for small $\\ell$. This is known as the problem of \\emph{cosmic variance}. We return to this point in Section \\ref{sec::discussion}.  In practice, the data are subject to various sources of measurement error, blurring, and unobserved parts of the sky. For example,the Milky Way, which is relatively bright, obscures the deep sky along a wide band. The spherical harmonics are no longer orthogonal over what is left of the sphere, which induces correlation and bias into the estimated $C_\\ell$s. There are in addition a host of other complications in measuring $Z$.  Our model, in vector form, is \\begin{equation} \\label{eq::linearModel} \\hat{C} = C + \\epsilon, \\end{equation} where $\\hat{C}$ is the \\emph{observed spectrum} and the noise vector $\\epsilon$, with covariance matrix $C \\aleph C^T$, incorporates the known sources of error, including measurement error. If there were no sky cut for the galaxy, $\\aleph$ would be diagonal, but in practice, it incorporates the various known sources of error. In practice, the unknown $C$ in the covariance matrix is replaced by a a pilot estimate, $C^0$. The choice of $C^0$ turns out to have surprisingly little effect on the results. We thus take the covariance matrix of $\\epsilon$ in equation (\\ref{eq::linearModel}) to be known and equal to $\\Sigma = C^0 \\aleph (C^0)^T$.  Another issue is that the observations are actually derived from a convolution of the $\\tilde C_\\ell$s with $\\ell$-dependent window functions; that is, the model is actually $\\hat C = K C + \\epsilon$ for some matrix $K$. But as Figure \\ref{fig::window} shows, the rows of $K$ are very nearly delta functions. (See Knox, 1999.) And in fact, incorporating these window functions has negligible effect on our results, so we disregard them in what follows.  \\begin{figure}[t] \\vbox { \\hglue -0.5in\\psfig{file=window-functions.eps,height=2.5in,angle=0} } \\caption{Bandpower windows from the matrix $K$ centered on (left to right) $\\ell=100,200,300,400$.} \\label{fig::window} \\end{figure} ",
        "conclusions": "\\label{sec::discussion}  \\subsection{Findings}  Our most striking finding is that the center of our nonparametric confidence ball using the WMAP data alone lies very close to the Concordance model fit over the range where the data are not noise dominated. Recall from equation (\\ref{eq::concord_like}) that the Concordance model incorporates four independent data sets, each with distinct likelihood forms. In contrast, the parametric model fit using only the WMAP data (with likelihood $\\cL_{\\ssrm WMAP}$) lies barely in the 73\\% confidence ball. Given that cosmologists often use 68\\% confidence intervals as their standard of evidence, this is tantamount to a rejection of the cosmological model that underlies that parametric fit.  This raises two points. First, it is remarkable that with a fully nonparametric method we have come very close to the Concordance model based on the WMAP data alone.  Second, that we obtained basically the same spectrum as the Concordance model calls into question the accuracy of the WMAP-only likelihood $\\cL_{\\ssrm WMAP}$.  \\subsection{Methods}  We have presented a nonparametric method for analyzing the CMB spectrum.  Our techniques have wide applicability to regression problems beyond cosmology.  By starting with a confidence ball, then probing the ball using functionals, one can address a variety of questions about the unknown function while maintaining correct coverage, despite multiplicity and post-hoc selection.  The method in this paper modifies the original Beran-D\\\"umbgen construction to account for heteroskedasticity. This modification yields a substantial reduction in the size of the confidence set. The resulting confidence set is also more useful in that it leads to tighter (looser) bounds in regions where the function is more (less) accurately measured.  One advantage of our approach, is that it allows one to separate the information in the data from the information in a model.  Although we did not pursue the full calculation here, we could intersect our confidence ball with the manifold of spectra from the parametric model as a way to combine data and model.  Specifically, we could use the cosmological model to generate spectra, but then test which spectra are consistent with the data by reference to our confidence ball.  This does not rely on likelihood asymptotics which, as we discuss below, are suspect in this problem.  Another advantage is that by extending this analysis to a constrained noparametric model (such as a three peak model) that contains the cosmological model, we can make the same inferences without being tied to the analytic form of the model. Our approach can then be used to check the model, make inferences under the model, and compare parametric to nonparametric inferences.  We should point out that cosmologists obtain confidence intervals for parameters in their (11-dimensional) model by integrating over the nusiance parameters and producing a marginal posterior.  However, the likelihood is ill-behaved, under-identified, and degenerate.  Moreover, in the physics literature, there does not seem to be a clear appreciation of the fact that interval estimates obtained this way need not have correct frequentist coverage.  There are several other advantages to our approach.  If a parameter is under-identified this will show up automatically as a wide confidence interval.  The intervals have correct asymptotic coverage and simultaneous validity over all parameters of interest.  There is no need to integrate or profile the likelihood function.  Finally, the asymptotic theory for (\\ref{eq::conf-ball}) is insensitive to the fact that the standard asymptotics for the likelihood approach fail.  \\subsection{Inferential Foundations}  Interestingly, there seems to be some confusion about the validity of frequentist inference in cosmology.  Since we have access to only one Universe -- and thus cannot replicated it --  some feel that it makes no sense to make frequentist inferences.  This represents a common misunderstanding about frequentist inference in general and confidence intervals in particular.  The frequency statements for confidence intervals refer to the procedure, not the target of the inference.  Our method for constructing confidence balls traps the true function 95 percent of the time, even over a sequence of different, unrelated problems. There is no need to replicate the given experiment, or Universe.  Complicating matters is the fact that the coefficients $a_{\\ell,m}$ of the temperature field are random and unknown.  To see the importance of this point, it is useful to make a finer distinction by defining the realized spectrum $\\tilde{C}_\\ell = (1/(2\\ell+1))\\sum_\\ell | a_{\\ell,m}|^2$, the ``true spectrum'' $C_\\ell=\\E(\\tilde{C}_\\ell )$ and the measured spectrum $\\hat{C}_\\ell$.  Note that all our inferences have actually been directed at the realized spectrum.  Some phyciststs find it disturbing to be making frequentist inferences about $\\tilde{C}_\\ell$ since it is a realization of a random variable rather than a parameter in the usual sense.  But this is no different than making inferences about a random effect in a standard random effects model.  These confusions have led to an interesting movement towards Bayesian methods in cosmology.  Of course, when used properly, Bayesian methods can be very effective. Currently however, the Bayesian interval estimates in the physics literature seem questionable, being based on unfettered use of marginalizing over high-dimensional, degenerate likelihoods using flat priors chosen mainly for convenience.  Indeed, an active area of research is finding corrections for such intervals to make them have correct coverage.  Moreover, the potentially poor coverage of the Bayesian interval seems not to have been widely recognized in the Physics literature."
    },
    "0410/astro-ph0410695_arXiv.txt": {
        "abstract": "I present preliminary results for a sample of $\\sim$700 red giants in $\\omega$~Cen, observed during the Ital-FLAMES Consortium GTO time in May 2003, for the Bologna Project on $\\omega$~Cen. Preliminary Fe and Ca abundances confirm previous results: while the metal-poor and intermediate populations show a normal halo $\\alpha$-enhancement of [$\\alpha$/Fe]$\\simeq$+0.3, the most metal-rich stars show a significantly lower [$\\alpha$/Fe]$\\simeq$+0.1. If the metal-rich stars have evolved within the cluster in a process of self-enrichment, the only way to lower their $\\alpha$-enhancement would be SNe type Ia intervention. ",
        "introduction": "The treatment of the large data volumes that can be obtained with large telescopes and multi-objects spectrographs, such as FLAMES at the ESO VLT, requires a major upgrade of the methods commonly used to reduce and analyze high and medium resolution spectra.  We are therefore developing a set of routines for an automatic or semi-automatic abundance analysis of stellar spectra based on equivalent widths (EW). The first product is DAOSPEC, a code developed by P. B. Stetson for automatic EW measurement (http://cadcwww.hia.nrc. ca/stetson/daospec/). The preliminary abundance analysis presented here is the first step of an iteractive and automatic procedure under development at the Bologna Observatory. ",
        "conclusions": ""
    },
    "0410/astro-ph0410376_arXiv.txt": {
        "abstract": "{  The transition from hadronic matter to quark matter in the core of neutron stars is likely to be associated with the appearance of a mixed phase, leading to a smooth variation of the star density profile. We discuss the results of a systematic study of the properties of the mixed pase upon Coulomb and surface effects. A state of the art nonrelativistic equation of state of nuclear matter has been used for the low density phase, while quark matter has been described within the MIT bag model, including the effect of perturbative one-gluon exchange interactions. The implications for neutron star structure are discussed. ",
        "introduction": "\\label{intro}  The possible occurrence of a core of deconfined quark matter in the interior of neutron stars has been investigated by a number of authors over the past three decades (for a recent review see, e.g., Weber \\cite{weber}). Due to the complexity of the underlying dynamics, theoretical approaches to the study of hybrid stars largely rely on models to describe the equation of state (EOS) of strongly interacting matter, in both the hadronic and quark sector, as well as on a set of assumptions on the nature of the phase transition.  Most calculations have been carried out using nuclear EOS obtained from either nonrelativistic nuclear many body theory (NMBT) or relativistic mean field theory (RMFT), while the deconfined phase is usually described within the MIT bag model (Chodos et al \\cite{bagmodel}).  In their pioneering work Baym and Chin (\\cite{baym}) employed the familiar Maxwell double tangent construction (see, e.g., Huang \\cite{huang}), which amounts to assume that the transition occur at constant pressure. Within this picture, charge neutral nuclear matter at energy density $\\epsilon_{NM}$ cohexists with charge neutral quark matter at energy density $\\epsilon_{QM}$, the two phases being separated by a sharp interface.  In the early 90s Glendenning (Glendenning \\cite{glend1,glend2}) first pointed out that the requirement that the two phases be individually charge neutral is in fact too restrictive. In a more general scenario charged nuclear and quark matter may share a common lepton background, thus giving rise to a mixed phase extending in space over a sizable fraction of the star.  The appearance of a mixed phase strongly affects the macroscopic properties of the star. A transition at constant pressure necessarily leads to the appearance of a discontinuity in the density profile, i.e. to a star consisting of a inner core of quark matter at energy density $\\epsilon_{QM}$ surrounded by nuclear matter at energy density $\\epsilon_{NM}$. On the other hand, the mixed phase allows for a smooth variation of the energy density, leading in turn to a smooth variation of the star density profile.  Whether the transition proceeds at constant pressure or according to Glendenning's picture depends upon i) the value of the Debye screening length, driving charge separation, and ii) the amount of electrostatic and surface energy needed for the formation of the structures of quark and nuclear matter in the mixed phase. If screening is too strong to allow for a uniform lepton density, or if the energy loss due to Coulomb and surface effects exceeds the gain in bulk energy, the standard scenario predicted by Maxwell construction turns out to be favorable.  The results of a detailed study carried out by Heiselberg et al (\\cite{HPS93}), suggest that the mixed phase is energetically favored over a sizable density range if the surface tension is less that $\\sim 70$ MeV/fm$^2$. In their calculations these authors adopted the MIT bag model (Chodos et al \\cite{bagmodel}) for the quark matter EOS, while nuclear matter was described using a somewhat oversimplified model, in which the energy-density included a quadratic compressional term and a symmetry term taken from a previous work of Lattimer et al (\\cite{compress}).  Over the past decade, the availability of new nucleon-nucleon potentials, resulting from accurate fits to nucleon-nucleon scattering data, and the improvement of the computational schemes made it possible to obtain a new generation of EOS within the framework of NMBT (Akmal \\& Pandharipande \\cite{apr}, Akmal et al \\cite{apr2}). In view of the fact that NMBT is a parameter free approach, whose dynamics is strongly cosntrained by nuclear data, and has been shown to possess a highly remarkable predictive power in theoretical studies of few nucleon systems (Pieper \\& Wiringa \\cite{pw}), it provides a natural candidate to describe neutron star matter in the nuclear phase.  In this paper we extend the work of Heiselberg et al (\\cite{HPS93}) carrying out a systematic study of the stabilty of the mixed phase. We adopt the state of the art EOS of nuclear matter, obtained by Akmal et al (\\cite{apr2}) within NMBT, and analyze the dependence of the results on i) the parameters entering the MIT bag model EOS, employed to describe quark matter, and ii) the value of the surface tension, driving both Coulomb and surface effects.  In Section \\ref{EOS} we summarize the main features of the model EOS of both nuclear and quark matter, while the implementation of Gibbs conditions in the case of two chemical potentials, leading to the appearance of the mixed phase, is discussed in Section \\ref{phase_transition}. Our main results are presented in Section \\ref{surface}, devoted to the role of Coulomb and surface energy. The implication of the appearance of the mixed phase for neutron star structure are outlined in Section \\ref{star}. Finally, the conclusion of our work are stated in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have carried out a study of the transition from nuclear matter to quark matter in the inner core of neutron stars, aimed at assessing whether the appearance of a mixed phase is energetically favorable and how the emerging picture depends upon the parameters entering the MIT bag model EOS.  In order to minimize the uncertainty associated with the description of the nuclear matter phase, we have adopted a EOS obtained from an {\\it ab initio} calculation, based on a dynamical model stongly contrained by experimental data and not involving any adjustable parameters (Akmal et al \\cite{apr2}).  Our results show that the effect of perturbative gluon exchange on the MIT bag model EOS is large and cannot be accounted for adjusting the value of the bag constant $B$. Neglecting interactions between quarks leads to a considerable softening of the EOS, resulting in a drastically lower transition density.  The fact that the setting $\\alpha_s=0$ lead to a much softer EOS is reflected by the rather small value of the maximum neutron star mass, barely exceeding the canonical value of 1.4 $M_\\odot$.  The stability of the mixed phase turns out to be strongly affected by surface and Coulomb effects. Using the softer quark matter EOS ($\\alpha_s=0, B=120$ MeV/fm$^3$) we find that even a very small value of the surface tension, $\\sigma = 2$ MeV/fm$^3$, produces a narrowing of the density region spanned by the mixed phase. With the harder EOS ($\\alpha_s=0.5, B=200$ MeV/fm$^3$) the cohexistence of neutral phases of nuclear and quark matter is energetically favored at all densities for $\\sigma = 10$ MeV/fm$^3$.  Comparison with the results of Heiselberg et al suggests that surface and Coulomb effect become larger when a realistic EOS is employed to describe the nuclear matter phase.  While the mass-central density relation appears to be largely unaffected by the occurrence of the mixed phase, its effect can be clearly seen in the $M(R)$ curve. However, as measurements of neutron star radii are plagued by large uncertainties, this feature is not likely to be exploitable to extract clearcut information from observations.  The most striking difference between the Maxwell and Gibbs picture of the phase transition shows up in the neutron star density profile, which in the case of transition at constant pressure exhibits a sharp discontinuity.  The presence of a density jump is known to affect neutron star dynamics, leading to appearance of a class of nonradial oscillation modes, called $g$-modes, associated with emission of gravitational radiation.  Early investigations of the $g$-modes focused on the discontinuities produced by the changes of chemical composition in the low density region of the neutron star crust, corresponding to a fractional distance from the surface $\\lsim$~10~\\% (Finn \\cite{finn}, McDermott \\cite{mcdermott}). These studies have been recently extended to the case of $g$-modes produced by a discontinuity located at much larger density and involving a much larger density jump, such as those associated with the transition to quark matter (Miniutti et al \\cite{gmodes}). Based on the results of calculations carried out using siple polytropic EOS, Miniutti et al (\\cite{gmodes}) argue that a simultanuoeus measurements of the frequencies of the fundamental $f$-mode and the $g$-mode would provide information on both size and location of the discontinuity.  Although it is unlikely that the first generation of laser interferometric antennas will detect gravitational waves emitted by an oscillating  neutron star, the new detectors currenltly under investigation (see, e.g., the EURO proposal (\\cite{EURO})) are expected to be much more sensitive at the relevant frequencies above 1-2 kHz. Hopefully, information on the neutron star matter EOS and the nature of the transition to quark matter may come gravitational wave astronomy."
    },
    "0410/astro-ph0410648_arXiv.txt": {
        "abstract": "Abundances of R Coronae Borealis stars (RCBs) and Extreme Helium stars (EHes) are discussed. Recent estimates of the $s$-process elements in hot extreme helium stars show enhancements of lighter $s$-process elements (Y, Zr) relative to heavier $s$-process elements, a characteristic shared by RCB stars. It also suggests that atleast some EHe stars went through an episode of $s$-process element synthesis in their earlier evolution. A majority of RCB stars show a high $^{12}$C/$^{13}$C ratio in their atmospheres. A recent analysis of the spectrum of minority RCB star V\\,CrA, however, shows a lower value between 4 $-$ 10. The implications  of these results are discussed. ",
        "introduction": "It is a great pleasure to be able to participate in this celebration   honouring Prof. David Lambert. I would like to convey to David greetings and   best wishes from his numerous friends and admirers in India. It is almost a hundred years since Ludendorff discovered the absence of Balmer lines in the  spectrum of the F type irregular variable R\\,CrB, the first recognised hydrogen deficient star. About 35 years later, the first helium star was  discovered by Popper.  The basic problem about the origins of these stars, namely,  a single (or double) star of intermediate mass becoming  a high luminosity hydrogen deficient star, still remains  a mystery. Presently, there are two principal hypotheses to account for their origins. The prime contender is the white dwarf merger scenario (or double degenerate, DD) in which a helium white dwarf secondary is accreted onto a CO (or He) white dwarf primary resulting in the ignition of a helium shell in the accreted envelope. The energy generated forces the star to expand to cool supergiant dimensions. The life time at this stage depends on the supply of helium in the shell and the mass of the envelope. The subsequent evolution would be like a canonical post-AGB contraction to the white dwarf track (Webbink 1984, Iben \\& Tutukov 1984, Saio \\& Jeffery 2000). The second scenario proposed is called the late thermal pulse (or final flash, FF), in which an AGB star in its final stages retains a helium shell of enough mass  as to get ignited for a last time after it descends on to the white dwarf track. The final thermal pulse forces the star to become a cool supergiant for a second time. The subsequent evolution would again resemble the canonical post-AGB evolution of single stars (Iben et al. 1983). Presently, there are no decisive observational tests predicted to choose either one (or any other alternative) of the scenarios.  Several groups of hydrogen deficient stars have now been identified ranging from high luminosity (low surface gravity) cooler HdC stars, RCBs, EHes to low luminosity (high gravity) hot PG\\,1159 stars (see Jeffery 1996 for a display of various groups in the log $g$, log $T_{\\rm eff}$ plane). It is not clear how these various groups of hydrogen deficient stars are related to each other $-$ e.g., EHes and RCBs. Do they represent sucessive stages in evolution? Even some groups are further subdivided and show diversity in properties e.g., majority and minority RCBs which show [S/Fe], [Si/Fe] values of about 0.5 and much higher, respectively (Lambert \\& Rao 1994, Asplund et al. 2000). Do they suggest different paths of evolution? In the last few years, stars like Sakurai's object and FG\\,Sge provided us examples of stellar evolution in real time by turning from a normal star to hydrogen deficient born-again giant (Asplund et al. 1997, Gonzalez et al. 1998, Duerbec et al. 2002). How are these born-again giants related to RCBs are some of the intriguing questions that the theory of intermediate mass star evolution is unable to provide. In keeping with the theme of the conference,`abundances as records of stellar evolution and nucleo- synthesis', I would like to explore in this presentation the interrelationships of RCBs and EHes and also address some of the issues raised above by seeking clues from surface chemical compositions. ",
        "conclusions": ""
    },
    "0410/astro-ph0410362_arXiv.txt": {
        "abstract": "The source of fluorine is not well understood, although core-collapse supernovae, Wolf-Rayet stars, and asymptotic giant branch stars have been suggested.  A search for evidence of the $\\nu$ process during Type II supernovae is presented.  Absorption from interstellar F~{\\small I} is seen in spectra of HD~208440 and HD~209339A acquired with the {\\it Far Ultraviolet Spectroscopic Explorer}. In order to extract the column density for F~{\\small I} from the line at 954 \\AA, absorption from H$_2$ has to be modeled and then removed.  Our analysis indicates that for H$_2$ column densities less than about $3 \\times 10^{20}$ cm$^{-2}$, the amount of F~{\\small I} can be determined from $\\lambda$954. For these two sight lines, there is no clear indication for enhanced F abundances resulting from the $\\nu$ process in a region shaped by past supernovae. ",
        "introduction": "While fluorine consists of the single stable isotope $^{19}$F, theories of stellar nucleosynthesis have not yet been able to identify clearly the origin of this element.  Three astrophysical sites are currently suggested as possible significant sources -- core-collapse supernovae (SN~II), Wolf-Rayet stars, and asymptotic giant branch (AGB) stars. In SN~II, the enormous flux of neutrinos escaping the core interact with $^{20}$Ne in the Ne-shell, causing evaporation of a proton or neutron and (in both cases) conversion of a small amount of $^{20}$Ne to $^{19}$F (Woosley et al. 1990).  This neutrino nucleosynthesis is often called the ``$\\nu$ process''. In Wolf-Rayet stars, $^{19}$F is produced inside the star mainly at the beginning of the He-burning phase; this fluorine will be destroyed later in the evolution of the star unless extensive mass loss removes the star's layers down to the regions where $^{19}$F has been synthesized and then ejected into space (Meynet \\& Arnould 2000).  For AGB stars, Jorissen, Smith, \\& Lambert (1992) suggested that $^{19}$F is a by-product of neutron captures during He-burning; Mowlavi, Jorissen, \\& Arnould (1998) provided a detailed description of the process. In our effort to understand the source of fluorine, we present the first measurements of interstellar F with the {\\it Far Ultraviolet Spectroscopic Explorer} ($FUSE$) via absorption from F~{\\small I} $\\lambda$954.  There are very few observations of F outside the solar system.  Observations of HF via rotational-vibrational lines in the infrared (Jorissen et al. 1992) are limited to the coolest giants. More recently, Cunha et al. (2003) measured HF absorption from red giants in the Large Magellanic Cloud and $\\omega$ Centauri.  They found that the F/O abundance ratio declined as the O abundance declined, a trend expected for $\\nu$-process synthesis of F but Wolf-Rayet stars remain a possibility. Liu (1998) obtained the only definitive fluorine abundance in a planetary nebula (NGC 4361) via [F {\\small IV}] $\\lambda$4060.  Interstellar F {\\small I} was detected through the line at 954 \\AA\\ toward $\\delta$ Sco by Snow \\& York (1981), who used the $Copernicus$ satellite; the measured equivalent width of 10.9$^{+4.5}_{-3.1}$ m\\AA\\ translates into an abundance, log~$\\epsilon$(F) = log~(F/H) + 12, of $\\sim$ 4.0.  The solar system abundance is 4.48 $\\pm$ 0.06 from meteorites and 4.60 $\\pm$ 0.30 from HF lines in sunspots (Anders \\& Grevesse 1989). Neufeld et al. (1997) detected HF absorption at 121.7 $\\mu$m with the $ISO$ satellite from interstellar gas in Sgr~B2, a Giant Molecular Cloud near the Galactic Center.  Chemical models (see also Zhu et al. 2002) indicate that most of the fluorine in molecular clouds should be in HF, while analysis of HF excitation suggests severe fluorine depletion assuming the solar system abundance is the appropriate measure for comparison.  The results on the abundance of interstellar F from HF are consistent with those for HCl in this cloud (Zmuidzinas et al. 1995), another hydrogen halide.  However, the apparent extreme levels of depletion for both species and the unknown reference for elemental abundances in the vicinity of the Galactic Center make it hard to interpret these results from a cosmochemical perspective. Of the suite of observations, only those of Jorissen et al. (1992) provide clear evidence for a site of F production: thermally pulsing AGB stars that also make $s$-process nuclides.  This paper describes the detection of F~{\\small I} absorption toward HD~208440 and HD~209339A, which are about 900 pc from the Sun (e.g., Pan et al. 2004a).  These stars were chosen for study because they belong to the Cepheus OB2 Association, an ideal site to search for $\\nu$-process products.  The stars in Cep OB2 formed in an expanding shell of gas created by previous members of the cluster NGC~7160 when their lives ended in SN~II (Patel et al. 1995; 1998). Absorption from H$_2$ is also present at 954 \\AA\\ and must be removed for an accurate measurement of F~{\\small I}; our analysis provides a limiting H$_2$ column density for new F~{\\small I} $\\lambda$954 searches.  The focus of our interpretation is on abundance ratios because elemental abundances are less accurately known, a result of the relatively large uncertainties in total hydrogen column densities.  In particular, we examine F/Cl, N/O, and F/O ratios.  Column densities for Cl~{\\small I} and {\\small II} and for N~{\\small I} come from $FUSE$ spectra, while absorption from O~{\\small I} and another strong line of Cl~{\\small I} comes from data acquired with STIS on the {\\it Hubble Space Telescope} ($HST$).  The F/Cl ratio is useful in ascertaining the presence of an anomolous F abundance because both halogens are expected to have similar levels of depletion onto interstellar grains.  The ratio of two volatile elements, N and O, is useful for probing the overall level of depletion.  With the information gleaned from the F/Cl and N/O ratios, we can explore what the F/O ratios say about the site of F nucleosynthesis. ",
        "conclusions": "We presented the first interstellar fluorine measurements with $FUSE$ and described the methodology to extract the F column density.  The strongest F~{\\small I} line at 954 \\AA\\ is blended with lines from H$_2$; a prescription for extracting information on F~{\\small I} from the blend was given.  In the process, we determined that the line at 954 \\AA\\ can be measured as long as $N$(H$_2$) was less than about $3 \\times 10^{20}$ cm$^{-2}$.  The inferred F/O ratio suggests that F is more depleted than O, consistent with expectations from condensation temperatures.  While our goal to find a clear signature for the $\\nu$ process from F measurements was not achieved in these first observations, we offered suggestions for furthering the search.  Additional F~{\\small I} observations are planned and determinations of the B/F ratio offer possibly the best means for seeking evidence for the $\\nu$ process in interstellar space."
    },
    "0410/astro-ph0410681_arXiv.txt": {
        "abstract": "{ The solar photospheric carbon abundance has been determined from [C\\,{\\sc i}], C\\,{\\sc i}, CH vibration-rotation, CH A-X electronic and C$_2$ Swan electronic lines by means of a time-dependent, 3D, hydrodynamical model of the solar atmosphere. Departures from LTE have been considered for the  C\\,{\\sc i} lines. These turned out to be of increasing importance for stronger lines and are crucial to remove a trend in LTE abundances with the strengths of the lines. Very gratifying agreement is found among all the atomic and molecular abundance diagnostics in spite of their widely different line formation sensitivities. The mean of the solar carbon abundance based on the four primary abundance indicators ([C\\,{\\sc i}], C\\,{\\sc i}, CH vibration-rotation, C$_2$ Swan) is ${\\rm log} \\, \\epsilon_{\\rm C} = 8.39 \\pm 0.05$, including our best estimate of possible systematic errors. Consistent results also come from the CH electronic lines, which we have relegated to a supporting role due to their sensitivity to the line broadening. The new 3D based solar C abundance is significantly lower than previously estimated in studies using 1D model atmospheres. ",
        "introduction": "Carbon is a crucial element for the emergence of life in the Universe as we know it. Furthermore, it is one of the most abundant elements in the cosmos behind only hydrogen, helium and oxygen in this respect. While it is clear that carbon is produced through the $3\\alpha$-reaction (Burbidge et al. 1957) in stellar interiors, a debate regarding the exact site of its production is still ongoing. Recently both low- and intermediate mass stars as well as massive stars have been advocated to be the primary sources of the element (e.g. Chiappini et al. 2003; Akerman et al. 2004). Carbon also plays a key role in the physics of the interstellar medium for example through its propensity to form dust (e.g. Dopita \\& Sutherland 2002; Lodders 2003).  In view of its rather special status in astrophysics, it is clearly of particular importance to ascertain the traditional abundance anchorage point, namely the solar carbon abundance. In the standard reference of Anders \\& Grevesse (1989) the solar carbon abundance is quoted as ${\\rm log} \\, \\epsilon_{\\rm C} = 8.56 \\pm 0.04$, which was subsequently revised to ${\\rm log} \\, \\epsilon_{\\rm C} = 8.52 \\pm 0.06$ by Grevesse \\& Sauval (1998) based on a re-analysis of various atomic and molecular lines using a temperature-modified Holweger-M\\\"uller (1974) 1D semi-empirical model atmosphere. Holweger (2001) preferred the slightly higher value of ${\\rm log} \\, \\epsilon_{\\rm C} = 8.59 \\pm 0.11$ based only on permitted C\\,{\\sc i} lines in the optical and infrared (IR). In sharp contrast to these results, Allende Prieto et al. (2002) found the much lower value ${\\rm log} \\, \\epsilon_{\\rm C} = 8.39 \\pm 0.04$ through a detailed study of the forbidden [C\\,{\\sc i}] line at 872.7\\,nm employing a state-of-the-art 3D time-dependent hydrodynamical model of the solar photosphere (Asplund et al. 2000b). This latter value is in good agreement with the evidence from nearby B stars (Sofia \\& Meyer 2001), local interstellar medium (Meyer et al. 1998; Andr\\'e et al. 2003), the solar corona and the solar wind (Murphy et al. 1997; Reames 1999), in particular in view of the recent downward revision of the solar oxygen abundance (Allende Prieto et al. 2001; Asplund et al. 2004, hereafter Paper IV).  While the [C\\,{\\sc i}] 872.7\\,nm line has a well-determined transition probability and should be robust against departures from LTE, its relative weakness and the possible presence of weak blends may cast some doubt on the reliability of the low carbon abundance derived by Allende Prieto et al. (2002). We here present a comprehensive study of other available carbon diagnostics, including permitted atomic and various molecular lines using the same 3D hydrodynamical solar model atmosphere. Our results corroborate the findings of Allende Prieto et al. and unambiguously point to a significant revision by $-0.17$\\,dex of the solar carbon abundance compared with the often-adopted value given in Anders \\& Grevesse (1989). It is noteworthy that only with a 3D analysis are the primary atomic and molecular abundance criteria in excellent agreement, which further strengthens our conclusions. ",
        "conclusions": "We have presented a comprehensive study of the available atomic and molecular lines to derive a new solar carbon abundance which is significantly lower than most previous analyses: ${\\rm log} \\, \\epsilon_{\\rm C} = 8.39 \\pm 0.05$ (C/H\\,$= 245 \\pm 30 \\cdot 10^{-6}$). The quoted uncertainty includes our best estimate of the remaining systematic errors. Our primary abundance indicators are the forbidden [C\\,{\\sc i}] 872.7\\,nm, high-excitation permitted C\\,{\\sc i}, CH vibration-rotation and C$_2$ Swan lines with a supporting role played by CH A-X electronic lines. A special effort has been made to select the most accurate and reliable atomic and molecular data such as transition probabilities and dissociation energies. For the C\\,{\\sc i} lines detailed non-LTE calculations in a 1D context have been performed, which revealed significant departures from LTE in the line formation. Finally, a novel feature of the present study is the use of a realistic 3D hydrodynamical solar model atmosphere, which furthermore has proved to be crucial in order to obtain consistent abundances between the various diagnostics. In sharp contrast, the 1D Holweger-M\\\"uller model atmosphere yields much higher abundances for the molecular lines than for the atomic lines, partly due to the higher temperatures in the line-forming regions and partly due to the neglect of temperature inhomogeneities. The excellent concordance between the various transitions with widely different temperature and pressure sensitivities is a very strong argument in favour of both the 3D model atmosphere as such and the new low carbon abundance.  Further support for the new low solar carbon abundance advocated herein comes from a comparison of the chemical composition in related environments. The revised solar photospheric carbon abundance is now in good agreement with those measured in nearby B stars with solar iron abundances: ${\\rm log} \\, \\epsilon_{\\rm C} = 8.28 \\pm 0.17$ or C/H\\,$= 190 \\pm 90 \\cdot 10^{-6}$ (Sofia \\& Meyer 2001). Furthermore, the solar carbon abundance is now in good agreement with estimates of the total (gas plus dust) local interstellar medium carbon abundance for realistic gas-to-dust ratios (Andr\\'e et al. 2003; Estiban et al. 2004). Finally, coupled with the previously re-determined solar photospheric oxygen abundance (Paper IV), the photospheric C/O ratio is $0.54 \\pm 0.09$, which is in good agreement with the measurements in solar flares ($0.42 \\pm 0.09$, Fludra et al. 1999; $0.54 \\pm 0.04$, Murphy et al. 1997) and solar wind particles ($0.47 \\pm 0.01$, Reames 1999)\\footnote{Since both carbon and oxygen are high ionization potential elements, to first order it is expected that the solar coronal C/O ratio is unaffected by the so-called first ionization potential (FIP) effect, which introduces a differential depletion of elements with low ionization potential ($\\chi_{\\rm ion} \\la 9$\\,eV) for as yet unexplained reasons, see review by Raymond (1999).}.  The new solar carbon abundance corresponds to a change of $-0.28$\\,dex relative to that advocated in Lambert (1978) and $-0.17$\\,dex compared with the preferred value in the standard source of Anders \\& Grevesse (1989). This change is primarily the outcome of improved models of the solar photosphere and the line formation processes, as well as more accurate atomic and molecular line data. The possibility of similar substantial systematic errors in standard 1D abundance analyses still today is certainly worth keeping in mind when interpreting the results of other abundance studies in terms of stellar nucleosynthesis and galactic evolution."
    },
    "0410/astro-ph0410154_arXiv.txt": {
        "abstract": "\\noindent We present results from an analysis of a \\chandra{} observation of the cluster of galaxies Abell~3581. We discover the presence of a point-source in the central dominant galaxy which is coincident with the core of the radio source PKS~1404-267. The emission from the intracluster medium is analysed, both as seen in projection on the sky, and after correcting for projection effects, to determine the spatial distribution of gas temperature, density and metallicity. We find that the cluster, despite hosting a moderately powerful radio source, shows a temperature decline to around $0.4~\\mysub{T}{max}$ within the central 5~kpc. The cluster is notable for the low entropy within its core. We test and validate the \\prog{xspec} \\texttt{projct} model for determining the intrinsic cluster gas properties. ",
        "introduction": "\\label{intro}  Abell~3581 is a nearby ($z\\sim 0.0218$), richness class 0, cluster of galaxies. The central dominant galaxy IC~4374 has an extensive optical emission-line filament system (\\citealt{DanzigerFocardi88}, \\citealt{JFN87}) like that seen around the central galaxy in more massive clusters (e.g. the Perseus cluster, \\citealt{Conseliceetal01}).  Abell~3581 was previously studied in X-rays, using data from the \\rosat{} and \\asca{} satellites by \\citet{Johnstoneetal98} who found that it had a cool intracluster medium (ICM) with $kT\\sim 2$~keV, and a short radiative cooling time throughout its core. Those data admitted a classical cooling flow with a significant amount of intrinsic absorption.  Recently, (e.g. \\citealt{Petersonetal01}, \\citealt{Johnstoneetal02}, \\citealt{Petersonetal03}, \\citealt{Kaastraetal04}, \\citealt{Sandersetal04}) it has been shown that clusters which have short central cooling times have much less gas at temperatures below about a third of the ambient cluster temperature compared with that expected from simple cooling. This inconsistency between theory and observations has led to the proposal of many different mechanisms to reheat the cooling gas.  One popular subset of these heating models appeals to energy input from the active nucleus in the central dominant galaxy, which is usually found to harbour a radio source (\\citealt{Burnsetal97}).  The central dominant galaxy in Abell~3581 hosts the powerful radio source PKS~1404-267. Since this is a cool cluster with a powerful radio source at its centre, it might be expected that if active galactic nuclei (AGN) are able to heat the ICM in clusters, this source might show the effects more clearly than richer clusters. However, although we will show that the ICM does not have significant amounts of gas below a temperature of $kT\\sim0.8$~keV, there seems to be no other clear evidence for heating of the ICM by the AGN.  We note that Abell~3581 has a similar mean temperature and luminosity to the cluster Abell~1983 which was recently studied with \\xmm{} by \\citet{PrattArnaud03}. The properties of the inner core of Abell~3581 are however quite different from those of Abell~1983. ",
        "conclusions": "So far we have considered heating in Abell~3581 only from the point of view of balancing the cooling from the ICM in the core of the cluster. However, there is a more global property of clusters that shows that heating has been important. The observed luminosity/temperature ($L_x/T_x$) relation for clusters, where $\\mysub{L}{x}\\propto \\mysub{T}{x}^3$ (e.g. \\citealt{ArnaudEvrard99}) rather than $\\mysub{T}{x}^2$ as expected if only gravitational heating is involved (\\citealt{Kaiser86}), indicates that some form of extra heating has taken place. The level corresponds to about 2-3~keV per particle (e.g. \\citealt{Wuetal00}).  Low temperature clusters and groups show a large dispersion in the $\\mysub{L}{x}/\\mysub{T}{x}$ relation (\\citealt{HelsdonPonman00}), which may mean that although heating is widespread, it is not the same in all objects. One possible origin for the heating, which has been suggested by many authors, is the radio source which is often associated with the central galaxy.  We have measured the unabsorbed X-ray luminosity in the central 60~kpc of both the Perseus cluster and Abell~3581 (excluding the point sources) and find values of $1.5\\times10^{44}\\ergps$ and $1.4\\times10^{43}\\ergps$ respectively in the 0.3-10~keV band, so the Perseus cluster is factor of ten brighter than PKS~1404-267 in this region. By comparison, the 1.4~GHz radio luminosity of the inner two radio lobes in the Perseus cluster is $2.0\\times10^{24}\\WHz^{-1}$ while the 1.4~GHz radio luminosity of the lobes in PKS~1404-267 is $2.5\\times10^{23}\\WHz^{-1}$, showing that the distributed radio power in the Perseus cluster is also a factor of ten brighter than in PKS~1404-267.  Abell~3581 is therefore particularly interesting as it is a low temperature cluster with a strong radio source where heating seems not to have been effective, even close to the active nucleus. \\citet{Roychowdhuryetal04} have suggested that heating by the AGN should also affect the entropy profile of a cluster beyond the coolest parts. Even if the radio source has offset radiative cooling in the very innermost regions, it is clear that it has not had a strong influence on the gas entropy at $0.1\\mysub{R}{200}$, compared with similar low temperature clusters.  Various explanations for this are possible. a) The radio source may be very young compared with the radiative cooling time of hundreds of Myrs, in which case radiative cooling would have dominated in the past. However, no rapidly cooling cluster has yet been found elsewhere. b) The major heating phase in clusters may have occurred 5 or more billion years ago, beyond $z\\sim1$, when the central active galaxy or galaxies were much more important. This may have been a variable process which has left some clusters without significant heating. c) Perhaps Abell~3581 is particularly old and thus dense, and has had few recent mergers. Cooling in its high density core may absorb all the kinetic energy produced by the present active nucleus so that little is available for the rest of the gas. d) The dense core may have been introduced into the cluster as a result of a merger, or e) The dense cool core may have caused the moderately powerful radio source by more effectively fuelling the central black hole. Radio source heating as inferred for the Perseus cluster (by the dissipation of sound waves; \\citealt{Fabianetal03}) may also apply, and scale, to Abell~3581.  The core of Abell~3581 appears to be similar to that of the Perseus cluster, except that it is about an order of magnitude less luminous and has a temperature that is about one third of that in the Perseus cluster. Somehow the core has avoided the heating, or at least the effects of heating, common to most other low temperature clusters and groups."
    },
    "0410/astro-ph0410224_arXiv.txt": {
        "abstract": "{ We discuss the possibility that astrophysical accretion disks are dynamically unstable to non-axisymmetric disturbances with characteristic scales much smaller than the vertical scale height. The instability is studied using three methods: one based on the energy integral, which allows the determination of a sufficient condition of stability, one using a WKB approach, which allows the determination of the necessary and sufficient  condition for instability and a last one by numerical solution. This linear instability  occurs in any inviscid stably stratified differential rotating fluid for rigid, stress-free or periodic boundary conditions, provided the angular velocity $\\Omega$ decreases outwards with radius $r$. At not too small stratification, its growth rate is a fraction of  $\\Omega$. The influence of viscous dissipation and thermal diffusivity on the instability is studied numerically, with emphasis on the case when $d \\ln \\Omega / d \\ln r =-3/2$ (Keplerian case). Strong stratification and large diffusivity are found to have a stabilizing effect. The corresponding critical stratification and Reynolds number for the onset of the instability in a typical disk are derived. We  propose that the spontaneous generation of these linear modes is the source of turbulence in disks, especially in weakly ionized disks. ",
        "introduction": "The simplest model of an accretion disk is that of a barotropic, axisymmetric rotating shear flow in hydrostatic vertical equilibrium, with a Keplerian velocity law. Realistic disks are also subject to baroclinic effects (when the rotation departs from cylindrical, i.e. when $\\Omega$ varies also with the axial coordinate $z$) and to a vertical stratification, induced either by the hydrostatic state or via the illumination of the surface of the disk due to the central object. If the stratification is unstable, it leads to turbulence via convective instability. When the stratification is stable, it is generally ignored or thought to be unimportant in the stability analysis, under the rationale that it can only {\\sl stabilize} the flow. Ignoring the stratification and baroclinic effects makes the accretion disk look like a simple differentially rotating shear flow, with an azimuthal keplerian angular velocity profile $\\Omega( r )\\propto r^{-3/2}$ . Its linear stability with respect to axisymmetric disturbances is governed by the Rayleigh criterion in the inviscid limit: \\EQ \\frac{d(r^2 \\Omega)^2}{dr}>0 \\quad \\hbox{for stability.} \\label{rayleigh} \\EN Flows obeying this criterion are called {\\sl centrifugally stable}. Keplerian flow, in which angular momentum increases outwards, falls into this category. Yet, there is observational evidence that astrophysical (putatively Keplerian) disks are turbulent ({\\it cf.} Hersant et al. 2003), and thus that a source of instability exists in these flows.\\  Leaving aside baroclinic effects, various mechanisms have been found able to destabilize centrifugally stable flows. They may or may not apply to astrophysical disks.  \\par i) Centrifugally stable flows can experience a globally subcritical bifurcation (Dauchot \\& Manneville 1997), induced by finite amplitude disturbances involving non-linear mechanisms not captured by the Rayleigh criterion (Dubrulle 1993). The transition threshold in this case is related to the amplitude of the external disturbance, as typically observed in a plane shear flow~ (Dauchot \\& Daviaud 1994 ).  Taylor-Couette experiments, with fluid sheared between two concentric rotating cylinders in the centrifugally stable regime, have revealed such transition. When the inner cylinder is at rest, the data of Wendt (1933) and Taylor (1936), re-analyzed by Richard and Zahn (1999), show that for the typical amplitude of the intrinsic disturbances of the experimental devices, turbulence subcritically sets in as soon as $ R=U d/\\nu > R^{nl}_c $, where $U$ is the relative velocity between the walls, $d$, the radial extent of the flow (the gap) and $\\nu$ the viscosity.  For small curvature, the threshold is essentially independent of the gap/radius ratio and in the limit of plane Couette flow, $R^{nl}_c \\simeq 1500$, a value in agreement with the minimal Reynolds number for which turbulence can be induced in plane Couette flow. For large curvature, $ R^{nl}_c $ scales with the square of the gap/radius ratio.  Finally, novel laboratory experiments were performed recently (Richard 2001), to refine the present analysis and to explore regimes with co-rotating cylinders. For the first time, the hysteretic behavior of the transition with the inner cylinder at rest has been described, a signature of subcriticality and turbulent regimes were detected when the angular velocity decreases outward, in the centrifugally stable region.\\   ii) Centrifugally stable flows can also be destabilized by compressibility effects via non-axisymmetric instabilities (Papaloizou \\& Pringle 1984). This instability occurs independently of boundary conditions as long as $d\\ln\\Omega/d\\ln r<-\\sqrt{3}$ (Papaloizou \\& Pringle, 1985; Glatzel, 1987). However, the existence of such an instability for Keplerian disks requires the presence of at least one sharp edge (Goldreich \\& Narayan, 1985). This condition may be unrealistic in standard accretion disks.\\   \\par iii) Centrifugally stable flows can be further destabilized by adding another restoring force, which acts as a  {\\sl catalyzer}. A first example is a vertical magnetic field (Velikhov 1959; Chandrasekhar 1960), which renders such flows unstable provided they are anticyclonic, i.e. if $d (\\Omega)^2 /dr <0$. The application of this mechanism to disks was first discussed by Balbus and Hawley (1991); they showed that  the stratification of the disk does not modify the result, and  that the maximal growth rate of instability in that case is of the order of the angular rotation velocity in the inviscid limit. The instability of a rotating flow subject to vertical magnetic field includes a surprising paradox: experimentally, it was found that in liquid metals, in the centrifugally unstable case, the magnetic field inhibits instabilities, and thus, has a {\\sl stabilizing} influence (Chandrasekhar 1960). In the centrifugally stable case, it creates an instability. Summarizing, it appears that a stabilizing factor acts upon a stable flow so as to generate instability.   iv) A similar behavior was observed in rotating flows subject to a vertical stable stratification (Whithjack \\& Chen 1974, Boubnov and Hopfinger, 1995). In the {\\sl centrifugally unstable} case, the stratification enhances the stability of the flow and tends to increase the critical Reynolds number of the inner cylinder (Fig. \\ref{fig:figure1}). In the {\\sl centrifugally stable} regime, the experimental stability curve crosses the critical line for the Rayleigh criterion, and turbulence sets in via a non-axisymmetric mechanism. This instability is present in the small gap and wide gap regime, showing that curvature effect do not play a role in the instability mechanism. A possible theoretical explanation for this instability has been given recently by Molemaker et al. (2001) and Yavneh et al (2001). They performed an analytical and numerical stability analysis of a rotating flow in the presence of a stable vertical temperature gradient, and discovered the existence of a linear non-axisymmetric instability for all anticyclonically sheared flows. In their work, they interpret this instability as being caused by the interaction of two edge modes through a mechanism of arrest and phase locking along the boundaries by the mean shear flow.\\   \\begin{figure}[hhh] \\centering \\includegraphics[width=8.5cm]{figure1.eps} \\caption[]{Neutral-stability curves in  a stratified Taylor-Couette flow experiment for different stratifications, in the domain $(R_2,R_1)$ where $R_2$ and $R_1$ are the Reynolds numbers based on the gap size and the velocity at the outer or inner cylinder. a) In the wide gap limit (ratio of inner cylinder to outer cylinder equal to $0.2$), experiment of Whithjack and Chen (1974): $N=1.5 s^{-1}$ (circles), $N=1.25 s^{-1}$ (diamonds) and $N=0.88s^{-1}$ (triangles); b) in the small gap limit (ratio of inner cylinder to outer cylinder equal to $0.8$), experiment of  Boubnov and Hopfinger (1995) $N=1.21 s^{-1}$ (circles), $N=0.89 s^{-1}$ (diamonds),  $N=0.54 s^{-1}$ (triangles) and $N=0 s^{-1}$ (crosses). The continuous line is the neutral-stability line for centrifugal instability  computed by Snyder (1968) in an un-stratified Taylor-Couette experiment. In the centrifugally unstable case, one sees that the stratification enhances stability. In the centrifugally stable case, stratification favors instability.} \\label{fig:figure1} \\end{figure}  We note that in astrophysical context, the stability of plane rotating shear flow has often been tackled through the so-called shearing sheet transformation introduced by Goldreich and Lynden-Bell (1965). In this approximation, the disk structure is equivalent to a rotating plane Couette flow, with shear being $-3/2$ of the rotation. The advantage of this kind of analysis is the possibility to investigate stability, by using a special class of disturbances for which the stability analysis of this flow can be reduced to study of ordinary differential equations. These modes, first introduced by Lord Kelvin in 1887, have the shape $u(x,y,z,t,k,\\beta,k_z)=u(t)\\exp(i(k-S\\beta t)x+i\\beta y+ik_z z)$. Here $S$ is the Couette shear rate, with the flow velocity in the $y$-direction varying with $x$. Individual non-axisymmetric such modes have been found to exhibit algebraic transient growth under various conditions, such as compressibility (Dubrulle \\& Knobloch 1992) or vertical stable and unstable stratification (Knobloch 1984; Korycansky 1992). However, any amount of dissipation causes ultimate decay of these modes, due to shear winding of the azimuthal wave-numbers. From a theoretical point of view, it can be shown that these modes are associated with the non-normality of the operator governing the perturbation dynamics. They may be the basis of a noise amplification mechanism, resulting in significant angular momentum flux (Ioannou \\& Kakouris 2001). However, these modes do not form a complete basis (for example, they cannot describe perturbations with periodic boundary conditions in $x$), and their study alone is not sufficient to determine the stability properties of the flow, as recognized by Korycansky (1992). A complete stability analysis requires solving the ``initial condition\" problem, looking at the stability properties of an infinite superposition of sheared modes, moving with the flow. However, even in the case of non-stratified rotating shear flow, this problem is difficult to handle analytically and calls for a numerical solution (Cambon et al. 1994). Therefore, the shearing-sheet/sheared mode method provides only a partial answer to the stability problem.\\  In the present paper, we resort to other approaches, in which boundary conditions must be fixed a priori, but which allow for a complete treatment of the instability. Our approach is complementary to the usual sheared mode approximation. We use two classical, analytically tractable approaches to reexamine the instability of a stratified astrophysical disk: one based on the energy method, the second based on the WKB approximation. The energy method is valid for a wide class of boundary conditions, namely rigid (vanishing velocity at the domain boundary), stress free (vanishing velocity derivative at the domain boundary) or periodic boundary conditions in the shear direction. Therefore, it does not apply to individual sheared modes, which satisfy none of these requirements.  This method leads to a sufficient condition for stability valid in the inviscid limit, and for perturbations with a characteristic scale much smaller than the vertical scale height (Section 2.1). The influence of curvature on this instability is discussed in Section 2.2. In Section 3, we use a method based on the WKB approximation to derive an explicit solution of the stability problem and exhibit unstable modes in the parameter space covering the sufficient condition for stability. This shows that the sufficient condition for stability is probably necessary. This theoretical study is completed in section 4 by a numerical study, probing instability regimes and the influence of dissipative processes on the instability. This study enlarges the numerical study of Yavneh et al. (2001) towards conditions more typical of astrophysical disks, namely Keplerian velocity profile and finite Prandtl number. In Section 5, we discuss the importance of this mechanism for turbulence generation in disks, and the similarity between this instability and the instability generated by a vertical magnetic field. Our conclusions follow in Section 6, where the interplay between the present shear instability and its baroclinic counterpart is discussed. ",
        "conclusions": "In this paper, we have used analytical and numerical methods to study the instability of Keplerian-like flow in the presence of a stable vertical stratification,  for perturbations obeying rigid, stress-free or periodic boundary conditions in the radial coordinates that lie less than one scale height apart so as to justify an incompressibility approximation. This instability is non-axisymmetric. Its growth-rate is a fraction of the rotation rate at small stratification. This `strato-rotational' instability is purely hydrodynamical and does not require the presence of any magnetic field, whatever small. Therefore, it is especially relevant in the context of weakly ionized disks, such as the primitive solar nebula. It is of course also relevant in other Keplerian disks, since the critical Reynolds number to trigger this instability is of the order of $10^3$, which is less than the critical Reynolds number for both the magneto-rotational instability and the finite-amplitude hydrodynamic instability. It may also interplay with baroclinic instabilities, resulting in mixed instabilities similar to those observed in oceans. However the question of which instability occurs first has little relevance, partly because we do not know the initial conditions, but mainly because there is no reason why the instability which has the fastest linear growth will dominate in the fully nonlinear state. What one should rather ask is what kind of turbulence exists at the high Reynolds numbers which characterize astrophysical disks. A partial answer to this question may already be found from Taylor-Couette laboratory measurements (Dubrulle et al. 2004), which provide information about the hydrodynamical regime of rotating shear flow, as well as the influence of magnetic field or stratification on the transport properties. When applied to circumstellar disks, this provides parameter free-predictions about accretion rates and fluctuations which are in good agreement with observations (Hersant et al. 2003). For the influence of processes more specific to astrophysics (radiative transport, etc.) one will probably have to wait until the numerical simulations are able to reach Reynolds numbers of order $10^6$ or more."
    },
    "0410/astro-ph0410538_arXiv.txt": {
        "abstract": "{ A new tool to extract cross-identifications and radio continuum spectra from radio catalogues contained in the VIZIER database of the CDS is presented. The code can handle radio surveys at different frequencies with different resolutions. It has been applied to 22 survey catalogues at 11 different frequencies containing a total of 3.5 million sources, which resulted in over $700\\,000$ independent radio cross-identifications and $\\sim 67\\,000$ independent radio spectra with more than two frequency points. A validation of the code has been performed using independent radio cross-correlations from the literature. The mean error of the determined spectral index is $\\pm 0.3$.  The code produces an output of variable format that can easily be adapted to the purpose of the user. ",
        "introduction": "}  The total number of records in radio-source catalogues has increased dramatically in the last two decades. Three major increases are due to R.~Dixon's 'Master Source List' in the seventies (for the first version see Dixon 1970; for an error list see Andernach 1989), to the 87GB (Gregory \\& Condon 1991), GB6 (Gregory et al. 1996), and PMN (Wright et al. 1994, 1996; Griffith et al. 1994, 1995) surveys in 1991, and to the release of NVSS (Condon et al. 1998), FIRST (White et al. 1998) and WENSS (Rengelink et al. 1997) in 1996/1997 (see also Andernach 1999). In order to study the nature of these celestial objects detected at radio frequencies and to take full advantage of the huge amount of information contained in the catalogues, one has to know their spectral energy distribution (SED) over as large a frequency range as available. Since each catalogue (except Dixon's master source list) contains information at a single frequency, cross-identification between different catalogues is essential for the study of these sources.  Radio source cross-identifications in the centimeter to meter wavelength domain are particularly difficult to obtain, because the underlying radio surveys can have huge differences in sensitivity and/or spatial resolution. Since the resolution depends on the observed frequency and the telescope diameter, low frequency surveys made with a small single-dish telescope can have resolutions up to tens of arcmins, while high-frequency observations with a large single dish telescope or an interferometer can have resolutions of a few arcsecs to tens of arcsecs.  On the other hand, the cross-identification of radio sources at different frequencies is made easier by using the fact that, in the vast majority of sources, the SED has a power-law distribution.  The radiation mechanism  is either synchrotron emission from relativistic electrons gyrating in a magnetic field, or emission of hot thermal electrons. Synchrotron emission produces a power law spectrum with a possible cut-off or reversal of the spectral index at low frequencies due to self-absorption or comptonization. The spectrum of thermal electrons is flat, at least in the optically thin domain. Over the frequency range in which the majority of radio surveys were made, the spectra are thus well defined by a power law, i.e. as a straight line in the $\\log$(flux density)--$\\log$(frequency) plot commonly used in radio continuum astronomy.  The cross-identification procedure assigns radio sources at different frequencies to one physical object. In this way accurate radio positions can be determined for these objects using the high frequency observations, their radiation processes can be studied, and a search for specific objects (data mining) becomes possible.  The VIZIER database\\footnote{http://vizier.u-strasbg.fr/viz-bin/VizieR} (Ochsenbein et al. 2000) at the Centre de Donn\\'ees astronomique de Strasbourg (CDS) contains approximately 500 catalogues with radio data. Of these 500, about 70 catalogues are from systematic surveys (for a list of the major surveys see Andernach 1999). Using VIZIER, only a cone search (where a central position and a radius is used) is possible on all, or on a subset of, these catalogues. This procedure gives a list of all radio sources within the search region without any cross-identification.  Cross-identifications of radio sources within SIMBAD\\footnote{http://simbad.u-strasbg.fr/sim-fid.pl} (Wenger et al. 2000) are essentially made on bibliographic grounds. Sources are only merged when a newly published radio catalogue gives alternate names, which were already known to SIMBAD. Thus there is a clear need for establishing links between the radio catalogues (cross-identifications) and for including the sources into SIMBAD.  \\subsection{Radio surveys}  There are two main kinds of radio surveys: (i) single dish and (ii) interferometric. Since interferometers cannot detect extended structures larger than the size corresponding to the angular resolution  of the shortest baselines (ten/tens of arcmin for compact configurations), only structures of this size, or smaller, can be detected in sources of large extent. This makes the identification and unique spectral index determination complex for very extended sources, if they were observed with both single-dish and interferometer telescopes.  In addition, the observations can be divided into two subclasses: (i)systematic surveys and (ii) surveys of a given source sample (e.g. SN remnants, H{\\sc ii} regions, etc.). While the first category is more suitable for integration into VIZIER, the second one is easier to feed into SIMBAD. Out of the 500 catalogues listed in VIZIER when choosing ``radio'', about 220 contain independent observations. Of these, about 80 represent systematic surveys and about 140 surveys of source samples.  The characteristics of radio surveys that have to be taken into account when searching for cross-identifications are \\begin{itemize} \\item frequency, \\item sky coverage, \\item angular resolution (half power beam width, HPBW), \\item sensitivity to point-like and extended sources. \\end{itemize}  From these surveys source catalogues are established. The sources are extracted from 2D maps using 2D Gaussian fits, which have in principle 4 parameters: (i) the center position, (ii) the major axis, (iii) the minor axis, and (iv) the position angle. One distinguishes 'map' parameters and 'sky' parameters of sources. 'Map' parameters are the extent along the major and minor axis from the Gaussian fit, whereas 'sky' parameters are the deconvolved ones. If $\\Theta$ is the survey resolution (HPBW) and $d$ the source extent on the map, then the true source extent is $d_{\\rm true} = \\sqrt{d^{2}-\\Theta^{2}}$.  The catalogues can have different formats. For a given source one may find positions, position errors, a name, the peak and integrated flux densities and their errors, the major/minor axis, the position angle, and various flags (confusion, extended source, warnings, etc.). The use of unified content descriptors (UCDs\\footnote{see http://vizier.u-strasbg.fr/doc/UCD/inVizieR.htx}), which is a classification scheme in which all the astronomical parameters accessible in VIZIER are stored, will make the data access much easier in the future.   \\subsection{Cross-identification of radio sources}   There are three categories of cross-identification between radio sources: \\begin{itemize} \\item cross-identification based on the proximity between two sources, \\item cross-identification based on the radio spectrum, \\item cross-identification based on other physical characteristics (e.g. galaxies, SN remnants, quasars, AGNs etc.). \\end{itemize} The last two kinds of cross-identifications are called 'value-added' cross-identification. In the present paper we present a tool for cross-identification called SPECFIND which uses a proximity criterion and assumes that all sources have a power law radio spectrum.  This paper is structured in the following way: Sect.~\\ref{sec:sql} describes the uniformisation of catalogue tables, which is required because the catalogue table entries often differ. The code algorithm is discussed in Sect.~\\ref{sec:general}, followed by a discussion on the detailed code structure (Sect.~\\ref{sec:detailed}). The code performance is presented in Sect.~\\ref{sec:performance} and the results are shown in Sect.~\\ref{sec:results}. The code is validated by the comparison of our spectral indices with independent estimates from the literature (Sect.~\\ref{sec:validation}). The summary and conclusions are given in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "}  SPECFIND is a very efficient tool to identify radio spectra using radio catalogues of different formats. SPECFIND can handle radio surveys of very different resolutions and sensitivities. It has been applied to 22 survey catalogues at 11 different frequencies containing a total of 3.5 million sources, leading to more than $700\\,000$ independent radio cross-identifications and $\\sim 67\\,000$ independent radio spectra with more than two independent frequencies. The code was tested and its results validated by a comparison between the spectral indices found by SPECFIND and those determined by other authors. The determined spectral indices have an error of about $\\pm0.3$. Negative spectral indices have smaller errors, while the error of positive spectral indices can be larger, mainly because of the occurrence of a peak in the spectrum. The code is quite rapid (less than 3~hr running time on a standard PC for 3.5 million sources) and since it is written in C, it can be run on virtually all PCs with at least 512~MB RAM. It produces an output of variable format that can be adapted easily to the purpose of the user. A special output to enter the cross-identifications into SIMBAD has been developed. The code has been optimised to find a maximum number of spectra with a relatively small number of misidentifications. It represents thus a promising tool to extract radio spectra from a large sample of radio catalogues, like those stored in VIZIER. Although at present only 22 catalogue entries can be accessed by SPECFIND, this number will be increased in the future. The advantage of this procedure is that radio spectra and cross-identifications are established in advance, and can be stored in a 'master list' within VIZIER, allowing a fast and efficient access. We expect to make the cross-identifications soon available via the VIZIER database."
    },
    "0410/astro-ph0410012_arXiv.txt": {
        "abstract": "{The {\\em Chandra} X-ray Observatory is finding a large number of cavities in the X-ray emitting intra-cluster medium which often coincide with the lobes of the central radio galaxy. We propose high-resolution Sunyaev-Zel'dovich (SZ) observations in order to infer the still unknown dynamically dominating component of the radio plasma bubbles. This work calculates the thermal and relativistic SZ emission of different compositions of these plasma bubbles while simultaneously allowing for the cluster's kinetic SZ effect.  As examples, we present simulations of an {\\em Atacama Large Millimeter Array (ALMA)} observation and of a {\\em Green Bank Telescope (GBT)} observation of the cores of the Perseus cluster and Abell~2052.  We predict a 5$\\sigma$ detection of the southern radio bubble of Perseus in a few hours with the {\\it GBT} and {\\it ALMA} while assuming a relativistic electron population within the bubble.  In Abell~2052, a similar detection would require a few ten hours with either telescope, the longer exposures mainly being the result of the higher redshift and the lower central temperature of this cluster.  Future high-sensitivity multi-frequency SZ observations will be able to infer the energy spectrum of the dynamically dominating electron population in order to measure its temperature or spectral characteristics.  This knowledge can yield indirect indications for an underlying radio jet model. ",
        "introduction": "\\label{sec:intro} The {\\em Chandra} X-ray Observatory is detecting numerous X-ray cavities in clusters of galaxies confirming pioneering detections of the {\\em ROSAT} satellite. Prominent examples include the \\object{Perseus cluster} \\citep{1993MNRAS.264L..25B, 2000MNRAS.318L..65F}, the \\object{Cygnus-A cluster} \\citep{1994MNRAS.270..173C}, the \\object{Hydra-A cluster} \\citep{2000ApJ...534L.135M}, \\object{Abell 2597} \\citep{2001ApJ...562L.149M}, \\object{Abell 4059} \\citep{1998ApJ...496..728H, 2002ApJ...569L..79H}, \\object{Abell 2199} \\citep{fabian2001moriond}, \\object{Abell 2052} \\citep{2001ApJ...558L..15B}, the vicinity of \\object{M84} in the \\object{Virgo cluster} \\citep{2001ApJ...547L.107F}, the \\object{RBS797 cluster} \\citep{2001A&A...376L..27S}, and the \\object{MKW3s cluster} \\citep{2002ApJ...567L..37M}.  They are produced by the release of radio plasma from active galactic nuclei (AGN) which are typically hosted by a cD galaxy located at the cluster center and mainly reside within the cool cores of galaxy clusters.  While radio synchrotron emission provides evidence for the existence of cosmic ray electrons (CRes) and magnetic fields, the detailed composition of the plasma bubble governing its dynamics is still unknown.  Minimum energy or equipartition estimates of the nonthermal pressure in the radio bubbles give values which are typically a factor of ten smaller than the pressures required to inflate and maintain the bubbles as determined from the surrounding X-ray gas \\citep[e.g.,][]{2001ApJ...558L..15B}.  This indicates that the standard minimum energy or equipartition radio arguments are missing the main component of the pressure and energy content of the radio lobes.  Possibilities include magnetic fields, cosmic ray proton (CRp) or CRe power-law distributions, or very hot thermal gas.  Solving this enigma would yield further insight into physical processes within cool cores \\citep{2003MNRAS.343..719D} as well as provide hints about the composition of relativistic outflows of radio galaxies because plasma bubbles represent the relic fluid of jets \\citep[e.g.][]{1998MNRAS.293..288C, 1998tx19.confE.407H, 2000ApJ...534..109S}.  Additionally, some of the clusters exhibit cavities in the X-ray emitting intra-cluster medium (ICM) without detectable high frequency radio emission, for instance in \\object{Perseus}, \\object{Abell 2597}, \\object{Abell 4059}, and the \\object{MKW3s cluster}. This category of X-ray cavities is also believed to be filled with radio plasma, but during the buoyant rise of the light radio plasma bubble in the cluster's potential \\citep{gull1973,2000A&A...356..788C,2001ApJ...554..261C,2001MNRAS.325..676B} the resulting adiabatic expansion and synchrotron/inverse Compton losses dwindles the observable radio emitting electron population producing a so-called {\\em ghost cavity} or {\\em radio ghost} \\citep{1999dtrp.conf..275E}. Possible entrainment of the ICM into the plasma bubble and subsequent Coulomb heating by CRes generates further uncertainty of the composition of the ghost cavity.  If the radio bubbles contain a significant amount of very hot thermal gas, this might be detected by X-ray observations.  However, this is quite difficult \\citep[e.g.,][]{2003ApJ...585..227B} due to the projected foreground and background cluster emission, and the fact that the X-ray emissivity is proportional to the square of the density.  If most of the pressure in the radio bubbles were due to the very hot, low density thermal gas, it would have a very low X-ray emissivity. Observationally, there has been a claim by \\citet{2002ApJ...567L..37M} to have seen hot X-ray emitting gas within the ghost cavity of the MKW3s cluster.  Obviously, it might be more useful to observe the radio bubbles with a technique which was sensitive to thermal gas pressure, rather than density squared, as pressure is the quantity which is missing.  For this reason, in this paper we propose high resolution Sunyaev-Zel'dovich (SZ) radio observations of the radio bubbles and radio ghosts in clusters, as the thermal SZ effect directly measures the thermal electron pressure in the gas. The thermal SZ effect arises because photons of the cosmic microwave background (CMB) experience inverse Compton collisions with thermal electrons of the hot plasma inside clusters of galaxies and are spectrally redistributed \\citep[e.g.][]{1972SZorig, 1980ARA&A..18..537S, 1995ARA&A..33..541R}. The proposed measurement is able to infer directly the composition of radio plasma bubbles and radio ghosts while indirectly obtaining indications for a specific underlying jet model.  The article is organized as follows: After basic definitions concerning the thermal, kinetic, and relativistic SZ effect in Sect.~\\ref{sec:rSZ}, we introduce a toy model in Sect.~\\ref{sec:toymodel} describing projected maps of the SZ flux decrement with spherically symmetric radio plasma bubbles.  The models for the cool core regions of the Perseus cluster and Abell~2052 are described in Sect.~\\ref{sec:perseus}.  Simulating an {\\em Atacama Large Millimeter Array (ALMA)} and a {\\em Green Bank Telescope (GBT)} observation of both clusters in Sect.~\\ref{sec:synthetic}, we examine whether the plasma bubbles are detectable by the SZ flux decrement.  Five physically different scenarios for the plasma composition of the bubbles are investigated exemplarily using three characteristic SZ frequencies in Sects.~\\ref{sec:composition} and \\ref{sec:kinSZ}. Finally, we discuss observing strategies for {\\em ALMA} and {\\em GBT}.  Throughout the paper, we assume a $\\Lambda$CDM cosmology and the Hubble parameter at the present time of $H_{0} = 70\\, h_\\rmn{70}~{\\rm km~s^{-1}~Mpc^{-1}}$. ",
        "conclusions": "This work provides a theoretical framework for studying the SZ decrement of radio plasma bubbles within clusters of galaxies. X-ray observations are proportional to the square of the thermal electron density and probe the core region of the cluster. On the other hand, the SZ effect is proportional to the thermal electron pressure enabling the detection of plasma bubbles further outwards of the respective clusters.  Assuming spherically symmetric plasma bubbles, we simulate an {\\em ALMA} and a {\\em GBT} observation of the cool core regions of the Perseus cluster and Abell~2052. In this context, we exemplarily investigate physically different scenarios of the composition of the plasma bubbles: As long as the bubble is dynamically dominated by relativistic electrons, protons, or magnetic fields, there exists a realistic chance to detect plasma bubbles in the SZ flux decrement. Non-detection of radio bubbles with the SZ effect at the position of X-ray cavities hints towards a dynamically dominant hot thermal electron population within radio plasma bubbles.  Detection of a non-thermal pressure support should be possible within a few hours observation with ALMA or GBT in the case of Perseus and a few ten hours in the case of Abell~2052.  Pursuing high-sensitivity multi-frequency SZ observations, it will be challenging but not impossible to infer the detailed nature of the different possible populations of the bubble.  For realistic observations, the frequency dependence of the relativistic SZ signal is contaminated by the presence of the kinetic SZ signal.  It would be optimal to have a frequency channel centered on the crossover frequency of the thermal SZ effect in order to infer the energy spectrum of the dynamically dominating population and to measure the temperature or spectral characteristics of the electron population."
    },
    "0410/hep-th0410033_arXiv.txt": {
        "abstract": "We consider a single Minkowski brane sandwiched in between two copies of anti-de Sitter space. We allow the bulk Planck mass and cosmological constant to differ on either side of the brane. Linearised perturbations about this background reveal that gravity can be modified in the infra-red. At intermediate scales, the braneworld propagator mimics four-dimensional GR in that it has the correct momentum dependance. However it has the wrong tensor structure. Beyond a source dependant scale, we show that quadratic brane bending contributions become important, and conspire to correct the tensor structure of the propagator. We argue that even higher order terms can consistently be ignored up to very high energies, and suggest that there is no problem with strong coupling. We also consider scalar and vector perturbations in the bulk,  checking for  scalar ghosts. ",
        "introduction": "How much do we believe in Einstein's General Theory of Relativity? We certainly don't believe it at energies beyond the Planck scale, where a quantum theory is required. In fact, experimental tests of gravity have only ever been conducted at much lower scales. For example, Newton's law has been tested down to about $0.2$ mm, or equivalently up to energies of around $10^{-3}$ eV. Solar system tests of GR are only valid outside of the Schwarzschild radius of the Sun.  Ultra-violet modifications of gravity beyond $10^{-3}$ eV occur in all higher dimensional theories, such as String Theory. Recently, the idea that gravity is also modified in the infra-red has received much interest~\\cite{Dvali:DGP, Kogan:multi5, Papazoglou:thesis, Gregory:GRS,Arkani-Hamed:ghost1,Arkani-Hamed:ghost2, Padilla:bigravity}. This is entirely plausible. Again, Newton's law has only been tested up to distances of around  $10^{26}$ cm, and solar system tests of GR do not extend much beyond the size of the solar system itself.  Infra-red modifications of gravity have some attractive features. They can be used to tackle the cosmological constant problem~\\cite{Dvali:diluting, Barvinsky:cc} by weakening gravity at low energies. To understand this in a simple way, suppose we have Newton's constant $G$ at intermediate scales, and $G_\\textrm{IR} \\ll G$ in the infra-red. Even for a large cosmological constant $\\La$, the contribution to the Hubble rate is considerably suppressed at large values of the scale factor \\be H^2 \\sim G_\\textrm{IR}\\La \\ll G\\La \\ee However, that is not to say that these models prohibit late  time inflation. On the contrary, these models generically give rise to Friedmann equations that are non-linear in $H^2$. These permit self accelerating solutions that can be tuned to the current Hubble rate~\\cite{Deffayet:cosmo,Lue:accn,Dvali:power,Damour:accn, Holdom:accn,Sahni:bwde}.  Unfortnately, these models are not without their problems. The most common of these is the vDVZ discontinuity that appears in the graviton propagator at all scales~\\cite{vanDam:VDVZ, Zakharov:VDVZ}. This represents a deviation from GR that will give the wrong prediction for light bending around the Sun. In some cases, it has been argued that this discontinuity reflects a breakdown in the linearised analysis~\\cite{Vainshtein:nonVDVZ,Deffayet:Vainshtein}, as opposed to a genuine sickness in the theory. Some theories predict ghost states, and therefore violate unitarity~\\cite{Pilo:ghost}. A sickness as bad as this usually leads to the death of a theory. We should emphasise, however, that some models are indeed ghost-free~\\cite{Dvali:DGP, Padilla:bigravity}.  In~\\cite{Padilla:asymm} we considered {\\it asymmetric} branes. By asymmetric, we mean that the gravitational parameters of the theory can differ on either side of the brane.  There are at least two ways in which this asymmetry might arise. Firstly, suppose we have some sort of wine bottle shaped compactification down to 5 dimensions. In the effective theory, the Planck scale at the fat end  of the bottle will be less than that at the thin end. Secondly, in~\\cite{Padilla:nested, Padilla:instantons} we showed how to construct a domain wall living entirely on the brane. In some cases~\\cite{Padilla:instantons}, the Planck scale on the brane differed on either side of the domain wall. Asymmetric branes admit self accelerating solutions~\\cite{Padilla:asymm}. It is natural to ask, therefore, if asymmetric brane models exhibit infra-red modifications of gravity.  In this paper, we will consider a positive tension brane. To the left of the brane, the bulk is described by the Einstein-Hilbert action with Planck mass, $M_L$, and negative cosmological constant, $\\La_L=-6/l_L^2$. To the right, we have Planck mass, $M_R$, and cosmological constant, $\\La_R=-6/l_R^2$. In general, we will take $M_L \\gg M_R$ and $l_L \\ll l_R$.  We will be particularly interested in Minkowski branes. Because of the asymmetry, we are allowed three classes of solution. We have (i) a Randall-Sundrum like solution where the warp factor in the metric decays away from the brane on both sides~\\cite{Randall:RS2}, (ii) an ``inverse'' RS solution where the warp factor grows on both sides, and (iii) a mixed solution where the warp factor decays on one side and grows on the other.  If we place a small matter source on the brane, we can investigate linearised perturbations about the background solutions. For the RS solution and the mixed solution we see the introduction of a new length scale \\be r=\\frac{M_L^3 l_L}{M_R^3} \\ee At intermediate energies ($1/r \\ll p \\ll 1/l_L$), the braneworld propagator is proportional to $1/p^2$, and will reproduce Newton's law. At lower energies ($1/l_R \\ll p \\ll 1/r$), the momentum dependance changes, so that Newton's law is modified in the infra-red. However, as in the DGP model~\\cite{Dvali:DGP}, the propagator suffers from a vDVZ discontinuity at intermediate scales, and will predict the wrong results for light bending around the Sun.  It turns out that the vDVZ discontinuity can be removed by including brane bending terms to {\\it quadratic} order. For a source of mass $m$, we discover yet another scale \\be p_*=\\left(\\frac{M_L^3 l_L}{mr^2}\\right)^\\third. \\ee For $p_* \\ll p \\ll 1/l_L$, the quadratic terms dominate over the linear ones. The result is that the braneworld propagator agrees exactly with four-dimensional General Relativity. We might be worried about the break down of the linearised analysis at this scale. In the DGP model, this breakdown has been linked to a strong coupling problem~\\cite{Luty:strong,Rubakov:strong, Arkani-Hamed:massive,Porrati:massive}. However, we will argue that there is nothing to worry about. This is because even though quadratic terms become important when  $p \\sim p_*$, even higher order terms are only important at much higher energies. It seems that we can indeed consistently modify four-dimensional GR in the infra-red.  The rest of this paper is organised as follows. In section~\\ref{setup}, we describe our set-up in detail, before deriving background solutions in section~\\ref{background}. In section~\\ref{sec:linear} we derive the linearised equations of motion and solve them. We analyse the solutions and establish the existance of the long distance scale $r$, aswell as the vDVZ discontinuity. In section~\\ref{sec:radion}, we consider scalar and vector perturbations in the bulk. Vector perturbations only exist for the inverse RS solution. Scalar perturbations exist whenever we include the AdS boundary. We derive the effective action to check that these scalar fields have a well behaved kinetic term. In other words, they are {\\it not} ghosts. In section~\\ref{sec:quadratic}, we go to quadratic order in the brane bending. We discover the new scale $p_*$, and are able to correct the tensor structure of the propagator. In section~\\ref{sec:cosmo}, we briefly show how the mixed case admits self accelerating solutions. Finally, section~\\ref{sec:discuss} contains some discussion. In particular, we argue that the quadratic analysis is valid up to very high scales, and that there is no strong coupling problem. ",
        "conclusions": "\\label{sec:discuss} We have considered in detail perturbations about a flat brane embedded in between two copies of anti-de Sitter space. We have abandoned $\\mathbb{Z}_2$ symmetry across the brane, giving rise to three distinct cases: (i) the RS case where the warp factor decays away from the brane on both sides, (ii) the inverse RS case where the warp factor grows on both sides and (iii) the mixed case where the warp factor decays to the left of the  brane, and grows to the right.  Let us now summarize our main results. The inverse RS case is a dead end: four dimensional Einstein gravity can {\\it never} be achieved. In contrast, take the RS case and the mixed case when $M_L \\gg M_R$ and $M_L^3 l_L \\ll M_R^3 l_R$. Consider a source of mass, $m \\in ( M_L^3 l_L^4/r^2, M_L^3 l_L r)$, where $r=M_L^3 l_L/M_R^3$. The metric on the brane is given by \\be \\label{summary} \\tilde h_{\\mu\\nu}^\\textrm{br}={\\begin{cases}\\frac{1}{m_{pl}^2p^2} \\left[ \\tilde T_{\\mu\\nu}-\\half \\tilde T \\eta_{\\mu\\nu} \\right] & \\textrm{when $p_* \\ll p \\ll p_\\textrm{uv}$}\\\\ \\frac{1}{m_{pl}^2p^2} \\left[ \\tilde T_{\\mu\\nu}-\\third \\tilde T \\eta_{\\mu\\nu} \\right] & \\textrm{when $1/r \\ll p \\ll p_*$}\\\\ \\frac{1}{M_R^3p} ~\\left[ \\tilde T_{\\mu\\nu}-\\third \\tilde T \\eta_{\\mu\\nu} \\right] & \\textrm{when $1/l_R \\ll p \\ll 1/r$} \\end{cases}} \\ee where $m_{pl}^2=\\half M^3_L l_L$, $p_*=(M^3_Ll_L/mr^2)^{1/3}$, and $p_\\textrm{uv}=\\textrm{min} \\{ 1/l_L, M^3_L l_L/m \\}$. Gravity on the brane mimics four-dimensional GR at intermediate scales, but is modified in the infra-red. This modification begins with a change in the tensor structure of the propagator. At even lower energies the momentum dependance also changes. Crucially, these models are free from ghosts.  We should also note that a cosmological brane in the mixed case will approach a de Sitter phase at late times, even when there is no additional matter. This does not happen for the RS case.  To illustrate our results more clearly, let us put some numbers in. Consider the mixed case with $M_L \\sim 1/l_L \\sim m_{pl} \\sim 10^{19}$ GeV, $M_R \\sim  10 - 100$ MeV, and $1/l_R \\ll 10^{-34}$ eV. Firstly, note that the current Hubble rate, $H_0 \\approx 2/r \\sim 10^{-34}$ eV, as desired. Now consider our source to be the Sun, with mass $m \\sim 10^{66}$ eV. Four-dimensional GR is reproduced at distances between $r_\\textrm{Schw}$ and $10^{21}$ cm. Here, the Schwarzschild radius of the Sun is given by $r_\\textrm{Schw} \\sim 1$ km. The large distance limit,  $10^{21}$ cm, extends well beyond the size of the solar system, so we have no conflict with experiment.   Of course,  the results (\\ref{summary}) were derived by going to quadratic order in perturbation theory. When $p_* \\ll p \\ll p_\\textrm{uv}$, we found that the quadratic brane bending term dominated over the linear terms. We might be worried that this means we have lost control of perturbation theory. However, we will now argue that we {\\it do} have control of perturbation theory, because terms beyond quadratic order are always small, right up to $p_\\textrm{uv}$.  Suppose, we have two GN coordinate systems, $x^a$ and $\\hat x^a$, where \\be \\hat x^\\mu=x^\\mu+\\xi^\\mu(x, z), \\qquad \\hat z=z+\\xi^z(x, z) \\ee In the ``unhatted'' coordinates, the brane is positioned at $z=0$. For simplicity we will demand that $\\xi^{\\mu}(x, 0)=0$. This just corresponds to a choice of gauge on the brane. In the ``hatted'' coordinates the brane is positioned at $\\hat z=\\xi^z(x, 0)=f(x)$.    Non-perturbatively, the metrics in each coordinate system are related as follows \\be g_{ab}(x, z)=\\hat g_{cd}(\\hat x, \\hat z) \\frac{ \\del \\hat x^c}{\\del  x^{a}} \\frac{ \\del \\hat x^d}{\\del  x^{b}} \\ee We shall now evaluate the $(\\mu \\nu)$ component at $z=0$ to derive the metric on the brane \\be \\label{nonpertbrane} g_{\\mu\\nu}(x, 0)=\\hat g_{\\mu\\nu}(x, f)+\\del_\\mu f\\del_\\nu f \\ee Suppose we write $\\hat g_{\\mu\\nu}=\\bar g_{\\mu\\nu}+\\chi_{\\mu\\nu}$, and Taylor expand the right-hand side of (\\ref{nonpertbrane}). We get $g_{\\mu\\nu}(x, 0)= g_{\\mu\\nu}^\\textrm{quad}(x)+\\delta g_{\\mu\\nu}(x)$, where \\be g_{\\mu\\nu}^\\textrm{quad}(x)=\\bar g_{\\mu\\nu}(x, 0)+f\\bar g_{\\mu\\nu}'(x, 0)+\\half f^2  \\bar g_{\\mu\\nu} ''(x, 0) +\\chi_{\\mu\\nu}(x, 0)+ f\\chi_{\\mu\\nu}'(x, 0) +\\del_\\mu f\\del_\\nu f \\ee and \\be \\delta g_{\\mu\\nu}(x) = \\sum_{n=3}^{\\infty} \\frac{f^n}{n!}\\bar g_{\\mu\\nu}^{(n)}(x, 0)+\\sum_{n=2}^{\\infty}\\frac{f^n}{n!}\\chi_{\\mu\\nu}^{(n)}(x, 0) \\ee Here prime denotes partial differentiation with respect to $z$. $ g_{\\mu\\nu}^\\textrm{quad}(x)$ includes all those terms that were considered in the analysis to quadratic order. $\\delta g_{\\mu\\nu}(x)$ contains all those terms that were assumed to be of higher order, and were consequently ignored. We shall now ask whether or not it was safe to ignore them.  Roughly speaking, we can associate each $z$ derivative with either a factor of $p$ or a factor of $1/l$. \tIt is obvious that as long as $pf \\ll 1$ and $f/l  \\ll 1$, then $\\delta g_{\\mu\\nu} \\ll g_{\\mu\\nu}^\\textrm{quad}$. In the quadratic analysis it is clear that $pf$ and $f/l$ are indeed small whenever $p \\ll p_\\textrm{uv}$. We conclude that it was entirely consistent to ignore all terms beyond quadratic order.  In analogy with the DGP model, we shall assume that a break down in perturbation theory gives rise to strong coupling. To derive the strong coupling scale we consider a source of mass $m \\sim m_{pl}$, and ask at what scale do we lose perturbative control~\\cite{Porrati:massive}. If this scale were $p_*$ we would be in trouble. We would expect strong coupling below distances of around $1000$ km. However, we have just shown that we don't expect to lose control until at least $p_\\textrm{uv}$, which typically corresponds to the Planck scale! These results suggest that there is no strong coupling problem here, although a more thorough analysis is clearly required.  Finally, note that the mixed case in particular appears to have much in common with the DGP model~\\cite{Dvali:DGP}. The expressions for $r$ and $p_*$ seem to have obvious DGP analogues. The cosmological behaviour is also very similar (see equation (\\ref{appFRW})). One can't help thinking that this is more than mere coincidence. With this in mind, consider the brane to be the common boundary to the decaying bulk and  to the growing bulk. If we were to calculate the boundary stress-energy tensor on the decaying side, we would expect there to be divergences. These must be cancelled by boundary counter terms that will probably take the form of localised curvature on the brane. If this is indeed the case, it would represent a new mechanism for obtaining induced curvature on a DGP brane. \\vskip .5in \\centerline{\\bf Acknowledgements} I would like to thank Christos Charmousis and John March-Russell for useful discussions. AP was funded by PPARC. \\medskip"
    },
    "0410/astro-ph0410297_arXiv.txt": {
        "abstract": "Over the last few years, high-quality X-ray imaging and spectroscopic data from {\\it Chandra\\/} and XMM-Newton have added greatly to the understanding of the physics of radio jets.  Here we describe the current state of knowledge with an emphasis on the underlying physics used to interpret multiwavelength data in terms of physical parameters. ",
        "introduction": "\\label{sec:introduction}  Jets in active galaxies emit over a wide range of energies from the radio to the $\\gamma$-ray.  Synchrotron radiation and inverse-Compton scattering are the two main radiation processes, with their relative importance depending on observing frequency and location within the jet.  The thermally-emitting medium into which the jets propagate plays a major role in the properties of the flow and the appearance of the jets.  In this chapter, we concentrate primarily on the interpretation of observations of resolved jet emission.  Much new information about jets has been gained over the past few years. The high spatial resolution of {\\it Chandra\\/} has been key to the study of X-ray jets, and the large throughput of XMM-Newton has assisted studies of the X-ray-emitting environments.  While historically the search for optical jet emission has been carried out using ground-based telescopes, a major problem has often been one of low contrast with light from the host galaxy.  The sharp focus of the Hubble Space Telescope (HST) helps to overcome this difficulty, and the HST is playing a particularly important role in optical polarization studies.  The recent opening of the Very Large Array (VLA) data archive has greatly assisted multiwavelength studies of radio jets.  It, together with the Australian Telescope Compact Array (ATCA) in the southern hemisphere, allows multifrequency radio mapping at angular resolutions well matched to {\\it Chandra\\/} and XMM-Newton. The combination of data from the radio to the X-ray has been a key element in advancing our understanding of jets.  Active-galaxy jets provide an exciting laboratory where the interplay of relativistic effects, plasma physics, and radiation mechanisms can be seen.  Our purpose in this chapter is to provide an introduction to the field and some pointers to recent research. ",
        "conclusions": ""
    },
    "0410/astro-ph0410068_arXiv.txt": {
        "abstract": "Changing the form of the star cluster mass function (CMF) can effectively change the upper end of the stellar initial mass function.  The yields of from supernovae are very sensitive to the mass of the progenitor star.  We show that by changing the parameters of the CMF, it is possible to change the yields of oxygen and magnesium by a factor of $\\sim 1.5$ and of metals in general by a factor of $\\sim 1.8$. ",
        "introduction": "Virtually all stars form in clusters (Clarke, Bonnell \\& Hillenbrand 2000; Lada \\&  Lada 2003).  The type-II supernovae (SNII) of the most massive  stars in these clusters are the main source of many heavy  elements, in particular oxygen and $\\alpha$-elements.  The yields of  heavy elements from these SNII depend upon the mass of the  progenitor (e.g. Tsujimoto et al. 1995; Woosley \\& Weaver 1995;  Hoffman et al. 1999; Limongi, Straniero \\& Chieffi 2000; Rauscher et al. 2002; Heger  et al. 2002).  Thus the yields of heavy elements will depend upon the initial mass function (IMF) of the most massive stars in clusters.  Of particular interest in the study of SNII products are oxygen and magnesium that are both thought to be pure SNII products and which are relatively easy to observe.  In large elliptical galaxies N/Fe, Na/Fe and Mg/Fe seem to be about  a factor of 2 supersolar (Henry \\& Worthey 1999), as do Mg/Fe and  O/Fe in stars of the Galactic halo and Thick disk (Prochaska et al. 2000). In the Galactic bulge,  anomalies in ratios such as Mg/Ca suggest changes in the initial mass  function (McWilliam \\& Rich 2004), which could also play a role in  the $\\alpha$/Fe ratios, although their behaviour is  usually attributed to differences in star formation history. In this investigation we are mainly concerned with the overall oxygen and magnesium yields, which can be deduced from the peak in the abundance distribution function of stars (Pagel 1997); this peak for Mg is evidently at a somewhat greater abundance in the Galactic Bulge and the central regions of large elliptical galaxies than the approximately solar value that it has in the solar neighbourhood.  In a recent paper Kroupa \\& Weidner (2003) showed that the (effective) upper IMF depends upon the cluster mass function (CMF).  Low-mass clusters are unlikely to contain very massive stars as they do not have the gas reservoirs from which to form them.  A $10^2 M_\\odot$ cluster cannot contain a star $> 100 M_\\odot$ and is unlikely to contain stars $>30 - 40 M_\\odot$.  A $10^5 M_\\odot$ cluster, on the other hand, is very likely to contain some stars  $> 100 M_\\odot$. This idea is supported by evidence for a truncated IMF in small clusters for stellar masses $> 10 M_\\odot$ (Thilker et al. 2002).  In the Solar neighbourhood the CMF is a power-law with slope $\\sim -2$ between $\\sim 10^2$ and $10^4 M_\\odot$ (Lada \\& Lada 2003).  There is evidence that the CMF varies with galaxy type (e.g. Kennicutt, Edgar \\& Hodge 1989; Thilker et al. 2002; Alonso-Herrero et al. 2002; Youngblood \\& Hunter 1999).  In this paper we discuss the effect on the yields of oxygen and magnesium from SNII due to changing the CMF (and so effectively changing the upper-end IMF).  In Section 2 we describe our method, in Section 3 we present the results, and in Section 4 we discuss their implications for heavy element yields in galaxies of different types and masses. ",
        "conclusions": "We have shown that the yields of O and Mg may vary by factors of $\\sim 1.5$, and of metals by $\\sim 1.8$, depending on the form of the cluster mass function.  This is due to the inability of low-mass clusters to  form very massive stars (cf. Kroupa \\& Weidner 2003) thus changing the fractions of contributors of different masses to enrichment.  In the Solar neighbourhood the CMF is a power-law with slope $\\sim -2$ between $\\sim 10^2$ and $10^4 M_\\odot$ (Lada \\& Lada 2003).  There is evidence that the HII luminosity function (and so presumably the CMF) varies with galaxy type (e.g. Kennicutt et al. 1989; Youngblood \\& Turner 1999) with early-type spirals having fewer and less bright HII regions than late-type galaxies. This suggests a steeper CMF and/or a lower-mass cut-off (although this conclusion is debated by Thilker et al. 2002; their sample contains mainly late-type galaxies, however). The upper limit of the HII luminosity function also varies with galaxy type, active star forming galaxies containing significantly more giant HII regions (e.g. Alonso-Herrero et al. 2002).  It might be expected that dwarf galaxies have a CMF that has a low upper-mass cut-off.  This is simply because dwarf galaxies have a far smaller reservoir of gas from which to form clusters, and the most massive clusters that they form would be expected to be fairly low. Thus dwarf galaxies would be expected to have sub-Solar yields of heavy elements.  The Solar Neighbourhood has a CMF that is close to ($2,50,10^4$) as shown by Lada \\& Lada (2003). This CMF produces close to the observed Solar abundances of O and metals, but under-produces Mg for these models.  Massive elliptical galaxies are thought to form from major merger events (Sanders \\& Mirabel 1996).  In such events, the star formation rate is thought to increase dramatically and super star clusters become a dominant mode of star formation.  Thus, a CMF with high upper and lower-mass cut-offs may be a reasonable model for these galaxies.  In such cases the yields could be a factor of $>1.5$ higher than in dwarf galaxies, and $>1.1$ higher than in normal spiral galaxies due solely to differences in the cluster mass functions."
    },
    "0410/astro-ph0410632_arXiv.txt": {
        "abstract": "We report on the first SCUBA detection of a Type 2 QSO at z=3.660 in the Chandra Deep Field South. This source is X-ray absorbed, shows only narrow emission lines in the optical spectrum and is detected in the sub-mm: it is the ideal candidate in an evolution scheme for AGN (e.g. Fabian (1999); Page et al. (2004)) of an early phase corresponding to the main growth of the host galaxy and formation of the central black hole. The overall photometry (from the radio to the X-ray energy band) of this source is well reproduced by the spectral energy distribution (SED) of NGC 6240, while it is incompatible with the spectrum of a Type 1 QSO (3C273) or a starburst galaxy (Arp 220). Its sub-mm (850 $\\mu$m) to X-ray (2 keV) spectral slope ($\\alpha_{\\rm SX}$) is close to the predicted value for a Compton-thick AGN in which only $1\\%$ of the nuclear emission emerges through scattering. Using the observed flux at 850 $\\mu$m we have derived a SFR=550--680 M$_{\\odot}$/yr and an estimate of the dust mass, M$_{\\rm dust}=4.2 \\times 10^8$ M$_{\\odot}$. ",
        "introduction": "The presence of massive black holes (MBHs) at the centers of most galaxies appears by now firmly established. A number of correlations have been observed between the masses of MBHs and the properties of the galactic bulges hosting them (e.g. \\cite{kormendy00}, \\cite{merritt01}). Since these correlations extend well beyond the direct dynamical influence of the MBH, it seems likely that a close link exists between the growth of the MBH and the formation of their host galactic bulges.  If these two events are co-eval then we would expect that powerful black holes are hosted in the centers of the most massive galaxies. SCUBA galaxies show a strong redshift clustering thus indicating that they are hosted by very high mass halos (\\cite{blain03}). This is confirmed by CO molecular gas emission line widths and possible rotation curves (\\cite{frayer98}; \\cite{genzel03}; \\cite{neri03}). The combination of deep X-ray observations and spectroscopic follow-up campaigns of SCUBA sources (\\cite{alexander04}) led to the discovery that at least $38 \\%$ of the SCUBA sources host an AGN, although in the majority of the cases the contribution of the nucleus is not bolometrically significant (i.e., $< 20 \\%$).\\\\ A possible scenario would predict that the black hole grows within a SCUBA galaxy until the obscuring gas is blown away revealing an AGN (e.g. \\cite{fabian99}). A similar evolution argument has been proposed for ULIRGs (\\cite{sanders96}), which can be considered the closest local analogs to SCUBA galaxies. We have started to investigate this scenario using a unique sample of heavily X-ray absorbed (log(N$_{\\rm H}) > 22$ cm$^{-2}$) and very luminous (log(L$_{\\rm X}) > 44$ erg s$^{-1}$) AGN, the so-called Type 2 QSOs, located in the Chandra Deep Field-South (CDF-S) and selected from \\cite{szokoly04}. In this letter we report on the first sub-mm detection of a high redshift (z=3.660) Type 2 QSO and discuss the implications for its nature. We will discuss in a future paper the properties of the full sample of Type 2 QSOs once our observational programme is completed. \\\\ We assume a cosmology with parameters $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$ and H$_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have a firm ($4 \\sigma$) SCUBA detection of a Type 2 QSO at redshift $z=3.660$ (CDFS-263). Its SED is consistent, from radio to X-ray wavelengths, with the shape of NGC 6240 while it can not be reproduced with the SED of 3C273 or Arp 220. Futhermore, we derive an $\\alpha_{\\rm SX}=1.29$. We conclude that CDFS-263 can be powered by a Compton thick AGN (we found N$_{\\rm H}\\sim 1-20 \\times 10^{23}$ cm$^{-2}$).  CDFS-263 is a strong candidate for an AGN in the initial phase (before the ``X-ray absorbed phase'') described by \\cite{page04} that corresponds to the main growth period of the host galaxy spheroid. It is a luminous X-ray source (L$_{\\rm X}>10^{44}$ erg s$^{-1}$) with high X-ray absorption (N$_{\\rm H} > 10^{23}$ cm$^{-2}$) and detected in the sub-mm, S$_{850}=4.8 \\pm 1.1$ mJy."
    },
    "0410/astro-ph0410318_arXiv.txt": {
        "abstract": "{High-resolution UCLES/AAT spectra of four B-type supergiants in the SMC  South East Wing have been analysed using non-LTE model atmosphere techniques to determine their atmospheric parameters and chemical compositions. The principle aim of this analysis was to determine whether the very low metal abundances ($-$1.1 dex compared with Galactic value) previously found in the Magellanic Inter Cloud region (ICR) were also present in SMC Wing. The chemical compositions of the four targets are similar to those found in other SMC objects and appear to be incompatible with those deduced previously for the ICR. Given  the close proximity of the Wing to the ICR, this is difficult to understand and some possible explanations are briefly discussed. ",
        "introduction": "Between the Large (LMC) and Small (SMC) Magellanic Clouds, two satellite galaxies of our own Milky Way, lies a  column of gas  known as the Magellanic Bridge. This Bridge greatly exceeds the galaxies' tidal limit, and is widely considered to be a remnant of the gravitational interactions between the Magellanic Clouds and the Milky Way. Numerical simulations by, for example, Gardiner \\& Noguchi (\\cite{Gar96}) and  Sawa, Fujimoto \\& Kumai (\\cite{Saw99}) imply that the tidal effects from a close encounter of the two Magellanic Clouds about $\\sim$200 Myr ago produced a tidal bridge and tail structure and subsequently triggered an episode of star formation.  This tail/bridge structure  is believed to constitute the inter-cloud region (ICR)  and would be seen as overlapping in the sky, with  the tail material modeled to be more distant than the bridge and with a  larger radial velocity (Gardiner \\& Noguchi \\cite{Gar96}). The SMC wing is located to the east of the SMC and enveloped by neutral hydrogen gas (see Fig.~\\ref{smcw_fig2}).  In optical images it appears as a large cloud of faint stars extending 6.5\\degr\\/ eastward from the SMC toward the LMC, and contains young stellar component (e.g.\\ Westerlund \\& Glaspey \\cite{Wes71}; Kunkel \\cite{Kun80}) that indicates recent star formation history. According to the predictions of the models (Kunkel et al.\\,\\cite{Kun94}), this Wing constitutes the densest regions of the bridge/tail material along the line-of-sight.  The tail/bridge material, if it is indeed the product of the past encounter(s) of the Clouds, would be expected to have the metallicity of its progenitors. However, a recent study by Rolleston et al.\\ (\\cite{Rol99}) of 3 B-type dwarfs in the ICR reported that they had a metallicity lower by $-$1.1 dex than that of our Galaxy, and this is also significantly lower than  the values found for the SMC and LMC (e.g.\\, Korn et al.\\ \\cite{Korn}; Kurt et al.\\ \\cite{Kur99}; Rolleston et al.\\ \\cite{Rol93, Rol96, Rol03};  Venn \\cite{Ven99}). A possible explanation is that the inter-cloud  material could have been stripped off from either of the galaxies at a  much earlier encounter (about $\\sim$8 Gyr ago; Kobulnicky \\& Skillman \\cite{Kob97}; Da Costa \\& Hatzidimitriou \\cite{DaC98}), when less nucleo-synthetic processing of interstellar medium (ISM) of the SMC had occurred. However, this would be inconsistent with the numerical simulations  discussed above that imply  this ICR was formed relatively recently and also with the presence of early-type stars.  Here we attempt to confirm and extend the results of Rolleston et al.\\ by investigating the chemical composition of the SMC wing. Using intermediate dispersion spectroscopy, we have identified suitable targets and here present high-resolution spectroscopy of four of these objects. By investigating the metallicity of these stars, we can address questions concerning the chemical homogeneity, or inhomogeneity, of the two distinct components, as well as possibly the tidal origin of the SMC wing material.  \\begin{table*} \\caption{Observational summary of the four program stars}\\label{tab:obs} \\begin{tabular}{lccrrrrrcc} \\hline\\hline\\noalign{\\smallskip} Star           &       RA,  Dec.         &  V   &  AAT/   & $T_{\\mathrm exp}$& $v_{\\mathrm lsr}$ & S/N\\\\ &       (2000.0)          & [mag]& UCLES   & [sec]            & [\\kms] & \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} ICR\\,02134-7431 &  02 13 26, $-$74 31 20  & 12.3 & 1999-10-25 &  7\\,500 & +197&  80 \\\\ SK\\,194         &  01 45 03, $-$74 31 32  & 11.7 & 2000-10-14 &  6\\,600 & +214& 140 \\\\ SK\\,202         &  01 53 03, $-$73 55 30  & 12.3 & 2000-10-14 & 25\\,200 & +180& 125 \\\\ SK\\,213         &  02 09 30, $-$74 26 00  & 13.7 & 2000-10-16 &  7\\,200 & +203&  60 \\\\ \\noalign{\\smallskip}\\hline \\end{tabular}\\\\ \\end{table*}  \\begin{figure} \\hspace*{-4mm}\\psfig{file=smcw_fig1.eps,width=9cm} \\caption{Normalised spectra of the 4 targets covering the region 4465 to 4485~\\AA, containing a \\ion{He}{i} and a \\ion{Mg}{ii} line. The spectra were rebinned to a pixel size of 0.1~\\AA. In SK\\,213 the \\ion{Mg}{ii} line has an equivalent width of approximately 70~m\\AA.} \\label{smcw_fig1} \\end{figure}  \\begin{figure} \\hspace*{-4mm}\\psfig{file=smcw_fig2.ps,angle=90,width=9cm} \\caption{Schematic diagram of the LMC, SMC and the Inter-Cloud Region. Solid lines define the stellar concentrations within the Magellanic Clouds, while the dashed lines show the H~{\\sc i} envelopes of the Magellanic system. The positions of our targets within the SMC wing are marked by solid black triangles. Other targets have been discussed by Rolleston et al.\\ (\\cite{Rol99}, circles;  \\cite{Rol03},  diamond). } \\label{smcw_fig2} \\end{figure} ",
        "conclusions": "The absolute abundances of metallic lines are summarised in Table~\\ref{tab:abs} on a logarithmic scale with [H] $\\equiv$ 12.0 dex. An observational uncertainty of $\\pm$10~m\\AA\\/ in the equivalent width estimate of an individual line would change an abundance estimate by typically $\\pm$0.15 dex. Where several lines due to an ionic species are observed the corresponding error in the mean abundance would be expected to be smaller. The abundance estimates may also be affected by errors in the adopted atmospheric parameters. The results, presented in Table~\\ref{tab:abs} for microturbulences of \\vt\\/ = 10 and 20 \\kms, lead to changes in the abundance estimates of typically $\\pm$0.1 dex, with larger uncertainties for in a few cases. For the effective temperature and gravity, numerical tests adopting error estimates of $\\pm$1\\,000~K and $\\pm$ 0.2\\,dex respectively, lead to changes in the abundance estimates of typically 0.1 to 0.3 dex.  Finally there will be uncertainties due to the intrinsic assumptions made in the model atmosphere calculations and in particular the neglect of the stellar wind. However, Trundle et al.\\  (\\cite{Tru04}) have used the non-LTE unified model atmosphere code {\\sc fastwind} first introduced by Santolaya-Rey, Puls \\& Herrero (\\cite{San97}) to analyse the spectra of SMC supergiants. For two targets, these results have been compared with those deduced using {\\sc tlusty} (see Dufton et al.\\  \\cite{Duf04}) and the comparison indicates that the neglect of the wind may not be a serious source of error.  Hence we conclude that an error estimate of $\\pm$0.3 dex might be appropriate to our abundance estimates.  The main purpose of this analysis is to investigate whether the very low metal abundances for the ICR found by Rolleston et al.\\ (\\cite{Rol99}) are also present in our SMC wing targets. To provide a comparison for our results, we have utilised three other recent studies of hot stars in the SMC. Trundle et al.\\  (\\cite{Tru04}, TLPD) studied 7 supergiants using the non-LTE unified model atmosphere code {\\sc fastwind}. Dufton et al.\\ (\\cite{Duf04}) have undertaken analysis of 9 supergiants using the same non-LTE grids as adopted here. As discussed above,  the two analyses yield similar results for two targets in common indicating that both TLPD and Dufton et al.\\ should form suitable comparisons. Lennon et al.\\ (\\cite{Len03}) corrected the LTE analysis of AV\\,304, a sharp lined B0V star, of Rolleston et al.\\ (\\cite{Rol03}) in order to allow for non-LTE effects. Although these corrections can  only be approximate, they would appear to provide estimates that agree well with a full non-LTE analysis currently being undertaken by the authors. The results of all these analyses are summarized in Table~\\ref{tab:abs}, the first two being for stars at a similar evolutionary to our targets, whilst the third should reflect the current chemical composition of the SMC interstellar medium. The three analyses yield different N (and to a lesser extent O) abundances estimates for reasons that are discussed below. However the excellent agreement for the O, Mg and Si abundances provide indirect support for their use as comparators for our SMC wing targets. Below we briefly discuss the estimates obtained for each element.  {\\bf Carbon:} The carbon abundances in the Wing targets appear to be lower for ICR\\,02134-7431 and SK\\,202 than for other SMC targets. The value for SK\\,213 appear relatively normal, whilst no \\ion{C}{ii} features were measurable in SK\\,194. However these results must be treated with caution for the estimates have used the doublet at 4267~\\AA. As discussed by, for example, Sigut (\\cite{Sigut}), this element suffers from significant non-LTE effects which are difficult to model. Additionally, Lennon et al.\\ (\\cite{Len03}) found that, even in a non-LTE analysis, this doublet appears to yield lower abundance estimates (by 0.3--0.4 dex) than is found from other \\ion{C}{ii} lines. Indeed the carbon abundance for AV\\,304 presented in Table~\\ref{tab:abs} has been corrected for this effect and this may explain the difference between the C abundance estimates in Dufton et al.\\ (\\cite{Duf04}) and for AV\\,304. Hence we do not believe that there is any significant evidence for the carbon abundance of our Wing targets being lower than that for the SMC in general.  {\\bf Nitrogen:} Nitrogen abundances were found for all four targets and range from 7.3 to 8.0 dex. TLPD also found a significant variation in their abundance estimates that they ascribed to different degrees of contamination by nucleo-synthetically processed material. This is also consistent with the far larger nitrogen abundances found in the supergiant samples than in AV\\,304 or in analysis of SMC \\ion{H}{ii} regions (see, for example, Kurt et al.\\ \\cite{Kur99}). Hence the N abundance estimates for our targets appear to be compatible with those for other evolved SMC supergiants.  {\\bf Oxygen:} The \\ion{O}{ii} spectrum is only observed in our hottest target with the abundance estimates being sensitive to the adopted microturbulence. However they are consistent with those found in other SMC B-type stars.  {\\bf Magnesium:} All the abundance estimates are based on the \\ion{Mg}{ii} doublet at 4481~\\AA. However they show little variation and are in excellent agreement with the other analyses.  {\\bf Silicon:} These abundance estimates are based on both \\ion{Si}{ii} and \\ion{Si}{iii} features. However, as the silicon ionization equilibrium has been used to constrain the effective temperature, the results have been combined into a single abundance estimate for each star. For our hottest target, SK\\,213, the results are affected by the choice of microturbulence but in general the estimates are compatible with other analyses of SMC objects.  The principle conclusion from the model atmosphere analysis is that  the SMC Wing targets appear to have similar chemical compositions to those of other objects in the SMC. In particular there is no evidence for the low metallicity found in the ICR targets by Rolleston et al.\\ (\\cite{Rol99}). From Fig.~\\ref{smcw_fig2}, it can be seen that all four targets are close to the inter-cloud region and it is arguable that two targets (ICR\\,02134-7431 and SK\\,213) are indeed part of the ICR or at the boundary between the Wing and the ICR. Given this close proximity, the difference in the chemical composition found by Rolleston et al.\\ and in the present study is surprising  especially as the ICR is thought to have formed from material from the SMC and LMC. The low metallicity found by Rolleston et al.\\ implies that the most likely source of the material is the SMC and then one is forced to the  conclusion that the ICR originated from a region close to (or indeed part of) the Wing but with a different chemical composition to that found for the parts of the Wing that we have sampled.  An alternative explanation is that the results in Rolleston et al.\\ are unreliable and that the ICR has a metallicity consistent with the SMC (or LMC). The low metallicity of the ICR was based on an LTE analysis of three targets with effective temperature in the range 20\\,000 to 24\\,000~K and logarithmic gravities of 3.6 to 3.8 dex; this would correspond to a spectral classification of early-B-type giants/dwarfs. The analysis was undertaken differentially with respect to two galactic stars with very similar atmospheric parameters, whilst the hottest ICR target (DI\\,1239) was also analysed relative to the SMC main sequence star AV\\,304, which has an effective temperature approximately 2\\,500~K higher. The results for all three stars were consistent and indicated that the metallicity of the ICR was between 1.0 and 1.2 dex lower than our Galaxy for the elements C, N, O, Mg and Si.  Additionally the abundances of O, Mg and Si were found to be 0.5 to 0.7 dex lower in DI\\,1239 compared with AV\\,304. One of the principal areas of uncertainty in the Rolleston et al.\\ analysis was the use an LTE approximation.   \\begin{table}  \\caption{A non-LTE analysis of the ICR stars using TLUSTY. For each star, the newly derived atmospheric parameters are followed by their absolute abudances along with the differential abundances relative to AV\\,304. Where more than one lines were used in the analysis, the standard deviation of the sample is $\\le$ 0.3 dex.}\\label{tab:icrs}  \\begin{tabular}{clllllll} \\hline\\hline\\noalign{\\smallskip} && \\multicolumn{2}{c}{DI\\,1162} & \\multicolumn{2}{c}{DI\\,1239} & \\multicolumn{2}{c}{DGIK\\,975}\\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\teff\\/ [K]    &&  21\\,500  &           & ~~\\,24\\,500  &               & ~~\\,21\\,000  &               \\\\ log~$g$ [cm]   &&  3.25     &           & ~~\\,3.70\t&               & ~~\\,3.35     &               \\\\ \\vt\\/   [\\kms] &&  5        &           & ~~\\,5\t&               & ~~\\,5\t       &               \\\\ \\noalign{\\smallskip} C\\,{\\sc ii}     &&  7.2 ( 2) & --0.3 ( 2)& ~~\\,7.5      & ~~\\,+0.2\t& $<$ 6.8 ( 2) & $<$ --0.6 ( 2)\\\\ N\\,{\\sc ii}     &&  6.5      & --0.1     & $<$7.3 ( 2) & $<$ +0.5      & ~~\\,6.5      & ~~\\,--0.1     \\\\ O\\,{\\sc ii}     &&  7.8 ( 6) & --0.3 ( 6)& ~~\\,7.6 (11) & ~~\\,--0.5 (10)& ~~\\,7.9 (10) & ~~\\,--0.2 (10)\\\\ Mg\\,{\\sc ii}    &&  6.3      & --0.5     & ~~\\,6.3      & ~~\\,--0.4\t& ~~\\,6.0      & ~~\\,--0.8     \\\\ Si\\,{\\sc ii/iii}&&  6.5 ( 5) & --0.5 ( 3)& ~~\\,6.0 ( 3) & ~~\\,--0.8 ( 3)& ~~\\,6.2 ( 5) & ~~\\,--0.9 ( 3)\\\\ {\\bf Mean$^*$}  &&           &{\\bf $-$0.4 $\\pm$ 0.2} & & {\\bf $-$0.4 $\\pm$ 0.4}&& {\\bf $-$0.6 $\\pm$ 0.3} \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\end{tabular}\\\\ \\smallskip * Excluing N. \\end{table}  To investigate this, we have re-analysed the ICR stars in Rolleston et al.\\ (1999) using the same non-LTE grids used in this study. The atmospheric parameters were re-deduced using the same methodology adopted by Rolleston et al., i.e.\\  the effective temperature was derived using the Balmer discontinuity and the surface gravity from the comparison of the observed Balmer line profiles with the non-LTE theoretical spectra.  Using these new atmospheric parameters, we have derived the absolute non-LTE abundance of the ICR stars and detailed line-by-line differential abundances relative to AV\\,304 (summary in Table~\\ref{tab:icrs}).  While the non-LTE atmospheic parameters are different from their LTE counterparts, the derived absolute abundances agree very well with the corresponding LTE values within the error. So do the differential abudances relative to the SMC standard AV\\,304.  More importantly, the metal deficiency seen in Rolleston et al.\\ is still evident in our non-LTE analysis. With an exception of C and N for which the ICR stars show both over- and under-abundances, other metal lines (O, Mg and Si) all show similar amount of under-abundace as seen in Rolleston et al.  Rolleston et al.\\ reported the ICR stars have metallicity $\\sim -$0.5 dex lower  than that of the SMC.  Excluding N, which might have been contamined from the dredged-up material, the average of our differential analysis gives a similar value.  The other main uncertainty is probably the choice of the microturbence. This quantity is difficult to estimate and its choice would affect the stronger metal lines in the Galactic stellar spectra more than those for the ICR targets. Hence, for example, adoption of larger microturbulences would reduce the magnitude of the underabundance found in the ICR. However, the microturbulences found by Rolleston et al.\\ from the \\ion{O}{ii} spectra range from 5 to 10 km s$^{-1}$ and appear reasonable for the stellar atmospheric parameters (see, for example, Hardorp and Scholz \\cite{Har70}; Gies \\& Lambert \\cite{Gie92}; Kilian et al.\\ \\cite{Kilian}).  Lehner et al.\\ (\\cite{Leh01}) has studied the interstellar ultraviolet absorption and \\ion{H}{i} emission spectrum towards a young star in Magellanic Bridge. These observations also appeared to support a low metallicity ($\\simeq -1.1$ dex) compared to that of the  Galaxy. Hence, although the failure to identify a low metallicity component in the SMC Wing is worrying, there would appear to be no substantial reason to discount the results of Rolleston et al.\\ (\\cite{Rol99}). Clearly further observations of targets in both the ICR and the SMC/ICR interface are crucial to understanding this discrepancy."
    },
    "0410/astro-ph0410404_arXiv.txt": {
        "abstract": "We present a performance analysis of SparsePak and the WIYN Bench Spectrograph for precision studies of stellar and ionized gas kinematics of external galaxies. We focus on spectrograph configurations with echelle and low-order gratings yielding spectral resolutions of $\\sim$10000 between 500-900nm. These configurations are of general relevance to the spectrograph performance.  Benchmarks include spectral resolution, sampling, vignetting, scattered light, and an estimate of the system absolute throughput.  Comparisons are made to other, existing, fiber feeds on the WIYN Bench Spectrograph. Vignetting and relative throughput are found to agree with a geometric model of the optical system. An aperture-correction protocol for spectrophotometric standard-star calibrations has been established using independent WIYN imaging data and the unique capabilities of the SparsePak fiber array. The WIYN point-spread-function is well-fit by a Moffat profile with a constant power-law outer slope of index -4.4. We use SparsePak commissioning data to debunk a long-standing myth concerning sky-subtraction with fibers: By properly treating the multi-fiber data as a ``long-slit'' it is possible to achieve precision sky subtraction with a signal-to-noise performance as good or better than conventional long-slit spectroscopy. No beam-switching is required, and hence the method is efficient. Finally, we give several examples of science measurements which SparsePak now makes routine. These include H$\\alpha$ velocity fields of low surface-brightness disks, gas and stellar velocity-fields of nearly face-on disks, and stellar absorption-line profiles of galaxy disks at spectral resolutions of $\\sim$24,000. ",
        "introduction": "\\subsection{The SparsePak Design}  In Paper I (Bershady \\etal\\ 2003) we describe the design, construction and laboratory calibration of SparsePak, a two-dimensional, formatted fiber array for the WIYN Telescope's Bench Spectrograph. \\footnote{The WIYN Observatory is a joint facility of the University of Wisconsin-Madison, Indiana University, Yale University, and the National Optical Astronomy Observatories.}  SparsePak is optimized for kinematic studies of nearby galaxies, and is designed to be a survey engine of stellar velocity dispersions in normal galaxy disks. SparsePak's size and pattern, combined with the versatility of the Bench Spectrograph, enable a wide range of science programs.  The technical premise for SparsePak was a fiber-optic feed with a two-dimensional sampling geometry and large light-gathering power for an existing spectrograph capable of yielding medium spectral resolutions ($5000<\\lambda/\\Delta\\lambda<20,000$) with large apertures. Within these constraints, SparsePak's design was optimized for galaxy kinematic studies to sample a large solid angle with these performance goals: (a) Velocity dispersions of $\\sigma \\sim $10 km s$^{-1}$ are resolvable (for reference, the stellar velocity dispersion of the old-disk stellar population in the Solar Neighborhood is roughly 20 km s$^{-1}$, Kuijken \\& Gilmore 1989); (b) such measurements are practical at the low surface-brightness levels found at several radial disk scale-lengths of nearby, normal spiral galaxies; and (c) observations are photon-limited at these faint light-levels -- at or below the continuum sky-brightness.  Meeting these performance goals is difficult. The sky continuum at 550nm yields only $2 \\times 10^{-2}$ photons s$^{-1}$ m$^{-2}$ arcsec$^{-2}$ \\AA$^{-1}$ at the top of the atmosphere at a dark site ($V = 21.8$ mag arcsec$^{-2}$). Each of SparsePak's 82 fibers subtends over 17 arcsec$^2$ in solid-angle. This is necessary for achieving photon-limited performance at medium spectral resolutions on a 4m-class telescope. Still, the demands on the spectrograph are severe: The spectrograph must be very efficient while achieving high dispersion and large demagnification. The Bench Spectrograph meets some, but not all of these requirements.  Since the spectrograph existed first, its capabilities drove the choice of the SparsePak fiber size. This, in turn, limits the spatial resolution and field-of-view of SparsePak, and hence the class of objects for which SparsePak is well suited.  SparsePak's $72 \\times 71$ arcsec sparsely-sampled ``source'' grid contains 75 of the 82 fibers. In this grid, fibers are spaced center-to-center by 9.85 arcsec, with the exception of an inner core of 17 fibers with a 5.6 arcsec spacing. Seven sky-fibers are placed 60 to 90 arcsec away from the center of the grid. This geometry is a compromise between sampling and field-of-view. The individual fiber foot-print and grid size sample several radial scale-lengths in normal spiral galaxies with recession velocities of (roughly) 2,000-10,000 km s$^{-1}$. Given the hexagonal fiber packing geometry, the grid can be densely sampled or critically over-sampled with only three pointings.  \\subsection{Advantages of SparsePak Bi-Dimensional Spectroscopy for Galaxy Kinematics}  SparsePak's two-dimensional fiber geometry provides significant advantages over long-slit spectroscopic observations of extended sources for a range of applications. These advantages overcome several key limitations plaguing studies of galaxy kinematics:  \\begin{itemize}  \\item {\\bf Signal-to-Noise.} Compared to long-slit spectroscopy, SparsePak samples more area at larger radii (where, for example, galaxies grow fainter). The linear sampling of long-slit spectroscopy has been a central problem for the study of disk stellar velocity fields and dispersions (e.g., Bottema 1997), and has made precision measurements impractical at radii of 2-3 disk scale-lengths. It is critical to probe these radii because at these galacto-centric distances the disk is expected to make a maximum contribution to the total enclosed galactic mass. In contrast to long-slit observations, with SparsePak it is possible to co-add fibers within annuli at many radial intervals. This improves the signal-to-noise in a manner comparable to measuring surface-photometry in two-dimensional images (see, for example, Bershady \\etal\\ 2002 and Verheijen \\etal\\ 2003).  \\item {\\bf Position.} Uncertainties in centering and offsets (for source mapping) can be substantially reduced with SparsePak observations {\\it a posteriori} by reconstructing spatial maps of the source from the two-dimensional spectral data (see below). This is either impossible or impractical with long-slit spectroscopy.  \\item {\\bf Rotation.} Uncertainties in source position angles are inconsequential for SparsePak observations since all angles are sampled simultaneously.  \\end{itemize}  The uncertainties of centering and position angles (PA) are nagging problems for optical, long-slit observations of galaxy velocity fields.  These problems are addressed by Fabry-Perot (FP) observations of ionized gas, but only for limited samples of galaxies (e.g. Schommer \\etal\\ 1993, Beauvais \\& Bothun 1999, Barnes \\& Sellwood 2003). For long-slit observations, all random errors in centering and position angle lead to systematic underestimates of the true rotation speeds of galaxies. This systematic effect leads to erroneous, systematic offsets in the Tully-Fisher relation, as derived from such data. The effects mimic -- and can be misinterpreted as -- an enhancement in the luminosity of a source for a given rotation speed (mass). Because of the steepness of the Tully-Fisher relation, a small systematic offset in velocity appear as large offsets in luminosity (e.g, a 10\\% error in rotation speed is equivalent to $\\sim$0.4 mag in luminosity, assuming a typical Tully-Fisher slope of 9-10 in the red--near-infrared; Verheijen 2001). The effects will be particularly pronounced for galaxies which display significant non-axisymmetric structure in either their light distribution or velocity fields. Such galaxies are prevalent among the later-type, lower-luminosity galaxies in the nearby universe, and appear to dominate at all luminosities at higher redshifts. Hence, a precise measurement of the rotation speed is essential for using, for example, the Tully-Fisher relation as a diagnostic of galaxy luminosity evolution as a function of mass.  The effects of poor spatial resolution or beam-smearing also degrade derived kinematic information. HI studies of galaxies can extend two-dimensional synthesis-maps beyond the optically detectable disks, but are both expensive in exposure time and limited in spatial resolution compared to optical observations of ionized gas. FP measurements are clearly a competitive option if the scientific premium requires the highest spatial resolution. SparsePak delivers coarser spatial sampling than most FP instruments, but it still offers finer sampling than most HI synthesis telescopes, and at much higher signal-to-noise. Because of its decreased spatial resolution, SparsePak provides significantly more light gathering power and hence higher sensitivity and efficiency for low-flux applications. This advantage is demonstrated in the rapid acquisition of samples for pilot studies of the ionized gas kinematics of barred spirals (Courteau \\etal\\ 2003), low surface-brightness disks (Swaters \\etal\\ 2003), and asymmetric systems (Kannappan \\etal\\ 2005).  While SparsePak is not a panacea for kinematic studies at all angular scales, the integral-field approach to spectroscopy is: For extended sources, high spatial resolution FP or integral-field spectroscopy (IFS) is superior to long-slit observations since 2-dimensional velocity fields provide independent spatial information not necessarily gleaned from photometry. In other words a ``well aligned slit'' may not always be meaningful. As noted in Paper I the overall efficiency of IFS vs. stepped, long-slit observations is significantly higher, and free of spatially-dependent systematic errors due to telescope offset problems or changing conditions. Unlike centering and PA mismatch, for SparsePak the problem of spatial scale can be solved a priori by choosing targets of the appropriate size or working on radial scales where beam-smearing is unimportant.  As FP observations of ionized-gas velocity fields have come of age (e.g., Palunas \\& Williams 2000, Garrido \\etal\\ 2002 and references therein), only one galaxy to date has been similarly mapped in stellar absorption (Debattista \\& Williams 2004). IFS has a greater combination of simultaneous spectral coverage and spectral resolution than FP observations. As such, IFS should be superior for stellar absorption-line kinematic observations of any kind because a significant range of wavelengths is needed for measuring cross-correlations. For example, the SAURON IFS instrument has measured stellar velocity fields for 72 early-type (E-Sa) galaxies, but only out to one effective radius and at spectral resolutions of $\\sim$1400 (de Zeeuw \\etal\\ 2002).  The advantage of IFS measurements is particularly pronounced for obtaining the highest spectral resolution data, since this requires sampling many weak (hence narrow) lines. In comparison to FP instruments and SAURON, as of the completion of this paper, SparsePak has been used to measure over a dozen stellar velocity and velocity-dispersion fields for late-type spiral galaxies out to 2-4 radial scale-lengths at spectral resolutions of $\\sim$10,000. This medium spectral-resolution kinematic mapping out to low surface-brightness is a unique capability of SparsePak.  \\bigskip  This paper reports the on-sky performance of the SparsePak and Bench Spectrograph. The detailed and systematic nature of the calibration measurements taken during commissioning allows us to (i) determine if our design goals are met, and (ii) establish a benchmark of SparsePak's present capabilities.  We focus on the performance of the Bench Spectrograph for a subset of gratings relevant to galaxy kinematic studies. Basic performance is quantified in terms of spectral resolution and throughput, the trade-offs between the two, and a measure of the absolute throughput of the Bench Spectrograph. The latter required a determination of the WIYN stellar point-spread function and the definition of aperture corrections as a function of internally-estimated seeing conditions. The final outcome yields a well-calibrated scheme for precision spectrophotometry with SparsePak.  We use SparsePak data to illustrate superior techniques for fiber-spectra data analysis.  We give attention to the optimization of spectral extraction and sky subtraction -- both topics which have received considerable attention in the literature. We show that problems with fiber-spectrograph sky subtraction is likely caused by poorly designed algorithms based on misconceptions about the data. We argue the results, presented here in the context of SparsePak, will be of general benefit to a wide range of fiber-fed instruments and science applications.  Finally, we give examples of commissioning science which highlight the new capabilities of SparsePak. Four significant results include stellar and ionized gas velocity fields in nearly face-on galaxies; a demonstration of the ability to register two-dimensional fiber-data to broad-band CCD images with sub-arcsecond positioning precision; a quantitative assessment of the relative merits of different spectral regions for stellar absorption-line kinematic measurements of systems with small internal velocities; and estimates of the limiting performance of SparsePak for stellar-kinematic studies of galaxy disks.  Commissioning observations and spectrograph configurations are described in \\S 2. The derived system resolution and sampling are found in \\S 3. The total system throughput, spectrograph vignetting, and fiber-to-fiber variations are presented in \\S 4. We examine the scattered light within the spectrograph in \\S5, and define methods of optimal spectral extraction in the specific context of SparsePak data. We demonstrate the ability to perform superior sky subtraction without the use of beam-switching in \\S 6.  Examples of early commissioning science of emission-line and absorption-line galaxy kinematics are presented in \\S 7. We conclude in \\S 8 with a summary of the most significant results. Five Appendices contain spectrograph characteristics; the system throughput budget; an aperture-correction protocol for spectrophotometry (including characterization of the WIYN point-spread-function); a new sky-subtraction algorithm; and an analysis tool for determining rapidly the kinematic PA of a galaxy. ",
        "conclusions": "We have presented the capabilities of SparsePak and the WIYN Bench Spectrograph in several configurations relevant for the study of galaxy kinematics, established procedures for conducting precision spectrophotometry of extended sources, and quantified and understood the throughput efficiency of the Bench Spectrograph. We have used SparsePak to implement significantly improved methods for sky subtraction. In this regard, performance for this fiber-fed system is comparable to long-slit, imaging-spectroscopic instruments. Finally, we have demonstrated SparsePak's on-sky capabilities based on observations of two, nearly face-on galaxies. Here we summarize each of these components.  The Bench Spectrograph offers a wide range of possible spectral resolutions across the visible band-pass, with trade-offs between resolution and system efficiency. We have focused on the high-end of the resolution range achievable with SparsePak: $\\Delta\\lambda/\\lambda\\geq5000$. Resolutions between 9,700 and 12,000 are typical with the echelle grating for wavelength between 500-900nm. Resolutions as high as 20,000-24,0000 are possible using the echelle in off-order (high incidence-angle) configurations.  These modes are 50\\% less efficient due to overfilling the grating and lower diffraction efficiency. Gain factors in resolution roughly equal loss factors in throughput. Resolutions of 4,000-6,000 are possible with 2nd-order gratings. Due to the smaller camera-grating distances used for the low-order gratings, made possible by larger camera-collimator angles, off-axis vignetting is $\\sim$20\\% smaller than for typical echelle configurations. This, combined with higher diffraction efficiencies make the low-order grating configuration the most efficient, but at the price of lower spectral resolution. The above resolutions are specific to SparsePak, but the trade-offs between spectral resolution and signal-to-noise are generic for the spectrograph.  In all cases we have explored, the spectral sampling is only 2.5-3.5 pixels (FWHM) even with SparsePak's large fiber diameters due to the large geometric and anamorphic demagnifications in the spectrograph. Spatial sampling is $\\sim$4 pixels (FWHM). The fiber spacing (roughly 10 pixels) is such that on-chip binning by a factor of 2 can be used in this dimension without signal degradation. This is useful for low-light-level applications where detector noise is significant.  The combined absolute throughput of the telescope, SparsePak and spectrograph has been established, based on measurements of a spectrophotometric standard star, to be 7\\% peak at 670nm, and 4\\% mean (over all used field angles -- fibers and wavelengths).  Our ability to measure reliably the absolute throughput is due to SparsePak's large fibers and array geometry. Large fibers result in small slit-losses for observations of stellar spectrophotometric standards. The two-dimensional format of the array has allowed us to develop a second-order, empirical formulation of aperture corrections with a precision better than 2\\%.  The modest system throughput stems in part from the exhibited, strong spatial vignetting function. We have well-matched this vignetting function with a geometric model that traces a realistic (laboratory-measured) fiber-output beam-profile through the optical system. On this basis we have been able to conclude (a) the vignetting is due to the lack of proper pupil placement within the spectrograph and large distances between collimator, grating and camera; and (b) averaged over all fibers, typically half of the light within the spectrograph is lost to vignetting at the central wavelength.  It is remarkable to ponder where the photons are lost overall. A detailed breakdown is documented, and plausibly understood.  The fiber introduce only a 20\\% loss, and are a minor contributor to the overall budget. By our estimate, half of the light is lost before it gets to the spectrograph (``top end'' losses due to atmosphere, 3 aluminum surfaces, and fiber surface losses and internal attenuation); half of this light is lost in geometric vignetting within the spectrograph (above); another factor of $\\sim2$ decrease comes from filters and gratings; and a final factor of $\\sim2$ decrease comes from surface-losses on spectrograph optics, dewar window, and CCD quantum efficiency.  We have grouped these losses together in this way purposefully. There is little that can be done to any single component in the first and last groupings to significantly increase the throughput; attention to improving many components is needed to make appreciable gains. In contrast, the middle two groupings offer the possibility of making significant throughput gains by shortening the collimator focal-length, properly placing the pupil (see Paper I), and introducing more efficient gratings (e.g., volume-phase holographic gratings, for which order-blocking filters are generally not needed). While collimator changes may result in decreased spectral resolution, the amplitude of the decrease would likely be of order or less than 30\\%, which would still be sufficient for galaxy kinematic studies. Since photon starvation is the critical limit, the trade of efficiency for spectral resolution is desirable.  A program to improve spectrograph throughput along these lines is underway and will be reported elsewhere.  We have also explored the implications of the fiber spacing for scattered light and optimal extraction of the spectral signal. The fibers are regularly spaced and sufficiently well separated for a clean beam extraction, with some degradation in the far red, CaII-triplet region at 860nm. At bluer wavelengths, cross-talk is below 1\\% when using extraction apertures capturing $\\sim$90\\% of the desired signal. At 860nm, the cross-talk rises to $\\sim$10\\% for the same captured signal. For SparsePak data, we find that weighted extractions offer little gain over an unweighted extraction, and therefore we advocate the unweighted extraction. An extraction threshold near 50\\% the peak is optimum.  One of the most striking results from the present work is the quality of the sky subtraction with the SparsePak array. This is due partly to the even placement of the sky fibers along the SparsePak slit, and possibly because the large fibers are well resolved and their monochromatic images are not dominated by the spectrograph OTF. Also important is our development of a new algorithm for data handling and fitting of the sky-fibers. The algorithm includes subtraction of spectral continuum and sky line-emission in separate stages. In the highest-performance implementation of this algorithm, a low-order function is fit to all fibers in each spectral channel to model the effects of field-dependent optical aberrations. By including an interactive clipping algorithm to remove source flux, we have shown that in certain cases almost all of the fibers can be used for sky subtraction, resulting in significant improvement in the S/N of the data. These cases include any situation where discrete source emission in the fibers has a large range of apparent Doppler shifts, e.g., in redshift surveys or galaxy kinematic studies where the spread of internal velocities is large. Here the performance appears to be as good as one might expect from beam-switching or nod-and-shuffle, but our method is twice as efficient.  Two other significant implications stem from our results on sky subtraction. First, there is a need to revisit the optimum number of sky fibers for survey work. We conclude more sky-fibers are needed than traditionally calculated from simple random-error analysis models because such analyses do not consider the effects of field-dependent optical aberrations. This is a critical issue for subtracting unresolved sky-lines. Second, we have realized the advantage of fiber-fed spectrograph over slitlet systems for ``multi-object'' spectroscopy.  The ability of the former to control the mapping of the telescope to spectrograph focal-planes allows for better -- or at least more efficient -- control over optical aberrations for sky-subtraction purposes. In a very general sense, this offsets the information loss introduced by FRD in the design and use of a spectrograph.  We have also given several examples of the types of science for which SparsePak was designed.  We have shown SparsePak easily generates data yielding high-resolution emission-line and stellar absorption-line velocity fields of nearly face-on galaxies. It appears SparsePak has a higher throughput than DensePak by about 20\\% -- above and beyond the 3$\\times$ gains from increased fiber solid-angle. A simple tool can be used to examine such SparsePak data and quickly assess, by eye, the kinematic position-angle of a galaxy.  SparsePak's sampling geometry also permits the {\\it a posteriori} registration of spectral data to normal CCD images via analysis of the spatial distribution of the spectral continuum. The precision of this registration is better than one-tenth the fiber diameter.  Our primary motivation for building SparsePak, a large-etendue, two-dimensional fiber array, capable of achieving resolutions ($\\sigma$) of 10 km s$^{-1}$ or better, is to measure the stellar motions and velocity dispersions in galaxy disks.  An analysis of kinematic data of a K giant star and the high-surface brightness blue galaxy (NGC 3982) indicates that this is indeed feasible.  To measure stellar velocity dispersions at radii of 2-3 disk scale-lengths in a ``normal'' disk requires $\\sim10$ hours. Given the paucity of such data, even these lengthy integrations will yield important results worthy of the effort. Bench Spectrograph throughput improvements will enable large-scale surveys. Due to the combination of larger intrinsic line-widths, higher backgrounds, and greater scattered light in the CaII-triplet region, we conclude the MgI region is preferable for stellar kinematic studies.  For emission-line galaxies the MgI regions also offers the [OIII]$\\lambda$5007 line within the same spectral window. Therefore in one setting it is possible to trace stellar and gaseous velocity fields, and hence the asymmetric drift, as well as line-of-sight velocity dispersions. Such measurements offer great opportunities for studying the dynamics and mass distributions of spiral galaxies.  SparsePak is now a facility instrument available to the public on the WIYN telescope. A web page on this instrument is maintained at: \\\\ http://www.astro.wisc.edu/$\\sim$mab/research/SparsePak/."
    },
    "0410/astro-ph0410722_arXiv.txt": {
        "abstract": "We present a structural analysis of nearby galaxies in spiral-spiral pairs in optical $BVRI$ bands and compare with the structures of isolated spiral galaxies and galaxies in ongoing mergers.  We use these comparisons to  determine how galaxy structure changes during galaxy interactions and mergers. We analyze light concentration ($C$), asymmetry ($A$), and clumpiness ($S$) parameters, and use the projections of $CAS$ parameter space to compare these samples. We find that the $CAS$ parameters of paired galaxies are correlated with the projected separations of the pair. For the widest and closest pairs, the $CAS$ parameters tend to be similar to those of isolated and ongoing major mergers (ULIRGs), respectively. Our results imply that galaxy $CAS$ morphology is a robust quantity that only changes significantly during a strong interaction or major merger. The typical time-scale for this change in our paired sample, based on dynamical friction arguments, is short, $\\tau \\approx 0.1 - 0.5$ Gyr. We find average enhancement factors for the spiral pair asymmetries and clumpiness values of $\\sim 2.2$ and 1.5. The $S$ parameter, which is related to star formation activity, has a moderate level of enhancement suggesting that this activity in modern spirals depends more on internal processes than on external conditions.  We furthermore test the statistical criterion for picking up interacting galaxies in an automated way by using the $A-S$ projection plane. The diversity of our spiral pair sample in the $CAS$ space suggests that structural/SF/morphological properties of interacting galaxies change abruptly only when the interaction becomes very strong and that the criteria for finding galaxies involved in major mergers from Conselice (2003) is effective. ",
        "introduction": "\\label{S1}  One of the major unsolved questions in modern astronomy is understanding galaxy formation and evolution. The hierarchical clustering scenario, motivated by the inflationary Cold Dark Matter (CDM) cosmology, has served as the main theoretical background to address this question (for a recent review see Firmani \\& Avila-Reese 2003). From the observational side, two types of information have been used: (i) detailed studies of local galaxies interpreted as fossil records of the formation process, and (ii) the rapidly increasing observational data of the universe at large redshifts. An important question regarding the latter item is the identification of properties which can be fairly compared among the galaxy populations at different redshifts. The correct identification and measurement of these properties is essential to infer empirically how the present-day galaxy populations formed.   Some optical morphological features along the Hubble-sequence have been used as the primary characteristic to describe local normal galaxies. Unfortunately, these features describe only the ``tip of the iceberg'' in galaxies. Some morphological features are thought to be related to transient internal dynamical and luminous (star formation) processes rather than to long-term fundamental physical properties of galaxies (van den Bergh 1998; Conselice 2003 and references therein). It has also been shown that gross morphological characteristics change with wavelength \\citep{Blo99}, making it even more difficult to use them as a basis for a physical classification of galaxies. On the other hand, a high fraction of distant galaxies look peculiar and cannot be placed onto the Hubble sequence \\citep{Gla95, Abr96, Con04} suggesting that there must be a physical cause for this change (Conselice 2003). These facts have motivated the search for quantitative physical indices, which would allow galaxies to be classified in different stages of their evolution \\citep{Ber00}. In this spirit, the latter authors and Conselice (2003) have proposed three structural indices to distinguish galaxies at different stages of evolution, namely the concentration of stellar light ($C$), the asymmetry in the light distribution ($A$), and a measure of its clumpiness ($S$).   According to the galaxy formation picture based on the hierarchical CDM cosmogony, disk galaxies assembly generically inside CDM halos. The main properties of the disks are determined largely by the halo mass aggregation history (and its related concentration), and by the halo angular momentum \\citep{Mo98,Fir00,Bos00,Zav03}. On the other hand, in the hierarchical picture, the morphological characteristics of galaxies may change several times during its evolution \\citep{Kau93,Bau96}.  Therefore, the structure of a galaxy at a given time reveals its present and past formation history.  The $CAS$ morphology system was designed and tested to reveal the major properties of galaxies (Conselice 2003). The $C$ parameter is related to the scale of a galaxy (mass, size) and its halo angular momentum, in particular the spin parameter, which is roughly constant in time for most of the halos. $C$ could slightly change (increase) with time depending on the galaxy mass accretion rate and relaxation history. Strong changes are expected only when dissipative major mergers happen (when a spheroid forms, for example). The $S$ parameter, which traces recent star formation (SF) \\citep{Con03}, is expected to change the most with time for a given galaxy as SF occurs. The SF rate history of isolated normal disk galaxies is predicted to increase smoothly, on average by a factor of $1.5-4$, from $z=0$ to $z\\sim 1-2$ and then to decrease slightly at higher redshifts \\citep{Avi03}. However, interactions may abruptly enhance SF (Hern\\'andez-Toledo et al. 2001), increasing $S$ in this case. Finally, the parameter $A$ is expected to change significantly during interactions and mergers \\citep{Con00a}.   The $CAS$ parameters therefore may trace out fundamental physical stages of galaxy evolution, as well as possible transient processes related to interactions. In the hierarchical scenario the rate of halo major mergers increases dramatically with redshift \\citep{Got01}. Observations show that in the local universe only about $2-8\\%$ of galaxies are in interacting pairs \\citep{Xu91,Patton00}, but this fraction increases significantly with redshift as shown for a galaxy sample up to $z\\sim 0.5$ \\citep{Patton02}. From these data and assuming short merging time scales ($<0.5$ Gyr), the latter authors infer that the merging rate increases as $(1+z)^{2.7}$. Thus, at $z\\sim 3$, the merging rate would be $\\sim 40$ times the present one. Interestingly, at $z\\sim 3$ the observed fraction of galaxies with asymmetric features typical of a recent merger was found to be as large as 50\\% for the most luminous and massive sources (Conselice et al. 2003a,b). The detailed study of the local interacting pairs is crucial for understanding the structures of high-redshift galaxies. In this spirit, the aim of the present paper is to study $CAS$ parameters and correlations for a local sample of spirals in pairs, and compare them with the $CAS$ parameters for a reference sample of non-interacting galaxies and a sample of merging galaxies, namely ultra luminous infrared galaxies (ULIRGs).  This further allows us to determine what physical processes are required to change the morphologies of galaxies.  We explore how different stages of interactions affect the structure, morphology, and SF properties of galaxies as revealed through $CAS$ parameters. In a subsequent paper \\citep[Paper II]{Her04}, the correlations of the $CAS$ parameters with other galaxy properties in pairs will be explored in order to gain insight on the significance of these parameters for understanding galaxy interactions and evolution.   To address these questions, we use a set of broad-band $BVRI$ observations of 66 spirals in interacting pairs \\citep{Her01}, a set of F555W/F606W and F814W HST observations for 66 ULIRGs (Conselice 2003 and references therein), as well as a comparison sample of $\\sim 90$ bright, large non-interacting local galaxies from the catalog of Frei et al. (1996).   The structure of the paper is as follows.  Section~\\ref{S2} summarizes the main characteristics of the samples that are relevant to this study. In \\S ~\\ref{S2.1} an overview of the $CAS$ parameters and their measurements is given. In Section 3 we present the $CAS$ parameters in the $B$, $V$, $R$ and $I$ bands for the (S+S) pairs and compare their loci in the $CAS$ volume with that of the reference samples of isolated galaxies and ultraluminous infrared galaxies (ULIRGs). The results are discussed and interpreted in \\S ~\\ref{S4}. In \\S ~\\ref{S5} we present our conclusions. ",
        "conclusions": "\\label{S5}  We have measured $CAS$ parameters in $BVRI$ bands for a sample of 66 galaxies in (S+S) pairs (Hern\\'andez-Toledo \\& Puerari 2001) and carried out a comparison analysis in two steps: 1) an inter-band internal comparison of the $CAS$ parameters in (S+S) pairs and 2) an $R$ band comparison in the $CAS$ structural space among (S+S) pairs and two samples with similarly compiled parameters, the (S)Frei control sample of isolated galaxies and 66 ultraluminous infrared galaxies \\citep{Con03}. This has allowed us to gain insight on how morphology and SF changes as a function of interaction/mergers and recognize time-scales and physical effects that drive the evolution of these galaxy features. Our main conclusions are as follow:   \\noindent (i) The differences in $CAS$ parameters for paired galaxies from the $B$ to $I$ bands is in general and on average small, suggesting that the optical $CAS$ parameters are stable galaxy properties until the beginning of a strong interaction. The average concentration $C$ parameter decreases on average from $I$ to $B$ bands, consistent with the finding that galaxy light distributions are more extended  (less concentrated) in bluer bands than in red bands \\citep{dJ96}.   \\noindent (ii) The $CAS$ parameters of paired galaxies change depending on the separation parameter $SEP$, in the sense that the $CAS$ parameters for the widest and closest paired galaxies tend to be closer to those of isolated and ongoing major mergers (ULIR) galaxies, respectively. The mean values of the $A$ and $S$ parameters of the closest ($SEP <1$) paired galaxies are 2.2 and 1.5 times larger than those of the isolated galaxies, respectively, i.e., even in strongly interacting galaxies, the SF activity (related to $S$) does not increase too much, suggesting that this activity is more related to internal processes rather than external ones.  The deviations in the $CAS$ planes of pair galaxies with respect to the isolated ones increase significantly on average (more than $2\\sigma$) only for the closest galaxies and those galaxies with clear interaction signatures (Figs. 7 and 8). Nevertheless, there are also paired galaxies with $SEP < 1$ but within the $2\\sigma$ deviations, and pairs with $SEP > 1$ but with deviations larger than $2\\sigma$. This diversity suggests that galaxies in our sample of (S+S) pairs are in different dynamical and morphological stages, from non-interacting to strong interacting, where the projected separation $SEP$ is only an approximate (statistical) criterion of interaction.  \\noindent (iii) Trends in $CAS$ space for the interacting galaxies differ from trends seen between $CAS$ values of isolated galaxies. Within isolated systems lower concentration values ($C$) correlate with larger $A$ and $S$ parameters. For galaxies in close pairs, these trends disappear and even revert for the most interacting galaxies: for several of them the larger $A$ and $S$, the higher is $C$.  The concentration index for strongly interacting galaxies may become higher owing to a ``flattening'' in the outer surface-brightness profile produced by interaction-induced {\\it extended} SF. In the $A-S$ plane, paired galaxies, in particular the closer ones, increase their $A$ with increasing $S$ more steeply than isolated galaxies.  \\noindent (iv) On average paired galaxies are slightly less concentrated, asymmetric, and patchy than ULIRGs, the significant difference being in the clumpiness. The latter display also a (statistically significant) larger scatter in the $CAS$ planes than the former. The major difference in trends between paired and ULIRG galaxies is in the $A-S$ plane: the most patchy ULIRGs have smaller asymmetries and slightly lower concentrations than the most patchy paired galaxies. The structural/SF/morphological properties of interacting galaxies change significantly only when the interaction becomes very strong.   \\noindent (v) A criterion for picking up interacting galaxies that is based only on $A>0.35$ is marginally acceptable for paired galaxies. In fact there are some interacting galaxies with $A<0.35$.  A more robust {\\it statistical} criterion uses both the $A$ and $S$ parameters: $|A-\\overline{A}_{\\rm isol}|>2\\sigma$, where $(A-\\overline{A}_{\\rm isol})$ is the deviation of the interacting-galaxy $A$ parameter from the predicted one from the linear regression to the isolated sample in the $A-S$ diagram, and $\\sigma$ is the scatter of this fitting.  For ongoing major mergers (ULIRGs) the $A > 0.35$ criteria is suitable.   The summary of this work is that the quantitative structural and SF parameters of galaxies seem to change significantly only during the strong interaction/merger phases and within $\\sim$ 0.2-0.5 Gyr. The use of the $CAS$ parameters to establish criteria for finding galaxies in strong interaction at high-redshift galaxies appears to hold up as it depends strongly on interaction stage (e.g., Conselice et al. 2003a,b) and changes quickly with time. Within the local sample of (S+S) pairs analyzed here, we find the entire range of structural types are represented from isolated-like structural indices to indices closer to ongoing major mergers. This suggests that morphology is a fairly robust quantify that tends to change rapidly within the major merger process on the order of hundreds of Myrs."
    },
    "0410/astro-ph0410208_arXiv.txt": {
        "abstract": "With presently known input physics and computer simulations in 1D, a self-consistent treatment of core collapse supernovae does not yet lead to successful explosions, while 2D models show some promise. Thus, there are strong indications that the delayed neutrino mechanism works combined with a multi-D convection treatment for unstable layers (possibly with the aid of rotation, magnetic fields and/or still existent uncertainties in neutrino opacities). On the other hand there is a need to provide correct nucleosynthesis abundances for the progressing field of galactic evolution and observations of low metallicity stars. The innermost ejecta is directly affected by the explosion mechanism, i.e.\\ most strongly the yields of Fe-group nuclei for which an induced piston or thermal bomb treatment will not provide the correct yields because the effect of neutrino interactions is not included. We apply parameterized variations to the neutrino scattering cross sections in order to mimic in 1D the possible increase of neutrino luminosities caused by uncertainties in proto-neutron star convection. Alternatively, parameterized variations are applied to the neutrino absorption cross sections on nucleons in the ``gain region'' to mimic the increase in neutrino energy deposition enabled by convective turnover. We find that both measures lead to similar results, causing explosions and a $Y_e>0.5$ in the innermost ejected layers, due to the combined effect of a short weak interaction time scale and a negligible electron degeneracy, unveiling the proton-neutron mass difference. We include all weak interactions (electron and positron capture, beta-decay, neutrino and antineutrino capture on nuclei, and neutrino and antineutrino capture on nucleons) and present first nucleosynthesis results for these innermost ejected layers to discuss how they improve predictions for Fe-group nuclei. The proton-rich environment results in enhanced abundances of $^{45}$Sc, $^{49}$Ti, and $^{64}$Zn as requested by chemical evolution studies and observations of low metallicity stars as well as appreciable production of nuclei in the mass range up to $A=80$. ",
        "introduction": "\\label{sec:introduction}   The problem of core collapse supernova explosions is an old one and the attempts to understand the mechanism have been ongoing for more than 30 years. The idea that a massive star proceeds through all burning stages from H to Si-burning, finally leading to the collapse of the resulting Fe-core to nuclear densities with formation of a neutron star has long been discussed \\citep[e.g.][]{Baade.Zwicky:1934,Oppenheimer.Snyder:1939,Arnett.Schramm:1973}. Since the 1960s the explosion mechanism has been related to neutrino emission from the hot collapsed core \\citep[e.g.][]{Colgate.White:1966,Bethe.Wilson:1985,Bethe:1990} interrupted by a period when it was speculated that the strength of the bounce at nuclear densities could permit shock waves with sufficient energies to lead to prompt explosions \\citep[e.g.][]{Baron.Cooperstein.Kahana:1985,Baron.Bethe.ea:1987}. However, introduction of previously neglected neutrino scattering processes were introduced (e.g. neutrino-electron scattering), which permitted the replacement of lost low energy neutrinos, led to a continuous energy leakage and to the death of the prompt shock within 10 ms after bounce \\citep[]{Bruenn:1989,Myra.Bludman:1989}.  Since then, and with the first neutrino detection from a core collapse supernova \\cite[SN1987A, see e.g.][]{Koshiba:1992,Burrows:1990}, the hope has been that further improvements would lead to successful explosions via energy deposition through neutrino and antineutrino captures on neutrons and protons ($\\nu_e + n \\rightarrow p + e^-$, $\\bar\\nu_e + p \\rightarrow n + e^+$). Two different paths were explored. 1.\\ Convective instabilities, but with still simplified neutrino transport, causing either (a) convective transport in the core and leading to higher neutrino luminosities \\citep[e.g.][]{Keil.Janka.Mueller:1996} or (b) higher energy deposition efficiencies in convective regions \\citep{Mayle.Wilson:1988,Herant.Benz.ea:1994,Janka.Mueller:1996,% Fryer.Warren:2004}. 2.\\ Improved neutrino transport schemes, leading to higher neutrino luminosities via the full solution of the Boltzmann transport equation for neutrino scattering and neutrino reactions \\citep{Mezzacappa.Bruenn:1993a,Mezzacappa.Bruenn:1993b,Messer.Mezzacappa.ea:1998}.  However, in recent years 1D spherically and 2D rotationally symmetric radiation-hydro calculations have not yet shown successful supernova explosions with the present knowledge of physical processes. \\citep{Rampp.Janka:2000,Mezzacappa.Liebendoerfer.ea:2001,% Liebendoerfer.Mezzacappa.ea:2001a,Liebendoerfer.Mezzacappa.ea:2001b,% Buras.Rampp.ea:2003,Janka.Buras.Rampp:2003,Hix.Messer.ea:2003,% Langanke.Martinez-Pinedo.ea:2003,Thompson.Burrows.Pinto:2003}. A very recent simulation of a 11.2~M$_{\\odot}$ core collapse shows a chance for successful weak explosions in a multi-D treatment with spectral neutrino transport \\citep{Janka.Buras.ea:2005}. This leaves us with two dilemmata. First, the fundamental problem that the supernova mechanism is still not understood. Second, there seems no way to predict the correct supernova nucleosynthesis yields. This is a problem in itself, but is also a limitation for the rapidly expanding field of galactic chemical evolution, which is being energized by the large amount of recent abundance observations from low metallicity stars \\citep[e.g.][]{Argast.Samland.ea:2002,Argast.Samland.ea:2004,% Sneden.Cowan.ea:2003,Cayrel.Depagne.ea:2004,Honda.Aoki.ea:2004,% Frebel.ea:2005}.  Supernova nucleosynthesis predictions have a long tradition \\citep[]{Woosley.Weaver:1986,Thielemann.Hashimoto.Nomoto:1990,% Woosley.Weaver:1995,Thielemann.Nomoto.Hashimoto:1996,% Nomoto.Hashimoto.ea:1997,% Nakamura.Umeda.ea:2001,Rauscher.Heger.ea:2002,% Chieffi.Limongi:2002a,% Chieffi.Limongi:2004,Umeda.Nomoto:2005}. But all of these predictions relied on an artificially introduced explosion, either via a piston or a thermal bomb \\citep[]{Aufderheide.Baron.Thielemann:1991} introduced into the progenitor star model. The mass cut between the ejecta and the remnant does not emerge from the simulations, but has to be determined from additional conditions. While the usage of artificially introduced explosions is justifiable for the outer stellar layers, provided we know the correct explosion energy to be dumped into the shock front (on the order of 10$^{51}$ erg seen in observations), it clearly is incorrect for the innermost ejected layers which should be directly related to the physical processes causing the explosion. This affects the Fe-group composition, discussed in detail in \\citet{Thielemann.Nomoto.Hashimoto:1996}, hereinafter TNH96, and \\citet{Nakamura.Umeda.ea:1999}, which was also recognized as a clear problem by \\citet{Chieffi.Limongi:2002a} and \\citet{Umeda.Nomoto:2002}. The problem is also linked to the so-called neutrino wind, emitted seconds after the supernova explosion, and considered as a possible source of the \\emph{r}-process to produce the heaviest elements via neutron captures \\citep{Takahashi.Witti.Janka:1994,Woosley.Wilson.ea:1994,Qian.Haxton.ea:1997,% Thompson.Burrows.Meyer:2001,Wanajo.Kajino.ea:2001,Terasawa.Sumiyoshi.ea:2002,% Thompson.Burrows.Pinto:2003}.  An indispensable quantity to correctly describe nucleosynthesis in the innermost ejecta is the electron fraction $Y_e=\\langle Z / A\\rangle$ in the layers undergoing explosive Si-burning. This $Y_e$ is set by the weak interactions in the explosively burning layers, i.e.\\ electron and positron captures, beta-decays, and neutrino or antineutrino captures. The dominant reactions in hot photodisintegrated matter, consisting mainly of neutrons and protons, are  \\begin{displaymath} \\begin{array}{l} \\nu_e + n \\rightleftarrows p + e^-\\\\ \\bar\\nu_e + p \\rightleftarrows n + e^+ . \\end{array} \\end{displaymath}  In section \\ref{sec:masscut} we will show that these reactions lead to significant changes in $Y_e$ toward equilibrium before the matter is ejected. We will also show that the resulting $Y_e$ in the innermost ejected layers is close to 0.5, in some areas even exceeding 0.5. This has been strongly postulated as a requirement in order not to violate abundance constraints from galactic evolution and solar abundances (TNH96).  The question arises how one could realistically simulate this behavior, given the existing problems with self-consistent explosions. Discussed improvements which could lead to successful supernova explosions are rotation and magnetic fields \\citep[e.g.][]{Thompson:2000,Thompson.Quataert.Burrows:2004} or uncertainties in neutrino opacities \\citep[see e.g.][]{Burrows.Reddy.Thompson:2004} or other microphysics properties. They would introduce additional mixing at the neutrino sphere and convective transport or change the neutrino luminosity via improved opacities. This indicates the two options for successful explosion as discussed above: (a) enhanced neutrino luminosities or (b) enhanced deposition efficiencies for neutrino capture in convective layers. These effects can be simulated in two ways: (a) Boosting the neutrino luminosity via a scaling (reduction) of the neutrino scattering cross sections on nucleons while keeping the electron/positron and neutrino/antineutrino capture cross sections on neutrons and protons at their original values. (b) Boosting the energy deposition efficiencies by enhancing the neutrino and antineutrino captures on neutrons and protons. Neither approachrepresents a self-consistent treatment, but no external energy is required to produce (i) a successful explosion with (ii) a consistently emerging mass cut between neutron star and ejecta. Moreover, our treatment guarantees that $Y_e$ is consistently determined by all weak interaction processes. ",
        "conclusions": "\\label{sec:conclusions}  Presently, self-consistent core collapse supernova simulations in 1D do not lead to successful explosions while 2D models show some promise. Remaining uncertainties in neutrino opacities and/or the expected strong influence of convection (due to hydrodynamic instabilities caused by entropy gradients and/or rotation and magnetic fields) are likely to change this result. They may lead either to higher neutrino luminosities or higher efficiencies of neutrino energy deposition via neutrino and antineutrino captures on nucleons. In order to examine the accompanied nucleosynthesis in 1D calculations of successful explosions, we performed simulations with variations in neutrino scattering cross sections on nucleons and/or neutrino and antineutrino captures on neutrons and protons. In both cases successful explosions emerge with an interesting evolution of the $Y_e$ gradient in the innermost ejecta, which were followed up by a postprocessing for nucleosynthesis purposes.  The detailed nucleosynthesis calculations with a consistent treatment of all weak interactions show an electron fraction $Y_e>0.5$, i.e.\\ a slightly proton-rich environment with relatively high entropies of up to $\\sim$50 $k_B$ per nucleon. This causes complete Si-burning with an alpha-rich (and proton-rich) freeze-out. About 0.0007 M$_\\odot$ of hydrogen remain in the innermost ejecta and do not stem from mixing this matter in from the hydrogen envelope. Such a proton-rich environment at relatively high entropies permits to produce also nuclei beyond $A$=64, up to $A$=80, with a major contribution to $^{64}$Zn. The rp-like process does not extend to masses beyond A=80--100 as the high entropies imply too small densities for a path at very small proton separation energies.  In addition, we find improvements within the Fe-group. The strong overabundances of $^{58,62}$Ni found in previous (too neutron-rich) environments are reduced. $^{45}$Sc and $^{49}$Ti are enhanced to permit predictions closer to solar proportions. Especially the emergence of $^{45}$Sc seems to be a solution to the previously not understood abundance of this only stable isotope of Sc. This discussion is also of interest with respect to $^{44}$Ti, made in the alpha-rich freeze-out in the inner explosive ejecta. $^{44}$Ti is sensitive to  $Y_e$ and reduced in the mass range where $^{45}$Sc is high. This will influence the overall predictions of $^{44}$Ti.  Values of $Y_e>0.5$ are due to the neutrino interactions with matter under electron non-degenerate conditions in a convectively unstable domain and thus related to the explosion mechanism. The effect of neutrinos decreases with $1/r^2$ and about 50-60\\% of the Fe-group ejecta (the outer part of explosive, complete Si-burning) is determined by values of $Y_e$ equal to or close to the initial values inherited from stellar evolution. In this first study we have tried to give an estimate for the overall Fe-group composition based on such a superposition of the present results for the innermost ejecta with those of TNH96. Future investigations will require to perform full nucleosynthesis calculation for complete stars based on these exploding models. They will also require a sensitivity test of the nucleosynthesis results to the scaling factors for neutrino-induced reactions discussed in Figure \\ref{fig:evol} and in combination with the position of the mass cut and the explosion energy. This should be considered in order to reproduce results for supernovae where detailed observational information in abundances, gamma-ray emitters and explosion energies is available."
    },
    "0410/astro-ph0410664_arXiv.txt": {
        "abstract": "{We investigate the thermal emission from magnetic neutron star surfaces in which the cohesive effects of the magnetic field have produced the condensation of the atmosphere and the external layers. This may happen for sufficiently cool ($T\\le 10^6$) atmospheres with moderately intense magnetic fields (about $10^{13}$ G for Fe atmospheres). The thermal emission from an isothermal bare surface of a neutron star shows no remarkable spectral features, but it is significantly depressed at energies below some threshold energy. However, since the thermal conductivity is very different in the normal and parallel directions to the magnetic field lines, the presence of the magnetic field is expected to produce a highly anisotropic temperature distribution, depending on the magnetic field geometry. In this case, the observed flux of such an object looks very similar to a BB spectrum, but depressed by a nearly constant factor at all energies. This results in a systematic underestimation of the area of the emitter (and therefore its size) by a factor 5-10 (2-3). ",
        "introduction": "Almost a decade after the discovery of the soft X-ray source RXJ185635-3754 (\\cite{WWN96}) using the ROSAT-PSPC, the thermal component associated with the direct emission from a neutron star's surface has been detected in more than 20 X-ray sources. In many cases it is superimposed on a power--law tail, but seven of these objects are well--characterized as simple blackbodies with temperatures ranging between 60 and 100 eV. The apparent small emitting surface of RXJ185635-3754, inferred from the best blackbody fit and its parallax (\\cite{WL02}) have led to speculation about its nature, and whether its apparent smallness can be considered as evidence that the object is a strange star: a self--bound object made of up, down and strange quarks (\\cite{Pons02,Drake02}), or a standard, misaligned pulsar. Although this new observational class (isolated compact stars) is probably the first real opportunity to place stringent constraints on the equation of state (EOS) of dense matter from astrophysical measurements (see \\cite{LP2001} for a review), one must be cautious before concluding that an apparently small X-ray source is a quark star, because the X-ray spectrum is not the only information available.  Using the Hubble Space Telescope, Walter \\& Matthews (\\cite{WM97}) subsequently identified an optical source at 6060 \\AA$\\,$ and 3000 \\AA, with a brightness only about 7 times brighter than an extrapolation of a 62 eV X-ray blackbody into the optical V band. The observed optical fluxes have been confirmed by subsequent observations from the 2-meter NTT (\\cite{Neu98}) and new HST measurements (\\cite{Pons02}). Remarkably, the other three isolated compact X-ray sources that have been detected in the optical band (RX J0720.4-3125, RX J1308.6+2127, and RX J1605.3+3249) also have a significant optical excess over the extrapolation of the X-ray blackbody (a factor 5 to 14). None of them have yet been detected as 1.4 GHz radio sources. Thus, RXJ185635-3754 is not an uncommon object, for it shares the same general observational properties of other isolated neutron stars (blackbody spectrum in X-ray, no evident spectral features, optical excess), except for the fact that four of them have fairly long periods (8-22 s), while RXJ185635-3754 is not variable, with a reported upper limit on the pulse amplitude of $<1.3\\%$ (\\cite{Bur03}).  Since blackbodies are no more than a simple approximation of the true emission mechanism, a step forward in understanding the thermal emission of neutron stars, consists of computing model atmospheres for low magnetic fields ($< 10^{11} G$) and emergent spectra for several compositions, masses and radius (\\cite{Rom87,Mil92,RR96}). The parameters that determine the shape of the observed spectrum are the atmospheric composition, the red-shifted temperature, $T_\\infty= T_{\\rm eff}/(1+z)$, the redshift $z$, the interstellar medium absorption, $n_H$, and the ratio $R_\\infty/d$, where $d$ is the distance to the object and $R_\\infty=R(1+z)$, $R$ being the radius of the star.  For a Planck spectrum the redshift factors contribute only to an overall scale factor, so that the redshift does not affect the models. For  realistic atmospheres with spectral features, however, the redshift is an important parameter because the identification of spectral lines would determine the $M/R$ ratio and establish relevant constraints on the EOS of dense matter. After a detailed investigation of the different parameters that affect the observed spectra, Pons et al. (2002) concluded that X-ray data alone do not allow us to establish severe constraints, given the large number of degrees of freedom. The combination of the X-ray spectra with the optical fluxes turned out to be much more restrictive, as was suggested earlier by \\cite{Pav96}. The broadband spectral energy distribution (SED) of isothermal heavy element atmospheres can fit the observed SED, but these models predict detectable absorption lines and edges. It was hoped that the more sensitive spectrographs on Chandra and XMM--Newton would be able to detect spectral features and help to understand the nature of this objects, but no lines or absorption edges have been seen in long exposures (\\cite{Bur01,Drake02}). Only the presence of broadband departures from a single blackbody spectrum has recently been suggested (\\cite{Bur03}).  The puzzle, then, is as follows: single temperature models based on heavy element atmospheres explain the broadband spectrum and the lack of variability of the X-ray spectrum, but they are in contradiction with the absence of spectral features. Single blackbody models, besides being unrealistic, cannot account for the systematic optical excess that is observed in 4 sources. An anisotropic surface temperature distribution (the simplest case would be two blackbody components) can reconcile X-ray and optical observations, because the quantity $(R_\\infty/D)^2$ must be interpreted as the solid angle subtended by the star's surface area which is visible at some distance $D$. Thus the true value of $R_\\infty$ will generally be larger if the assumption that the temperature is uniform on the surface is relaxed. However, this is barely consistent with the lack of pulsation of RXJ185635-3754, unless its magnetic and rotation axes are oriented in a very particular way.  Nevertheless, there is an important ingredient missing in the above discussion, the likely existence of strong magnetic fields. The strong (10$^{12-13}$G) magnetic fields of pulsars suggest that most, if not all, neutron stars should have similarly strong fields. However, detailed atmosphere models with strong magnetic fields are only available for hydrogen (\\cite{Zav95}), because reliable opacities and EOS have not yet been developed for heavier elements and because of the complexity of modeling magnetic atmospheres with arbitrary magnetic field structures.  For heavy-element dominated atmospheres, only approximate treatments of magnetic Fe atmospheres exist (\\cite{RRM97}), and the results show that the spectra are globally much closer to a blackbody than for light element atmospheres.  There is a case to be made that RX~J185635-3754 is magnetized. The deep VLT image of the target released by van Kerkwijk \\& Kulkarni (\\cite{KK00}) shows what looks like a classic bow-shock nebula. The presence of a bow-shock suggests that this is a magnetized neutron star with a relativistic wind, as seen in the pulsars PSR 1957+20 (Kulkarni \\& Hester \\cite{KH88}) and PSR 2224+65 (Cordes, Romani, \\& Lundgren \\cite{CRL93}).  An alternative possibility that explains naturally the absence of spectral features is the emission from a solid surface. This was a common idea 20-30 years ago (\\cite{Bri80}, B80 in the following) until the existence of a thin gaseous atmosphere was appreciated and model atmospheres became more popular. However, at sufficiently low temperatures, highly magnetized neutron stars may undergo a phase transition that turns the gaseous atmosphere into a solid (\\cite{Lai01}). The critical temperature below which the atmosphere condensates depends on the composition and the magnetic field. For example, for typical magnetic field strengths of $10^{13}$ G, a Fe atmosphere will condensate for $T< 0.1$ keV while a H atmosphere needs temperatures lower than 0.03 keV to undergo the phase transition to the {\\it metallic} state (\\cite{Lai01}). Notice that effective temperatures of the observed isolated neutron stars fall in this temperature range, therefore they should plausibly be in the solid state if the dominant element in the atmosphere is Fe. In such a metallic neutron star surface made of nuclei with atomic number $Z$ and atomic weight $A$, the pressure vanishes at a finite density \\begin{equation} \\rho_s \\approx 560~A Z^{-3/5} B_{12}^{6/5} {\\rm g\\, cm}^{-3} \\label{rhos} \\end{equation} where $B_{12}$ is the magnetic field in units of $10^{12}$ G.  This idea has been recently revisited by Turolla et al. (2004), who found that the emitted spectrum is strongly depressed at low energies, thus making more difficult the reconciliation between observational data and the condensed surface model. However, in this last work the emissivity is calculated neglecting one of the transmission modes in some cases. This simplification is inaccurate and can modify the emitted spectrum, as we will discuss in the text. Besides, the sole presence of a strong magnetic field, necessary to condensate the atmosphere, results naturally in an anisotropic surface temperature (\\cite{GKP04}), which must be calculated consistently with the magnetic field structure. In this paper our aim is to study the thermal emission from solid surfaces of neutron stars and its implications on the observational properties of isolated compact objects.  The paper is organized as follows: Sect. 2 is devoted to discussing the calculation of the emissivity starting from the description of the dielectric tensor for the condensed neutron star surface. In Sect. 3 we show results, including integrated spectra for different strengths and geometries of the magnetic field. Final remarks and a summary of the main conclusions are discussed in Sect. 4. ",
        "conclusions": ""
    },
    "0410/astro-ph0410452_arXiv.txt": {
        "abstract": "We report the first survey of chemical abundances in M and K dwarf stars using atomic absorption lines in high resolution spectra. We have measured Fe and Ti abundances in 35 M and K dwarf stars using equivalent widths measured from $\\lambda / \\Delta \\lambda \\approx 33\\,000$ spectra. Our analysis takes advantage of recent improvements in model atmospheres of low-temperature dwarf stars.  The stars have temperatures between 3300 and 4700 K, with most cooler than 4100 K. They cover an iron abundance range of $\\rm -2.44 < [Fe/H] < +0.16$. Our measurements show [Ti/Fe] decreasing with increasing [Fe/H], a trend similar to that measured for warmer stars where abundance analysis techniques have been tested more thoroughly.  This study is a step toward the observational calibration of procedures to estimate the metallicity of low-mass dwarf stars using photometric and low-resolution spectral indices. ",
        "introduction": "Low mass (M dwarf and cooler) main sequence stars are by far the most numerous stars in the Galaxy and make up most of its baryonic mass. However, there have been few detailed chemical abundance studies of these stars with spectra of sufficient resolution for atomic absorption lines to be measured individually. None of these included more than a few stars.  This was largely a result of their low intrinsic luminosity and the molecular lines present in their spectra. Their faintness meant that few of these stars were bright enough for high signal to noise, high resolution spectra to be measured easily without a large telescope. The molecular bands present in their spectra complicated the calculation of stellar model atmospheres and cause line blends which make it difficult to measure atomic line strengths in large regions of the visible spectrum.  There have been abundance studies which used photometry or which fitted broad molecular features in low-resolution spectra of M and K dwarfs, but their results are not as certain as they would be if they were calibrated with stars for which more precise abundances were based on higher-resolution spectra. There have also been studies which used synthetic spectrum fitting of low resolution spectra to measure metallicities of low mass dwarfs. These have given rough metallicity estimates, but have not produced elemental abundances with the precision which is possible using higher resolution spectra.  Recent advances in model atmospheres of low mass dwarf stars, largely as a result of the improved treatment of molecular opacity, have provided the opportunity to determine abundances of M and K dwarf stars through analysis of their atomic spectral lines.  Our study of chemical abundances in Kapteyn's Star (HD 33793) \\citep{ww04} showed us that it is possible to locate spectral regions in these stars where molecular bands are not present or are sufficiently weak for accurate equivalent widths of atomic lines to be measured.  The survey of abundance estimates we report here will provide some of the data necessary to calibrate metallicity indices for cool dwarf stars, for example the TiO and CaH indices of \\citet{rhg95} and \\citet{hgr96}.  With calibrated indices it should be possible to estimate the metallicity of a much larger number of cool dwarfs, including those too distant and faint to study using high resolution spectra.  When a large volume-limited sample is available, it will be possible to determine if the `G dwarf problem' \\citep{vdb62, at76} continues to lower-mass main sequence stars so that low-metallicity M dwarfs are more scarce than expected in our Galaxy. ",
        "conclusions": "We have calculated the Fe and Ti abundances of 35 M and K dwarf stars using atomic line equivalent widths measured using spectra with high spectral resolution and high signal-to-noise ratios. We believe this more than triples the number of dwarf stars in the temperature range of our sample for which chemical abundances have been measured using atomic lines.  Our abundance estimates should be more reliable than those of previous studies because we have used updated model atmospheres which include molecular opacity data which is more complete than what was available in any previous models.  The stars we studied have temperatures between 3300 and 4700 K, with a median temperature of 3950 K.  They have [Fe/H] abundances between $-2.4$ and $+0.2$ The abundance ratio [Ti/Fe] decreases with increasing [Fe/H], showing a trend similar to what has been observed for Ti and other $\\alpha$-elements in warmer stars \\citep{w62}.  The previous studies using atomic lines and high resolution spectra to measure abundances in dwarf stars in the temperature range of our stars are sparse and included few stars.  \\citet{m76} studied Kapteyn's Star.  He reported $\\rm [Fe/H] = -0.5 \\pm 0.3$, which is 0.5 dex larger than our result. \\citet{s94} studied six M and K dwarfs using two Fe~{\\sc i} lines.  One star, HD 95735 (Gl 411), was also included in our study.  The Fe abundance we derived for the star is 0.42 or 0.57 dex smaller, depending on which gravity Savanov used.  Several researchers have estimated metallicities of low mass dwarfs by fitting synthetic spectra to low resolution observed spectra, $400 < \\lambda / \\Delta \\lambda < 6000$, which included atomic and molecular lines \\citep{j95,jlah96,viti97,leg02}. The spectra were for the most part in the near infrared.  The metallicities derived were not determined as precisely as is possible using higher resolution spectra, but the method shows what analysis is possible for fainter stars where low resolution spectra are all that is available.  The trend we see in [Ti/Fe] vs [Fe/H] is similar to that seen for warmer stars in the Galaxy. Figure~\\ref{F2} compares our data to the abundances found by \\citet{f00}, who observed field halo and disk stars, most of which are dwarf stars between 5000 and 6500 K.  We have made the correction necessary to account for the different assumed solar abundances. Our abundance values do not appear out of place among the larger set, although our average [Ti/Fe] is higher by about 0.1 dex at most [Fe/H] values. We see a similar relation between our data and that compiled by \\citet{gcc03}. The similar [Ti/Fe] trend was expected since stars of all masses are presumably made from clouds with the same compositions.  Because we selected stars for observation to cover a wide range of metallicity, rather than to statistically represent the low mass dwarfs in a given volume, our results say nothing about the relative numbers of stars with different metallicities.  One goal of this work will be to find a combination of photometric and low-resolution spectral indices which can be used to determine the metallicity of these low mass dwarfs, and to calibrate this method using our observationally determined metallicities.  Because the molecular band strengths measured by the spectral indices depend on both the chemical composition and the temperature of a star, the method will need to be sensitive to both properties. We plan to obtain low resolution spectra of the stars in our list for which TiO and CaH indices have not been measured.  By compiling a statistically significant sample of observationally calibrated metallicity estimates for low mass dwarfs it will be possible to determine the metallicity distribution of these stars in the Galaxy.  This will provide observational evidence of whether there is a K and M dwarf problem in our Galaxy similar to the G dwarf problem, where low-metallicity stars are more scarce than predicted by Galactic star formation and chemical evolution models. The answer to this question, positive or negative, will provide important constraints for models of the chemical enrichment of the Galaxy."
    },
    "0410/astro-ph0410178_arXiv.txt": {
        "abstract": "We report the near-infrared (IR) identification of the likely counterpart to \\x, a  highly-absorbed Galactic X-ray source recently suspected to belong to the rare class of OB supergiant-neutron star binary systems. Our {\\it JHK$_s$}-band imaging of the field reveals the existence within the X-ray error boxes of a near-IR source consistent with an early-type star lying at $d$$\\sim$7 kpc and suffering $A_V$$\\sim$16 mag of extinction, the latter value being in good agreement with the hydrogen column density derived from a modelling of the X-ray spectrum. Our follow-up, near-IR spectroscopic observations confirm the nature of this candidate and lead to a late O-type supergiant classification, thereby supporting the identification of a new Galactic OB-supergiant X-ray binary. ",
        "introduction": "The class of High-Mass X-Ray Binaries (HMXRBs) is defined as systems made up of a compact object (generally a neutron star) accreting material from the wind of an early-type companion, either an OB supergiant or a Be star (in which case X-ray outbursts occur when the neutron star passes through the circumstellar disk). Although the vast majority of the HMXRBs are identified as Be/X-ray binaries, the few systems with an OB supergiant primary are of particular interest. First, these objects might ultimately lead to the formation of a neutron star-black hole system and are thus important in the context of massive binary star evolution. Second, the disturbance induced by the orbital motion of the collapsed object through the primary stellar outflow can be used to probe the structure and physical properties of radiation-driven winds (e.g., Haberl, White \\& Kallman 1989; Kaper, Hammerschlag-Hensberge \\& van Loon 1993). To date, however, only a handful of OB-supergiant X-ray binary systems have been optically identified (Liu, van Paradijs \\& van den Heuvel 2000). Therefore, there is a need to uncover other members of this important population.  The Galactic X-ray source \\xs was first detected using the {\\it Uhuru} satellite (Forman \\etal 1978) and has been persistently catalogued ever since by several X-ray missions. These observations are summarized by Wen, Remillard \\& Bradt (2000; hereafter WRB) who present a variability study of this source using the All-Sky Monitor (ASM) onboard the {\\it Rossi X-Ray Timing Explorer (RXTE)} satellite. The X-ray power ($<$$L_X$$>$$\\sim$5 $\\times$ 10$^{36}$ ergs s$^{-1}$ in the 5--100 keV energy band for an adopted distance of 7 kpc and $A_V$$\\sim$15 mag) showed a coherent, energy-dependent sinusoidal modulation ($\\pm$30 per cent in the 5--12 keV energy band) over 3 yrs with a period of 4.400$\\pm$0.001 days. This value is best interpreted as the orbital period of a binary system, but seems too long to be associated with a low-mass X-ray binary. The X-ray luminosity is also typical of other HMXRBs, but the smooth and regular nature of the X-ray light curve argues against a transient Be-neutron star binary system. These arguments suggest that the mass-donor star is an OB-type supergiant, while the hard X-ray spectrum supports a neutron star companion. Levine \\etal (2004) recently discussed pointed observations of \\xs carried out with the Proportional Counter Array (PCA) and the High Energy X-ray Timing Experiment (HEXTE) onboard {\\it RXTE}. Their analysis confirmed the modulation of the X-ray flux according to the 4.400-d orbital period, but also revealed the existence of pulsed emission with a period of 605 s. The location of \\xs in the pulse period/orbital period plane clearly indicates that it is a wind-fed accretion system (Corbet 1986). The mass function derived from the Doppler delay curve, along with an estimate of the orbital inclination angle derived from a modelling of the orbital phase-dependence of the X-ray flux ($38^o \\la i \\la 72^o$), led to a mass of the primary in the range 9--31 M$_{\\odot}$ and an upper limit on its size of about 22 R$_{\\odot}$. After assumptions about the outflow dynamics, the wind column density estimates were used to derive a wind mass-loss rate, $\\dot{M}$, exceeding 1.3 $\\times$ 10$^{-6}$ M$_{\\odot}$ yr$^{-1}$. Based on the inferred stellar parameters and high mass-loss rate, they further hypothesized that the mass donor could be a Wolf-Rayet star.  The {\\it Einstein} imaging proportional counter (IPC) has localized this X-ray source at 19$^{\\rm h}$10$^{\\rm m}$46$^{\\rm s}$ and +07$^{\\circ}$36$'$07$''$ (J2000) with an uncertainty of about 50 arcsec. Additional constraints on the position are provided by observations made with the modulation collimator experiment (A-3) onboard the {\\it HEAO 1} satellite (see WRB).\\footnote{Note that the position of the grid of {\\it HEAO 1} A-3 ``diamonds'' in fig.1 of WRB is incorrect (L. W. Wen 2001, private communication). See Fig.1 for the correct position.} This relatively small error box offers promising prospects for seeking the counterpart, although the very high dust obscuration resulting both from the presumably large distance and from the location close to the Galactic plane ($b$$\\sim$--0.8$^{\\circ}$) prevents any follow-up studies in the optical. Here we present near-IR imaging and spectroscopic observations identifying the likely counterpart to this X-ray source as a late O-type supergiant. ",
        "conclusions": "The use of several diagnostic lines in the {\\it HK} windows has led to a robust classification of the likely counterpart to \\xs as a late O-type star, while \\brs in emission strongly points to a supergiant (Hanson \\etal 1996). The proposed O7.5--O9.5 If spectroscopic classification is internally consistent with the near-IR photometric properties and also compatible with the stellar parameters independently derived by Levine \\etal (2004) from X-ray data, provided that the star satisfies log\\,($L_{\\star}$/$L_{\\odot}$) $\\ga$ 5.2. As mentioned previously, the X-ray data do not argue in favour of a Be/X-ray binary system. The lack of strong near-IR emission lines in \\xs strongly supports this conclusion (Everall \\etal 1993). Despite the prohibitively high visual extinction, further advances in our understanding of this system may come from radio or time-resolved near-IR observations, the latter both in the photometric and spectroscopic modes. Further X-ray observations with higher spatial resolution, as well as near-IR spectroscopic data for other possible candidates in the field (notably candidate C; see Figs 1 and 2) are also needed for an unambiguous identification of the counterpart."
    },
    "0410/astro-ph0410334_arXiv.txt": {
        "abstract": "The Friedman equation for the universe with arbitrary curvature $(k = 0, \\pm 1)$, filled with mutually noninteracting pressureless dust, radiation, cosmological constant, and strings is considered. We assume the string domination scenario for the evolution of the latter component. Moreover, we discuss the simplest possibility for the scaling of the string energy density: $\\varrho \\propto R^{-2}$. For such models we write down the explicit solution of the Friedman equation. We realize that corresponding cosmological models do not essentially differ from those without strings. We find an analytic formula for the radial coordinate $\\chi$ of a galaxy with a redshift z and express it in terms of astronomical parameters. This relation is then used for the derivation of the astrophysical formulas such as luminosity distance, angular diameter, and source counts, which may serve for testing the string-dominated universe. It seems that the most sensitive test, at least from the formal point of view, is the formula for the number of galaxies $N(z)$ corresponding to a given value of the redshift. We show that the maximum of $N(z)$ strongly depends on the density of strings, especially if the density is large enough to explain the $\\Omega$ problem. Other tests are also proposed. ",
        "introduction": "One of the consequences of the inflationary scenario is the near flatness of the universe, which means that the energy density is very close to the critical value $\\Omega = \\varrho/\\varrho_c = 1$ (Guth 1981; for observational aspects, see also Loh 1986; Loh and Spillar 1986a,b). However, observation of galaxies gives us a value roughly one order less. A lack of observable mass necessary for approaching the critical density is usually called the dark matter problem. people propose many solutions to this problem (Turner 1987). One of the simplest solutons is the assumption that the cosmological constant does not vanish. Other possibilities are massive neutrinos, axions, or heavy superpartners of the usual particles, for example, photinos. Several years ago, when the theories of grand unification (GUTs) were used for the description of the early universe, one more candidate appeared to solve the mentioned problem - cosmic strings.  In gauge theories with spontaneous symmetry breaking, the phase transition at critical temperature (above which symmetry can be restored) can give rise to the nontrivial vacuum structure of the universe (Zel'dovich {\\it et al.} 1974). These are domain walls, strings, or monopoles, depending on the topology of the manifold of degenerate vacua (Kibble 1976). More complicated topological objects such as walls bounded by strings or monopoles connected by strings can also be formed (Vachaspati and Vilenkin 1984). However, we have good reason to believe that, up to the present time, in our visible part of the universe only strings, or eventually an infinite network of strings with monopoles in vertices, could survive (Vilenkin 1981a; Vachaspati and Vilenkin 1987).  It is well-known that there are two extreme configurations to which a system of strings may evolve: a scaling configuration or a string-dominated universe (SDU). The former possibility has been extensively examined in the context of galaxy formation (Zel'dovich 1980; Vilenkin 1981b; Vilenkin and Shafi 1983; Turok and Brandenberger 1986; Kibble 1986). The latter one is also very attractive because it may give the explanation for the missing mass (Vilenkin 1984a; Turner 1985; Kibble 1986). We shall not discuss both scenarios, referring instead to the literature (Vilenkin 1985; Kibble 1985; Bennet 1986a,b; Scherer and Frieman 1986; Aryal {\\it et al.} 1986). Our considerations are based on the result of Turok and Bhattacharjee (1984), who have shown that, neglecting interaction, the energy density of a network of strings scales as $\\varrho_s \\propto R^{-n}$, where $2\\leq n \\leq 3$ (see also Kibble (1986) and Vachaspati and Vilenkin (1987). The case $n=2$ and the astronomical constraints on SDU have been examined by Gott and Rees (1987). Earlier, Gott (1985) had given exact solution and presented interesting considerations concerning the gravitational field of strings in the context of gravitational lensing (see also Hiscock (1985)).  The purpose of the present paper is to calculate some atrophysical formulas for the homogeneous and isotropic universe with arbitrary curvature $(k=0,\\pm 1)$ filed with pressureless dust, radiation, cosmological constant, and the system of strings. For simplicity, we consider only an extreme case when $n=2$ (it corresponds, for example, to the set of randomly oriented straight strings or to the tangled network of strings which conformally stretches by the expansion). This case is particularly interesting because it allows treatment of all mentioned components of the universe simultaneously in an analytic way. Generally, this is not the case if $2\\leq n \\leq 3$. Although the case under consideration is probably too idealistic, it gives some aspects of observational problems in the universe with strings and in some sense completes the discussion of the string-dominated universe given by Vilenkin, Kibble, and Kardashev (1986), Charlton and Turner (1987), or Gott and Rees (1987).   Ther exist several observational tests of the cosmic-string scenario: gravitational lensing (Vilenkin 1984b), anisotropy in the microwave-background radiation (Kaiser and Stebbins 1984), and gravitational radiation emitted by decaying loops (Hogan and Rees 1984). The tests proposed in the present paper are of a different type and are not relevant for detecting a single string. Our tests could be applied only for the universe in which the network of strings is sufficiently dense in order to influence astrophysical observables such as luminosity distance, angular diameter, or source counts. Moreover, the strings under consideration should be relatively light (Vilenkin 1984a).  We do not want to make a statement pro or con regarding the string scenario. Our task is to propose tests that could be used for the verification of the hypothesis concerning a universe dominated by a network of strings. It seems to us that at least one of them, namely source counts, is very promising for this purpose. We leave consideration of the less trivial case of the string-domination scenario (where $3>n>2$ and the universe behaves very like a matter-dominated one) for the future.  In the next section, we introduce some definitions and we describe the Friedmann equation for the universe with strings in a form that allows use of the methods and results (with some modifications) from our previous papers concerning the case without strings (D\\c{a}browski and Stelmach 1986a,b; 1987a). Next, we give analytic solutions of the Friedman equation in terms of Weierstrass elliptic and nonelliptic functions. In Sec. III, we find the relations between astronomical parameters and derive some astrophysical formulas for the universe with strings. Section IV is devoted to the discussion of the tests for the existence of the network of cosmic strings. We pay much attention to the case where the formulas are given by elementary functions. Because of their simplicity, these cases are of special interest. In Sec. V we summarize the results. ",
        "conclusions": "The aim of the paper was to examine some observational implications of the universe filled with a sufficiently dense network of strings. The averaged energy density of strings was assumed to scale as $\\varrho_s \\sim R^{-2}$. Hence, we assumed the simplest possibility when the system of strings stretches conformally by the expansion and may come to dominate the universe. Besides cosmic strings, the universe contains non-interacting pressureless dust, radiation, and a cosmological constant. The mathematical formalism, which could describe such general models, is based on Weierstrass functions. We have given solutions of the Friedman equation and the principal astrophysical formulas for observable quantities in the models. The first conclusion from the results obtained was the confirmation of the general conviction that the influence of the cosmic strings on the evolution of the universe is relatively small in spite of the sophisticated equation of state for the strings (pressure given by strings is negative). On the basis of the derived formulas, we proposed some tests, which may be used for the verification, if not for the entire string scenario, at least for the string-domination scenario (of course, only in the case $\\varrho_s \\sim R^{-2}$. We showed that the luminosity distance and consequently the apparent-magnitude relations are not sensitive tests if one takes small accuracy of present observational data into account. In the next test, we examined the influence of the cosmic strings on the redshift $z_{min}$, for which the angular size of galaxies takes a minimum value. We noticed that in the string-dominated universe $z_{min}$ is shifted compared to the case without strings, and the value of this shift is of order $\\frac{1}{2}$. We checked also that the small $\\Lambda$ term (reasonable value for our present universe) does not change the result too much. The last test considered in the paper was connected with the source counts. We proved that the number of sources corresponding to a given redshift in a homogeneous, expanding universe possesses a maximum depending strongly on the strings. Strictly speaking, this dependence is not so distinct for a low energy density of strings. Then $z_{max}$ is roughly equal to $2$. However, in string-dominated universe, for example in the case when strings form invisible dark matter giving $\\Omega=1$, $z_{max}$ dramatically increases to $e^2 - 1$ for $\\sigma_{s0}\\to \\frac{1}{2}$, where $e=2.71828\\ldots$. From Fig.\\ref{fig3} it follows that the string-dominated scenario (for example, where $90\\%$ of the total energy is in strings and the rest is in matter) should give $z_{max}=4$. One should also emphasize that $z_{max} \\approx 2$ still allows existence of cosmic strings up to $40\\%$ of the total energy (in our considerations, we assumed that the number of sources in a comoving volume is constant in time). In spite of the strong dependence of $z_{max}$ on the density of strings $\\sigma_{s0}$, practical application of the presented test does not seem to be possible up to now because of large values of redshifts which should be taken into account $(z>2)$. First, the problem of faint galaxies exists and there is a limit to how faint a galaxy one can detect. The second problem is even more severe and is connected with evolutionary effects of galaxies for large $z$. The full formula for this test should involve the appropriate corrections.  In conclusion, we would like to remind the reader that throughout the paper we considered only the simplest system of cosmic strings scaling as $\\varrho_s \\sim R^{-2}$. Discussion of observational implications of other, maybe more realistic, string-domination scenarios, where $\\varrho_s \\sim R^{-n}$ and $n>2$, will be published in the future.      \\appendix \\vspace{0.6cm}"
    },
    "0410/astro-ph0410428_arXiv.txt": {
        "abstract": "In an X-ray binary system, the companion star feeds the compact neutron star with plasma materials via accretions. The spinning neutron star is likely covered with a thin ``magnetized ocean\" and may support {\\it magnetohydrodynamic (MHD) tidal waves}. While modulating the thermal properties of the ocean, MHD tidal waves periodically shake the base of the stellar magnetosphere that traps energetic particles, including radiating relativistic electrons. For a radio pulsar, MHD tidal waves in the stellar surface layer may modulate radio emission processes and leave indelible signatures on timescales different from the spin period. Accretion activities are capable of exciting these waves but may also obstruct or obscure their detections meanwhile. Under fortuitous conditions, MHD tidal waves might be detectable and offer valuable means to probe properties of the underlying neutron star. Similar situations may also occur for a cataclysmic variable -- an accretion binary system that contains a rotating magnetic white dwarf. This Letter presents the theory for MHD tidal waves in the magnetized ocean of a rotating degenerate star and emphasizes their potential diagnostics in X-ray and radio emissions. ",
        "introduction": "For a magnetized spinning degenerate compact star covered with a thin plasma layer referred to as an {\\it ocean}, various magnetohydrodynamic (MHD) tidal waves, excited by relevant processes, may exist. This scenario is plausible for an accretion binary system in which the Roche-overflow from a companion star accretes onto a spinning magnetized compact star. This Letter describes the theory of MHD tidal waves that propagate in a thin magnetized ocean of a spinning compact star. When combined with pertinent dynamic processes and emission mechanisms, such MHD tidal waves may leave indelible signatures of electromagnetic emissions in a wide spectrum ranging from the radio to GeV gamma rays and could provide key diagnostics that might offer clues and hints of otherwise complex phenomena. In general, MHD tidal waves may be excited gravitationally, thermally, electromagnetically, or by impacts of accreting materials. While accretion activities may stir up and sustain these waves, wave signals might be undesirably entangled with or overwhelmed by accretion noises. Once observationally established, these waves may be utilized to probe the structures in the {\\it surface layers} of neutron stars and white dwarfs by extending or adapting the helioseismological techniques (Papaloizou \\& Pringle 1978; van Horn 1980; Saio 1982; Lou 2000, 2001b). Furthermore, damping properties of these waves might provide information of the stellar constituent materials and diffusion processes (Lou 1995). ",
        "conclusions": ""
    },
    "0410/astro-ph0410102_arXiv.txt": {
        "abstract": "Detection of X-ray emission from classical novae in their post-outburst stages provides crucial information about the nova phenomenon. The soft X-ray emission gives a direct insight into the remaining hot nuclear burning shell. A numerical model for steady H-burning white dwarf envelopes has been developed to study the post-outburst phases of classical novae. Properties of the X-ray emission of the post-nova white dwarf are obtained, and the results are compared to ROSAT observations of V1974 Cyg. ",
        "introduction": "White dwarf envelope models with steady hydrogen burning have been calculated for realistic compositions of white dwarf envelopes in classical novae, and their evolution has been approximated as a sequence of steady states. Three chemical compositions for ONe novae have been considered, corresponding to models from \\cite{jh98} with 25\\%, 50\\% and 75\\% mixing between the solar accreted matter and the degenerate core. Results show that, for a given white dwarf mass and composition, the envelope evolves in the HR diagram increasing its effective temperature along a plateau of quasi-constant luminosity, while its mass slowly decreases due to hydrogen burning. When the envelope mass is reduced to the minimum value for steady hydrogen burning, the source turns-off. This occurs just when the maximum effective temperature is reached.  Both the luminosity and the envelope mass during the plateau depend on the white dwarf mass and hydrogen mass fraction. The behaviors of luminosity and envelope mass as a function of white dwarf mass and hydrogen mass fraction can be approximated by the expressions  \\begin{center} \\(L_{\\rm plateau}(L_{\\odot })\\simeq 5.95\\times 10^{4}\\left( M_{\\rm c}/M_{\\odot} -0.536X_{\\rm H}-0.14\\right)\\)  \\(\\log M^{\\rm ONe}_{\\rm plateau}(M_{\\odot })\\simeq 0.42X_{\\rm H}-\\left( M_{\\rm c}/M_{\\odot}-0.13\\right) ^{3}-5.26\\) \\end{center} ",
        "conclusions": ""
    },
    "0410/astro-ph0410272_arXiv.txt": {
        "abstract": "A significant fraction of the accreting black holes powering  high redshift AGN are obscured by large columns of dust and gas. For this reason, luminous type 2 quasars can be efficiently discovered combining hard X--ray and near--infrared observations. We will briefly discuss the most recent results. ",
        "introduction": "The space density of luminous quasars in the distant Universe represents a key observational ingredient to understand the formation and evolution of supermassive black holes (SMBHs). The hard X--ray energy range (above 2 keV) is well suited for this purpose as it provides an unbiased view of the obscured accretion power which is responsible for the majority of the SMBH energy output recorded in the X--ray background (XRB) spectrum.  The capabilities of both {\\it Chandra} and XMM--{\\it Newton} in performing sensitive X--ray surveys have been continuously exploited in the last four years. As a result, large samples of X--ray sources spanning a wide range of fluxes (from a few 10$^{-13}$ down to about 10$^{-16}$ erg cm$^{-2}$ s$^{-1}$) are available for reliable statistical studies on the extragalactic X--ray source population. Despite extensive campaigns of spectroscopic follow--up observations with the largest, ground--based telescopes, the X--ray source classification remains challenging. The most important reason is that the optical counterparts of many X--ray sources are too faint even for 8--10 m class telescopes (Fig.~\\ref{rhx}). The spectroscopic completeness of the deepest {\\it Chandra} surveys (CDFS and CDFN) is of the order of 50--70 \\%. Deep multiband photometry is usually employed to estimate redshifts and, over restricted portions of the sky, a much higher completeness level could be achieved.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=0.6\\textwidth]{rx.ps} \\end{center} \\caption[]{The 2--10 keV flux vs. the R--band magnitude for X--ray selected sources in the Chandra (blue small squares) and HELLAS2XMM (red large open circles)} \\label{rhx} \\end{figure}  A detailed discussion on the optical identification breakdown is beyond the purposes of this paper and can be found elsewhere (\\cite{barger03}, \\cite{szokoly04}, \\cite{fiore03}). The general picture emerging from the spectroscopic observations can be summarized as follow:  \\smallskip \\par\\noindent $\\bullet$ Unobscured broad--line AGN are found up to redshift 5 and constitute the dominant population at relatively bright X--ray fluxes ($>$ 10$^{-14}$ erg cm$^{-2}$ s$^{-1}$) \\cite{bauer04}  \\smallskip \\par\\noindent $\\bullet$ The space density of obscured AGN increases towards fainter fluxes. Clear signatures of X--ray and/or optical obscuration are typically detected (though some mismatch between optical and X--ray absorbing properties is present). The luminosity and redshift distribution of obscured sources is skewed toward lower values and most of them are classified as Seyfert galaxies ($L_X < 10^{44}$ erg s$^{-1}$ and $z < $ 1--1.5) rather than quasars.  \\smallskip \\par\\noindent $\\bullet$ The discovery of relatively high X--ray luminosity AGN hosted in the nuclei of optically normal galaxies. Whether they are heavily absorbed Seyfert galaxies, low--luminosity AGN whose emission is diluted by the host galaxy starlight or something more exotic is still subject of debate \\cite{marcy03}.   \\par The most important findings emerging from the identification process concern the redshift and absorption distribution of X--ray selected sources. The former is peaked at lower redshifts ($z \\simeq$ 1) than expected on the basis of population synthesis models for the XRB \\cite{gilli}. The latter seems to indicate a lack of obscured, luminous sources with respect to model predictions.  On the one hand, the results described above favour a late formation of the XRB (i.e. \\cite{steffen04}) and call for a major revision of the AGN synthesis models which are based on the unified scheme. On the other hand, it has been suggested \\cite{treister} that selection effects against the identification of obscured AGN combined with the small solid angle covered by deep surveys prevent from a reliable comparison with model predictions. ",
        "conclusions": ""
    },
    "0410/astro-ph0410099_arXiv.txt": {
        "abstract": "{We present X-ray spectral and timing analysis of members of the young open cluster IC 2391 observed with the XMM-Newton observatory. We detected  99 X-ray sources by analysing the summed  data obtained from MOS1, MOS2 and pn detectors of the EPIC camera; 24 of them are members, or probable members, of the cluster. Stars of all spectral types have been detected, from the early-types to the late-M dwarfs.  Despite the capability of the instrument to recognize up to 3 thermal components, the X-ray spectra of the G, K and M members of the cluster are well described with two thermal components (at kT$_1 \\sim$ 0.3-0.5 keV and kT$_2 \\sim$ 1.0-1.2 keV respectively) while the X-ray spectra of F members require only a softer 1-T model.  The Kolmogorov-Smirnov test applied to the X-ray photon time series shows that approximately 46\\% of the members of IC 2391 are variable with a confidence level $>$99\\%. The comparison of our data with those obtained with ROSAT/PSPC, nine years earlier, and ROSAT/HRI, seven years earlier, shows that there is no evidence of significant  \\v on these time scales, suggesting  that long-term variations due to activity cycles similar to that on the Sun are not common, if present at all, among these young stars. ",
        "introduction": "Most of our knowledge on the long-term evolution of coronal emission derives from observations of open clusters, which supply large, chemically homogeneous, and well dated samples.  First X-ray studies of open clusters with the {\\em Einstein} Observatory \\citep{Stern81, Cai85, Mice85, Mice88, Mice90} showed that the average level of X-ray luminosity decays with age. ROSAT observations have extended the {\\em Einstein} results both by enlarging the number of  known X-ray emitters in a given cluster and covering clusters over a much  wider age range \\citep[e.g.,][]{Ster95, Gag95, Mice96, Mice99, Stau94, Ran95, Ran96, Pro96, Jeff97, Ra98, Mice00, Fra00,Scio00, Bar02}. The comparison of different open clusters shows that the X-ray emission decreases from younger clusters (like IC 2391, IC 2602 at $\\sim$ 30 Myr) to intermediate age clusters (like $\\alpha$ Per, Pleiades and NGC2516 at $\\sim$ 50 - 100 Myr), to older clusters (like Hyades, Coma, Praesepe at $\\sim$ 500-700 Myr).  There is ample consensus that X-ray luminosity evolves during the stellar lifetime mainly as a result of the spin down due to magnetic braking. The coronal emission is influenced by several stellar properties but the rotation, in particular, plays a fundamental role; since rotation evolves with age so does coronal activity.  Data on open  cluster have shown that the most  active stars, with rotation velocities above $\\sim$ 15-20 km/s, reach a maximum X-ray luminosity such that L$_{{\\rm x}}$/L$_{{\\rm bol}}$ saturates at $\\sim$ 10$^{-3}$, where L$_{\\rm bol}$ is the star's bolometric luminosity \\citep[e.g.][]{Vi84, Piz03}. Saturation lacks a clear  interpretation: it could be the effect of dynamo saturation, or it could correspond to the total filling  of the star's surface by active regions, as originally suggested by \\citet{Vi84} or it could be the result of other, yet unknown, phenomena. Furthermore ROSAT observations  have shown that stars rotating even faster ({\\em  v~sin~i~}$>$ 100  km/s) exhibit a level of X-ray luminosity 3-5 times below the saturation level \\citep{Ra98}. The origin of this phenomenon, named supersaturation by \\citet{Pro96}, is still unexplained and several hypotheses have been formulated. Supersaturation could be the result of an overall decrease of dynamo efficiency at very high rotation rates or the consequence of the redistribution of the radiative output due to the lower stellar \"effective\" gravity because of the enhanced rotation (Randich 1998). Recently, \\citet{Mar03} has found evidence of rotational modulation in a supersaturated star of IC 2391 with a very short photometric period of 0.22 days. The observed modulation implies that the star is not completely covered by active regions. As an explanation, \\citet{Jar04} has suggested that the X-ray emitting corona of rapid rotators suffers a centrifugal stripping of coronal loops, producing open field regions that become dark in X-rays.\\\\ The IC 2391 open cluster is an ideal laboratory for testing activity-age relations in fast rotators, because it is young enough ($\\sim$ 30 Myr) that for solar-type members to have just arrived on the zero-age main-sequence (ZAMS). At a  relatively nearby distance (162 pc), this  southern cluster contains about 172 known members or probable members, ranging in spectral type from B to M, enabling the simultaneous study of the different processes driving X-ray emission from stars of different internal structure. The first X-ray observation of the cluster was performed with the ROSAT/PSPC \\citep{Pa93, Pa96}, followed by a 56 ksec ROSAT/HRI observation \\citep{Si98}. These observations have shown that among the late-type members of the cluster  the spread of X-ray luminosities is as large as a factor $\\sim $20. A similar spread is observed in the rotational periods, implying that the stars arrive on the ZAMS with a wide range both in angular momentum and activity level values. However, the limited sensitivity of the ROSAT observations  has allowed the detection of only a small fraction of dM cluster members. On the other hand some dM stars in IC 2391 have been observed and catalogued only recently.  In this paper we report the results of the analysis of the XMM-Newton observation of IC2391, which covers a field of 30${\\arcmin} \\times 30{\\arcmin}$ in the central region of the cluster. The XMM-Newton/EPIC higher sensitivity and  spectral resolution, combined with its larger spectral band and  more continuous time coverage with  respect  to that of ROSAT,  allow us to improve our knowledge of spectral and timing characteristics of the cluster coronae. Unfortunately the cluster size is $\\sim$ 4 square degrees, therefore our observation covers only a small fraction of the known members.  The structure of the paper is the following:  Sect. 2 describes the  XMM-Newton/EPIC observation  and the data  analysis; Sect. 3 and 4 report  the spectral and time analysis, respectively; Sect. 5 summarizes the main results. ",
        "conclusions": ""
    },
    "0410/astro-ph0410266_arXiv.txt": {
        "abstract": "Shock acceleration is a ubiquitous phenomenon in astrophysical plasmas.  Plasma waves and their associated instabilities (e.g., Buneman, Weibel and other two-stream instabilities) created in collisionless shocks are responsible for particle (electron, positron, and ion) acceleration. Using a 3-D relativistic electromagnetic particle (REMP) code, we have investigated particle acceleration associated with a relativistic  jet front propagating into an ambient plasma. We find small differences in the results for no ambient and modest ambient magnetic fields. Simulations show that the Weibel instability created in the collisionless shock front accelerates jet and ambient particles both perpendicular and parallel to the jet propagation direction. The small scale magnetic field structure generated by the Weibel instability is appropriate to the generation of ``jitter'' radiation from deflected electrons (positrons) as opposed to synchrotron radiation. The jitter radiation resulting from small scale magnetic field structures may be important for understanding the complex time structure and spectral evolution observed in gamma-ray bursts or other astrophysical sources containing relativistic jets and relativistic collisionless shocks. ",
        "introduction": "\\label{int}  \\vspace*{-0.6cm} Nonthermal radiation observed from astrophysical systems containing relativistic jets and shocks, e.g., active galactic nuclei (AGNs), gamma-ray bursts (GRBs), and Galactic microquasar systems usually has power-law emission spectra. In most of these systems, the emission is thought to be generated by accelerated electrons through the synchrotron and/or inverse Compton mechanisms. Radiation from these systems is observed in the radio through the gamma-ray region. Radiation in optical and higher frequencies typically requires particle re-acceleration in order to counter radiative losses.  Particle-in-cell (PIC) simulations can shed light on the physical mechanism of particle acceleration that occurs in the complicated dynamics within relativistic shocks.  Recent PIC simulations using injected relativistic electron-ion jets show that acceleration occurs within the downstream jet, rather than by the scattering of particles back and forth across the shock as in Fermi acceleration (Frederiksen et al.\\ 2003, 2004; Hededal et al. 2004; Nishikawa et al.\\ 2003, 2004, 2005). Silva et al.\\ (2003) have presented simulations of the collision of two inter-penetrating electron-positron plasma shells as a model of an astrophysical collisionless shock. In the electron-positron simulations performed with counter-streaming jets (Silva et al.\\ 2003), shock dynamics involving the propagating jet head (where Fermi acceleration may take place) was not investigated. In general, these independent simulations have confirmed that relativistic jets excite the Weibel instability (Weibel 1959).  The Weibel instability generates current filaments and associated magnetic fields (Medvedev and Loeb 1999), and accelerates electrons (Silva et al.\\ 2003; Frederiksen et al.\\ 2003, 2004; Nishikawa et al.\\ 2003, Hededal et al. 2004).  In this paper we present new simulation results of particle acceleration and magnetic field generation for relativistic electron-positron and electron-ion shocks using 3-D relativistic electromagnetic particle-in-cell (REMP) simulations. In our new simulations, electron-positron and electron-ion relativistic jets are injected into electron-positron and electron-ion plasmas in order to study the dynamics of a relativistic collisionless shock both with and without an initial ambient magnetic field.   \\vspace*{-0.7cm} ",
        "conclusions": "\\vspace*{-0.5cm} We have performed self-consistent, three-dimensional relativistic particle simulations of relativistic electron-positron and electron-ion jets propagating into magnetized and unmagnetized electron-positron and electron-ion ambient plasmas. The main acceleration of electrons takes place in the region behind the shock front. Processes in the relativistic collisionless shock are dominated by structures produced by the Weibel instability.  This instability is excited in the downstream region behind the jet head, where electron density perturbations lead to the formation of current filaments. The nonuniform electric field and magnetic field structures associated with these current filaments decelerate the jet electrons and positrons, while accelerating the ambient electrons and positrons, and accelerating (heating) the jet and ambient electrons and positrons in the transverse direction.  Other simulations with different skin depths and plasma frequencies show that the growth rate and spatial structure of current filaments generated by the Weibel instability scale with the plasma frequency and the skin depth (Nishikawa et al. 2004).  An additional simulation in which an electron-ion jet is injected into a ambient plasma with perpendicular magnetic field shows magnetic reconnection due to the generation of an antiparallel magnetic field generated by bending of jet electron trajectories, consequently jet electrons are subject to strong non-thermal acceleration (Hededal and Nishikawa 2004).  These simulation studies have provided new insights for particle acceleration and magnetic field generation. Further research is required to develop radiation models based on these microscopic processes.  {\\bf Acknowledgments:} K. Nishikawa is a NRC Senior Research Fellow at NASA Marshall Space Flight Center. This research (K.N.) is partially supported by the National Science Foundation awards ATM-0100997, and INT-9981508. P. Hardee acknowledges partial support by a National Space Science and Technology (NSSTC/NASA) award.  The simulations have been performed on ORIGIN 2000 and IBM p690 (Copper) at the National Center for Supercomputing Applications (NCSA) which is supported by the National Science Foundation.  \\vspace*{-1.0cm}"
    },
    "0410/astro-ph0410116_arXiv.txt": {
        "abstract": "We discuss constraints on black hole spin and spin-related astrophysics as derived from X-ray spectroscopy.  After a brief discussion about the robustness with which X-ray spectroscopy can be used to probe strong gravity, we summarize how these techniques can constrain black hole spin.  In particular, we highlight {\\it XMM-Newton} studies of the Seyfert galaxy MCG--6-30-15 and the stellar-mass black hole GX~339--4.  The broad X-ray iron line profile, together with reasonable and general astrophysical assumptions, allow a non-rotating black hole to be rejected in both of these sources.  If we make the stronger assertion of no emission from within the innermost stable circular orbit, the MCG--6-30-15 data constrain the dimensionless spin parameter to be $a>0.93$. Furthermore, these {\\it XMM-Newton} data are already providing evidence for exotic spin-related astrophysics in the central regions of this object.  We conclude with a discussion of the impact that {\\it Constellation-X} will have on the study of strong gravity and black hole spin. ",
        "introduction": "With evidence for the existence of black holes and the dynamical measurement of their mass becoming almost pass\\'e, an increasing focus is being placed on detecting the effects of black hole spin.  Spin truly is a creature of the relativistic Universe, and the observational investigation of spin puts us one step closer to being able to genuinely test strong-field General Relativity (GR).  Even if GR passes all of these tests (which, of course, would be the most ``boring'' possibility), black hole spin gives us crucial insight into how black holes of all masses are born, and may well be an important ingredient in powering some of the most energetic sources in the Universe.  At the current time, the best evidence for the effects of black hole spin come from X-ray observations, both timing and spectroscopy. X-ray variability studies, particularly investigations of quasi-periodic oscillations (QPOs) have produced tantalizing hints that we might be witnessing the effects of black hole spin (Stella, Vietri \\& Morsink 1999; Strohmayer 2001).  However, the lack of any agreed upon theoretical framework for the high-frequency QPOs prevents us from drawing robust conclusions at this time.  For this reason, the most compelling studies of black hole spin have originated from X-ray spectroscopy.  In this contribution, we describe constraints on black hole spin from X-ray spectroscopy.  We will align our discussion around three questions; ``Have we seen the effects of strong gravity at all?'', ``Have we seen the effects of black hole spin?'', and ``Can we probe the exotic astrophysics associated with spinning black holes?''  For the impatient reader, the answers to these questions are ``Yes!'', ``Very probably'', and ``We're maybe just starting to...''.  We conclude by discussing future prospects for probing black hole spin and testing strong-field GR with both X-rays and gravitational waves. ",
        "conclusions": "Current data are already allowing us to probe black hole physics within a few gravitational radii of the event horizon, and may well be giving us the first observational glimpses of physics within the ergosphere.  But this is just the beginning of X-ray astronomy's exploration of strong gravity, not the end of the road.  The enormous throughput of {\\it Constellation-X} will allow us to probe detailed time variability of the iron line.  Dynamical timescale line variability, an easy goal for {\\it Constellation-X}, will allow us to follow non-axisymmetric structures in the disk as they orbit (Armitage \\& Reynolds 2003; also see Iwasawa, Miniutti \\& Fabian [2003] for the first hint of such structure in {\\it XMM-Newton} data).  This gives us a direct probe of an almost Keplerian orbit close into a black hole. Furthermore, line variability on the light crossing time will allow us to probe relativistic reverberation signatures (Reynolds et al. 1999; Young \\& Reynolds 2000), essentially giving us a direct probe of the null geodesics in the space-time.  Together, these variability signatures will allow true tests of strong-field GR.  There is no compelling reason to believe that GR fails on the macroscopic scales probed by either X-ray or gravitational wave studies of astrophysical black holes.  In the event that GR is verified, both X-ray and gravitational wave observations will allow unambiguous measurements of black hole spins.  Gravitational wave observations with {\\it LISA} of a stellar mass black hole spiraling into a $10^6\\,M_\\odot$ black hole (a so-called Extreme Mass Ratio Inspiral; EMRI) will allow precision measurement of the supermassive black hole's spin as well as tests of the no-hair theorem and the Kerr metric.  The event rates of such sources is quite uncertain, however, partially due to the recent evidence for ``anti-hierarchical'' black hole growth (e.g., Marconi et al.  2004) and its implications for the number density of $10^6\\,M_\\odot$ black holes in the cosmic past. X-ray spectroscopy with {\\it Constellation-X} provides a crucial parallel track of study in which we can obtain measurements of black hole spin across the whole mass range of astrophysical black holes (i.e., stellar, intermediate, and supermassive) using spectral features that are already known to exist.  Only then can the demographics and astrophysical relevance of black hole spin truly be gauged.  We thank the conference organizers for a stimulating and highly enjoyable meeting.  We also thank the National Science Foundation (US) for support under grant AST0205990."
    },
    "0410/astro-ph0410320_arXiv.txt": {
        "abstract": "The cataclysmic variable SDSSJ~015543.40+002807.2 is confirmed to be a magnetic system of the AM Herculis type.  With an orbital period of 87.13 minutes, it is also the shortest-period eclipsing Polar known.  Monitoring with {\\it XMM-Newton\\/} finds a high-state light curve dominated by a single X-ray emitting accretion pole located slightly prograde of the secondary star.  The hard X-ray spectrum is typical of radial shocks on magnetic white dwarfs ($kT\\sim10$~keV), and there is evidence for a soft X-ray component consistent with reprocessing from the stellar surface.  The optical circular polarization is weak ($v\\lesssim3\\%$) when the accretion rate is high ($m_V\\sim15.5$), due to optically-thick cyclotron emission and the apparent competition between two accreting poles.  However, in low states ($m_V\\sim18$), the polarization increases smoothly to the blue, reaching 20\\% at 4200~\\AA, and the flux spectrum displays a rich set of thermally-broadened cyclotron harmonics that indicate a polar field of 29~MG.  The phase interval preceding the 6.5~min eclipse depicts the development of P-Cygni components followed by complete absorption reversals in the emission lines.  This phenomenon is not unexpected for a strongly accreting magnetic system viewed through the cool base of the funnel, and high-quality spectroscopy through this interval will likely lead to new insights into the dynamics of magnetic coupling and gas flow onto the white dwarf. ",
        "introduction": "Despite a catalog of nearly 80 systems (e.g., Ritter \\& Kolb 2003), discoveries of new AM Herculis systems, or Polars, continue to provide insights into the phenomenon of accretion onto compact magnetic objects through the exploration of interesting geometries and new regions of parameter space.  For example, we have only recently begun to appreciate the degree that a very strong field ($B \\gtrsim 100$~MG) on the white dwarf primary star reduces the observed circular polarization -- the hallmark of strongly magnetic cataclysmic variables (mCVs) -- by shifting the higher optical-depth, low-harmonic cyclotron emission features into the optical portion of the spectrum (Schmidt et al. 1996; 2001).  Unusually high specific mass-transfer rates ($\\dot m \\gtrsim 10$~g~cm$^{-2}$~s$^{-1}$) can depress the accretion spot below the surrounding photosphere, and thereby also yield low polarization (e.g., Stockman 1988).  By contrast, very low-$\\dot m$ systems ($<10^{-2}$~g~cm$^{-2}$~s$^{-1}$) have begun to appear in deep spectroscopic surveys (Reimers, Hagen, \\& Hopp, 1999; Reimers \\& Hagen 2000; Szkody et al. 2003), where the reduced Doppler broadening and high contrast to underlying sources of emission give rise to striking cyclotron harmonic features that can be mistaken for quasar emission lines. Due to the low optical depth, circular polarization in these features can exceed 70\\%.  Eclipsing geometries are particularly valuable for all binary types because the transition timings and profiles allow the locations and sizes of emission regions to be mapped within the systems.  The accretion binary SDSS~J015543.40+002807.2 (hereafter SDSS~J0155+0028) was discovered by Szkody et al. (2002) in the Early Release Data of the Sloan Digital Sky Survey (SDSS).  The SDSS spectrum was saturated due to its brightness, but follow-up time-resolved spectroscopy with the Apache Point Observatory in November 2000 showed strong \\ion{He}{2} $\\lambda$4686 emission and a large amplitude radial velocity curve with an orbital period of 87 min. The most intriguing aspect of the system was that the Balmer and \\ion{He}{1} (but not the \\ion{He}{2}) emission lines went into absorption for part of the orbit, while the continuum did not change during those phases. In addition, the phasing of the absorption did not correspond to the spectroscopic zero-crossing of the radial velocities, such as would result from an eclipse of the white dwarf by the secondary star. SDSS~J0155+0028 was also weakly detected in {\\it ROSAT\\/} survey observations. Based on these characteristics, the system was proposed to be a likely high inclination mCV.  If found to be strongly polarized, SDSS~J0155+0028 would be the shortest-period eclipsing AM Her system known.  Optical photometry obtained when the system was bright ($m_V\\approx15.2$) by students at the Nordic-Baltic Research Course (Dubkova, Kudryavtseva \\& Hirv 2003) in 2002 August supported this picture by revealing a deep (5.9 mag), narrow ($\\sim$400~s) eclipse superposed on a broad hump.  This is characteristic of a magnetic accretion system dominated by a single pole in which the accretion spot is self-eclipsed by the white dwarf for less than half the orbital period.  Further photometry of SDSS~J0155+0028 by Woudt, Warner, \\& Pretorius (2004) on 2003 October 4$-$5 found the system in a low accretion state ($m_V \\sim 18$) but still strongly variable on a period of $87.13\\pm0.02$ min. The brightness hump has a different shape and is narrower in width than in the high state observations.  In this paper we report the results of {\\it XMM-Newton\\/} X-ray and ground-based optical photometry, spectroscopy and spectropolarimetry that confirm SDSS~J0155+0028 to be a magnetic CV. The system is characterized by wide swings in mass-transfer rate and an unusual viewing geometry that offers new avenues for studying the structure of accretion columns in Polars. ",
        "conclusions": "With a period of 87.13~min, SDSS~J0155+0028 is the shortest-period eclipsing Polar known.  Observations to date indicate that the system is highly active ($m_V = 15-16$) a large fraction of the time, and reveal an X-ray light curve and spectrum, as well as a low-state optical cyclotron emission spectrum, that are characteristic of active Polars ($kT \\sim 10$~keV; $B\\sim30$~MG). The $>$5~mag deep, $\\sim$6.5~min duration eclipse has yielded an accurate ephemeris, but photometry with higher time resolution and spectrophotometry with better spectral resolution will provide a more detailed understanding of the stellar components, the extent of the funnel, and the accretion spot location(s) on the white dwarf surface.  Particularly interesting is an absorption-line phase that immediately precedes the main eclipse.  This is interpreted as the result of a favorable viewing perspective of the optically-thick accretion flow through the cool base of the funnel, which must lie virtually in the plane of the orbit. Existing observations have yielded a simple description of the gas flow, but future high-quality spectrophotometry through the absorption-line phase promises to be a powerful technique for mapping the trajectory and properties of the accretion flow in this magnetic CV."
    },
    "0410/hep-th0410247_arXiv.txt": {
        "abstract": "We discuss the cosmological evolution of gravitational waves (GWs) after inflation in a brane-world cosmology embedded in five-dimensional anti-de Sitter (AdS${}_5$) bulk spacetime. In a brane-world scenario, the evolution of GWs is affected by the non-standard cosmological expansion and the excitation of the Kaluza-Klein modes (KK-modes), which are significant in the high-energy regime of the universe. We numerically solve the wave equation of GWs in the Poincar\\'e coordinates of the AdS${}_5$ spacetime. Using a plausible initial condition from inflation, we find that, while the behavior of GWs in the bulk is sensitive to the transition time from inflation to the radiation dominated epoch, the amplitude of GWs on the brane is insensitive to this time if the transition occurs early enough before horizon re-entry. As a result, the amplitude of GWs is suppressed by the excitation of KK-modes which escape from the brane into the bulk, and the effect may compensate the enhancement of the GWs by the non-standard cosmological expansion. Based on this, the influence of the high-energy effects on the stochastic background of GWs is discussed. ",
        "introduction": "\\label{sec: introduction} The stochastic background of gravitational waves (SBGW) generated during the accelerating phase of the early universe is one of the most fundamental predictions of the inflationary scenario and this can provide a direct way to probe the very early universe. In particular, information about extra-dimensions is expected to be imprinted on the SBGW. Motivated by recent developments in particle physics, the possibility that our universe is described by brane with three spatial dimensions embedded in a higher-dimensional space has been extensively discussed. According to this scenario, gravity propagates in the extra spatial dimensions, while the standard model particles are confined to the three dimensional brane. At low-energy scales, four-dimensional general relativity is successfully recovered and the extra-dimensional effects should be fairly small. On the other hand, at high-energy scales, the localization of gravity is not always guaranteed and a significant deviation from the standard four-dimensional theory is expected. If this is true, the SBGW can provide a direct probe of the extra-dimensional effects.  In this letter, we focus on the brane-world model proposed by Randall and Sundrum \\cite{RS2}. In this model, a three dimensional brane is embedded in the five dimensional anti-de Sitter (AdS${}_5$) bulk spacetime (see Ref.\\cite{Maartens} for a comprehensive review of this brane-world scenario). The important model parameter is the AdS${}_5$ curvature radius $\\ell$ that determines the scale on which Newton's force law is modified. Table-top experiments on Newton's force law impose the constraint on $\\ell$ that $\\ell<0.1$mm \\cite{Chi}. In the context of cosmology, the evolution of the universe is significantly modified by extra-dimensional effects when the Hubble horizon becomes shorter than the AdS${}_5$ curvature scale, $H \\ell >1$. Thus there is  a critical frequency corresponding to the wavelength of the gravitational waves (GWs) that cross the horizon when $H \\ell =1$. This frequency is given by $f_{\\rm crit}=0.2$mHz$\\,(0.1$mm$/\\ell)^{1/2}$ \\cite{Hogan}. Since the GWs with short wavelengths re-enter the horizon at high energies, it is expected that the SBGW above the critical frequency is crucially affected by the extra-dimensional effects. Theoretically, there are two main effects in the high-energy regime: (i) the cosmological expansion becomes slow compared with that in four-dimensional general relativity, which enhances the amplitude of the SBGW and (ii) the excitation of Kaluza-Klein modes (KK-modes) suppresses the amplitude of the GWs on the brane. An interesting point is that these two effects affect the amplitude of the SBGW in an opposite way. Thus, quantitative calculations of the effects of the KK-modes excitation are essential in predicting the spectrum of SBGW.  There have been many attempts to understand the effect of the KK-mode excitation. One way is to work with Gaussian normal coordinates \\cite{ELMW,Battye1}. Using these coordinates, we have performed a numerical study of the high-energy effects and showed that the excitation of KK-modes dominates over the effect of the non-standard cosmic expansion \\cite{TH}. However, the Gaussian normal coordinates have a coordinate singularity and the analysis is limited to the relatively low-energy universe $(H \\ell <1)$. Recently, Ichiki and Nakamura  have performed a numerical calculation in the high-energy regime $(H \\ell >1)$ using a null coordinate system \\cite{IN1,IN2}. While they reached the same conclusion as our previous study, they took the initial condition so that the perturbation is constant along a null hypersurface. At present, however, it remains unclear whether the initial condition imposed by them is appropriate for the one determined from inflation.   In this letter, we use the Poincar\\'e coordinate system to perform numerical calculations. The main advantage of this coordinate system is that it covers the whole bulk spacetime. Using these coordinates, several analytic and semi-analytic methods ahve been proposed \\cite{Tanaka,KK,Battye2}, which work well in the low-energy regime. Also, the quantum fluctuations in the inflationary epoch have been discussed in the Poincar\\'e coordinate system \\cite{Gor,Koba}. Thus, we can set more plausible initial conditions from the inflationary universe than those in the Gaussian normal coordinates.   We set up the basic equations and prepare the numerical calculation in Section 2. The numerical results are presented in Section 3. We first discuss the validity of the initial conditions. Then we attempt to construct the spectrum of SBGW in the high energy regime of the universe. Finally, Section 4 is devoted to summary and discussion. ",
        "conclusions": "\\label{sec: summary}   We numerically studied the behavior of GWs in the high-energy regime of the universe in the Randall-Sundrum brane-world model. In contrast to the previous study in the GN coordinates, we used the Poincar\\'e coordinate system in which the brane is moving in the AdS${}_5$ background. We numerically investigated the evolution of GWs in the radiation dominated universe.   We have confirmed that the constant mode, $h(z,\\tau)=$const., in the Poincar\\'e coordinates just after inflation is a plausible initial condition from inflation. In the radiation dominated epoch, however, the constant mode does not hold in the bulk. Thus, in principle, we have to start our numerical calculations just after inflation. Fortunately, however, the resultant amplitude of the GWs on the brane becomes insensitive to the initial time $t_0$ of our numerical simulation if we set $t_0$ sufficiently early. In other words, our numerical result is insensitive to the transition time from inflation to the radiation dominated epoch if the perturbation re-enters the horizon long enough after the transition.   Based on these discussions,  we turn to focus on the frequency dependence of the high-energy effects in order to predict the spectrum of SBGW. We found that the excitation of KK-modes compensates the effect of non-standard cosmological expansion. As a result, the density parameter of SBGW scales as $\\omegagw\\,\\propto\\,f^{0}$ above the critical frequency $f_{\\text{crit}}=0.2 \\text{mHz}\\,(0.1\\text{mm}/\\ell)^{1/2}$. The result seems indistinguishable from the prediction in the standard four-dimensional theory and slightly contradicts with the result by Ichiki and Nakamura \\cite{IN2}. The most striking result of our work is that the behavior of the GWs in the bulk sensitively depends on the transition time. This feature implies that the constant initial condition in the Poincar\\'e coordinate does not agree with the constant initial condition on the null hypersurface even on super-horizon scales. In this sense, the result obtained by Ichiki and Nakamura \\cite{IN2}  may not necessarily agree with our results.  In addition, the present numerical analysis is restricted to the frequency  range $0.3\\lesssim f/f_{\\rm crit}\\lesssim 15$. It is thus premature to discuss the detectability of the SBGW by extrapolating our numerical results  to the high-frequency bands observed via future GW detectors. Nevertheless, the significance of KK-mode excitations above the critical frequency holds the clue to probe the signature of extra-dimensions and/or to constrain on brane-world models. In this respect, a more quantitative and precise prediction for the spectrum of SBGW should deserve a further investigation.  Recently, Kobayashi and Tanaka \\cite{KT} analytically estimated the effect of KK-mode excitation at low-energy scales. According to their result, the amplitude of the GWs on the brane depends on the transition time as well as the energy scale $\\epsilon_*$. This would not be a contradiction with the present numerical result because we set the initial condition in the high-energy regime $H \\ell >1$, where the low-energy approximation used by the authors does not work. Thus, in order to understand our result, we need an analytical study of the KK-mode excitation in the high-energy regime. This is a challenging future work.  \\begin{ack} We would like to thank Roy Maartens for reading our manuscript carefully and for his hospitality during TH's visit to Portsmouth. We would like to thank Kiyotomo Ichiki for discussions and comments. KK and TH acknowledge the support from the Japan Society for Promotion of Science(JSPS) Research Fellowships. AT is supported by a Grant-in-Aid for Scientific Research from the JSPS(No.14740157). \\end{ack}"
    },
    "0410/astro-ph0410305.txt": {
        "abstract": "We present results from a high resolution X-ray imaging study of nearby LINERs observed by ACIS on board {\\it Chandra}.  This study complements and extends previous X-ray studies of LINERs, focusing on the under-explored population of nearby dust-enshrouded infrared-bright LINERs.  The sample consists of 15 IR-bright LINERs (L$_{\\rm FIR}$/L$_{\\rm B}$ $>$ 3), with distances that range from 11 to 26 Mpc.  Combining our sample with previous {\\it Chandra} studies we find that $\\sim$ 51\\% (28/55) of the LINERs display compact hard X-ray cores. The nuclear 2-10 keV luminosities of the galaxies in this expanded sample range from $\\sim$ 2 $\\times$ 10$^{38}$ ergs s$^{-1}$ to $\\sim$ 2 $\\times$ 10$^{44}$ ergs s$^{-1}$.  We find that the most extreme IR-faint LINERs are exclusively AGN.  The fraction of LINERs containing AGN appears to decrease with IR-brightness and increase again at the highest values of L$_{\\rm FIR}$/L$_{\\rm B}$.  We find that of the 24  LINERs showing compact nuclear hard X-ray cores in the expanded sample which were observed at H$\\alpha$ wavelengths, only 8 actually show evidence of a broad line.  Similarly, of the 14 LINERs showing compact nuclear hard X-ray cores with corresponding radio observations, only 8 display a compact flat spectrum radio core.  These findings emphasize the need for high-resolution X-ray imaging observations in the study of IR-bright LINERs.  Finally, we find an intriguing trend in the Eddington ratio vs. L$_{\\rm FIR}$ and L$_{\\rm FIR}$/L$_{\\rm B}$ for the AGN-LINERs in the expanded sample that extends over seven orders of magnitude in L/L$_{\\rm Edd}$.  This correlation may imply a link between black hole growth, as measured by the Eddington ratio, and the star formation rate (SFR), as measured by the far-IR luminosity and IR-brightness ratio.  If the far-IR luminosity is an indicator of the molecular gas content in our sample of LINERs, our results may further indicate that the mass accretion rate scales with the host galaxy's fuel supply.  We discuss the potential implications of our results in the framework of black hole growth and AGN fueling in low luminosity AGN. ",
        "introduction": "One third to as many as one half of all galaxies in the local Universe are classified as Low Ionization Nuclear Emission Line Regions (LINERs), displaying narrow optical emission lines of low ionization uncharacteristic of photoionization from normal stars (e.g., Heckman 1980; Ho, Filippenko, \\& Sargent 1997a; Stauffer 1982).  Since the low excitation optical line spectrum can be produced by a range of processes including shock heating (e.g. Dopita et al. 1996) through cloud-cloud collisions in galaxy mergers or surrounding active galactic nuclei (AGN), cooling flows, and photoionization by nonthermal energy sources (e.g. Ferland \\& Netzer 1983) or hot stars (e.g. Terlevich \\& Melnick 1985), the nature of the central engine in LINERs has been under debate for the past several decades. Because LINERs dominate the population of extragalactic objects in the nearby Universe, determining the nature of their central engine has considerable bearing on a number of issues of astrophysical interest.  Indeed, since normal galaxies are now known ubiquitously to host quiescent black holes (Gebhart et al. 2000; Ferrarese \\& Merritt 2000), the low mass accretion rates inferred for many confirmed accretion-powered LINERs may suggest that these objects capture the population of AGN just before accretion onto the black hole ``turns off\".  As a consequence, establishing the number of LINERs containing accreting black holes (``AGN-LINERs\"), their luminosities, accretion rates, and the relationship of these quantities with the properties of the parent galaxy will provide critical insight into some of the most fundamental questions in extragalactic research today.   The vast majority of LINER galaxies are dust-enshrouded, emitting the bulk of their radiation at infrared (IR) wavelengths.  Using the compilation by Carillo et al. (1999) of all known LINERs studied in the literature, approximately 80\\% of all nearby LINERs have far-IR luminosities exceeding that in the B-band\\footnote[1]{Far-infrared luminosities (in units of solar luminosities: L$\\odot$) correspond to the 40-500$\\mu $m wavelength interval and were calculated using the IRAS 60 and 100 $\\mu $m fluxes according to the prescription: L$_{\\rm FIR}$=1.26$\\times$10$^{-14}$(2.58f$_{60}$+f$_{100}$) in W m$^{-2}$ (Sanders \\& Mirabel 1996), L$_{B}$: B magnitude see Carrillo et al (1999)}. However, almost all detailed multiwavelength investigations of nearby LINERs have been based on optically-selected samples, which generally exclude the IR-bright population (e.g. radio: Nagar et al. 2002, X-ray: Terashima et al. 1998, 2003, Ho et al. 2001; optical: Ho et al.1997b).  While considerable progress has been made on understanding these sources, very little is known about the more numerous IR-bright population. Any study of the AGN detection rate and properties of accretion powered LINERs based on optically-selected samples will therefore be biased toward the IR-faint sources.  Furthermore, IR-bright LINERs, characterized by dusty host galaxies with higher star formation rates, may possibly capture a distinct subset or evolutionary state of the LINER population.   By expanding the sample of LINERs studied to include a larger range of IR-brightness ratios, an understanding of possible evolutionary sequences as well as the physical significance of IR-brightness in the LINER class can for the first time be explored.   To investigate the properties of the IR-bright population, we undertook a {\\it Chandra} ``snapshot\" survey of a sample of nearby IR-bright LINERs.  Traditional optical or ultraviolet studies of IR-bright LINERs can be inconclusive, since highly obscured AGN in dusty galaxies can be hidden at these wavelengths. Sensitive hard X-ray observations at high spatial resolution with {\\it Chandra} can greatly aid in determining accurately the number of LINERs containing accreting black holes, providing a robust probe of obscured AGN out to column densities of a few times 10$^{24}$ cm$^{-2}$.  Detection of a single compact hard X-ray source coincident with the nucleus in the most nearby galaxies would be the telltale sign of an accreting central black hole.  In this paper, we build on our previous archival {\\it Chandra} study of LINERs (Satyapal, Sambruna, \\& Dudik 2004; henceforth SSD04) as well as previous {\\it Chandra} studies of LINERs (e.g. Ho et al. 2001, henceforth H01; Eracleous et al. 2002;) in order to increase the sample of IR-bright LINERs and obtain a more comprehensive understanding of this important class of objects.   The outline of this paper is the following.  In Section 2, we summarize the sample selection criteria for the new observations presented in this work and describe the basic properties of the LINER sample used in our analysis.  In Section 3, we present our observations and data reduction techniques, together with a tabulation of X-ray luminosities and morphological descriptions.  In Section 4, we discuss the analysis and assumptions adopted in deriving black hole masses, bolometric luminosities, and Eddington ratios. In Section 5 we present our results, including a discussion of our derived AGN detection rates, a comparison with other AGN indicators, X-ray morphological findings, luminosities, and correlations between various quantities.  In Section 6 we discuss the implications of our correlations and summarize our conclusions in Section 7. ",
        "conclusions": "We have examined the properties of a sample of nearby IR-bright LINERs using proprietary and archival X-ray observations from {\\it Chandra}.  Our main results are as follows.  \\begin{enumerate} \\item Twenty-eight out of 55 LINERs (51\\%) show compact hard nuclear cores coincident with the radio or {\\it 2MASS} nucleus, with a luminosity L$_X$ (2-10keV) $\\geq$ 2 $\\times$ 10$^{38}$ ergs s$^{-1}$.  The nuclear 2-10 keV luminosity for this expanded sample of galaxies ranges from $\\sim$ 2 $\\times$ 10$^{38}$ ergs s$^{-1}$ to $\\sim$ 2 $\\times$ 10$^{44}$ ergs s$^{-1}$.   \\item We find that the most IR-faint LINERs are exclusively AGN.  The fraction of LINERs containing AGN appears to decrease with IR-brightness and increase again at the highest values of L$_{\\rm FIR}$/L$_{\\rm B}$.  \\item Of the LINERs with hard X-ray nuclear cores and observations at H$\\alpha$ wavelengths, only 33\\% (8/24) show broad lines.  Similarly, only roughly half of the LINERs observed in the radio (8/14) show a flat spectrum radio core.  These findings emphasize the need for high-resolution X-ray imaging observations in the study of IR bright LINERs.   \\item We find a surprising trend in the Eddington ratio as a function of L$_{\\rm FIR}$/L$_{\\rm B}$ and L$_{\\rm FIR}$ that extends over seven orders of magnitude in L/L$_{\\rm Edd}$. This may imply a fundamental link between mass accretion rate, as measured by the Eddington ratio, and star formation rate (SFR), as measured by the IR-brightness.  The correlation may further indicate that the accretion rate scales with fuel supply, a finding that has important implications for AGN fueling in low luminosity AGN. \\end{enumerate}   We are very thankful to Davide Donato, Alex Rinn \\& Brian O'Halloran for their invaluable help in the data analysis, and for numerous instructive discussions and to Joe Weingartner, for his enlightening and thoughtful comments.  We are also very grateful for the insightful suggestions from the referee, which helped improve this work.  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.  SS gratefully acknowledges financial support from NASA grant NAG5-11432 and NAG03-4134X.  RMS gratefully acknowledges financial support from NASA LTSA grant NAG5-10708 and from the Clare Boothe Luce Program of the Henry Luce Foundation."
    },
    "0410/astro-ph0410446_arXiv.txt": {
        "abstract": "Based on observations of the most abundant sulfur-bearing molecules (H$_2$S, CS, NS, SO, H$_2$CS, OCS, and SO$_2$) carried out with the IRAM 30m telescope \\footnote{Based on observations carried out with the IRAM 30m telescope. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).} and SEST \\footnote{Based on observations with the Swedish/ESO submillimetre telescope (SEST) at the European Southern Observatory, Chile. ESO N$^{\\rm o}$70.B-0022.}, we present the first analysis of the sulfur chemistry of an extragalactic source, the nuclear region of the starburst galaxy NGC\\,253. This is the first time that H$_2$S and, tentatively, H$_2$CS are detected towards the nucleus of a starburst galaxy. Source averaged fractional abundances of these molecules are a few 10$^{-9}$, except for CS and OCS which are more abundant (10$^{-8}$). Sulfur isotopic ratios, $^{32}$S/$^{34}$S$\\sim$8$\\pm 2$ and $^{34}$S/$^{33}$S$>$9, are measured through observations of $^{13}$CS, C$^{34}$S, and C$^{33}$S. A comparison with the observed relative abundances towards different prototypical Galactic sources suggests that the chemical composition of NGC\\,253 is similar to that found towards the molecular clouds complexes like Sgr\\,B2 in the nuclear region of the Milky Way. The large overabundance of OCS compared to the predictions of time-dependent sulfur chemistry models supports the idea that OCS is likely injected into the gas phase from the grain mantles by low velocity shocks. ",
        "introduction": "The formation and evolution of stars strongly affects the chemistry of their molecular environment. In particular, sulfur-bearing molecules show enhanced abundances in the hot molecular cores, the condensations associated with massive protostars. The most accepted idea is that during the collapse of a star forming cloud, a large fraction of the available sulfur freezes out by accretion onto grain mantles and hydrogenates in the form of H$_2$S. The high temperatures and turbulent motions in these objects produced by the newly formed stars induce the evaporation and disruption of grain mantles, releasing the H$_2$S into the gas phase. After evaporation, this molecule is believed to drive a fast high temperature gas-phase chemistry leading first to the formation of SO and SO$_2$ and subsequently of other sulfur-bearing molecules such as CS and H$_2$CS \\citep{Charnley97,Hatchell98}. It is still uncertain whether the release of H$_2$S is caused exclusively by thermal evaporation or whether shocks play a role in the disruption of the grain mantles \\citep{Hatchell02}. In low mass star-forming regions, where the temperatures are not high enough to evaporate the grain mantles, there is evidence that shocks may be responsible for injecting the H$_2$S from ices into the gas-phase \\citep{Buckle03}. Sulfur-bearing molecules are considered to be an important tool for studying the presence of shocks within massive star-forming regions. Time-dependent chemical models have roughly succeeded to reproduce the observed abundances of the most common sulfur-bearing species both in high and low mass star-forming regions \\citep{Hatchell98,Buckle03}.  In starburst galaxies like NGC\\,253, only the most abundant sulfur-bearing molecules have been detected so far, namely CS, OCS, SO \\citep{Henkel85,Mauers95,Petu92}, and recently also SO$_2$ and NS \\citep{Martin03}. Extragalactic H$_2$S has only been detected toward the Large Magellanic Cloud, where H$_2$CS is also tentatively detected \\citep{Heikkila99}. In this paper we present the first detection of H$_2$S and the tentative detection of H$_2$CS in NGC\\,253, allowing us to investigate for the first time the sulfur chemistry in its nuclear region using a set of the most abundant sulfur-bearing species. ",
        "conclusions": "\\subsection{Molecular fractional abundances} We find that most of the sulfur-bearing species detected towards the nuclear region of NGC\\,253 show similar fractional abundances of a few 10$^{-9}$ (Table~\\ref{tab:MolecDensity}). Only CS and OCS have abundances a factor of $\\sim5$ higher. In order to understand the mechanism that dominates the chemistry in the nuclear environment of NGC\\,253 we summarize in Table~\\ref{tab:comparison} the fractional abundances of all the sulfur-bearing molecules detected towards this starburst galaxy and those observed towards selected Galactic sources (hot cores, Galactic center clouds, dark clouds, photodissociation regions (PDRs) and shocks driven by massive protostars) which are dominated by different kinds of chemistry. The abundance ratios discussed on this section are also shown in Table~\\ref{tab:comparison}.  The SO$_2$ abundance in PDRs is much lower,by a factor of 40, than that found in NGC\\,253. Thus, abundance ratios such as H$_2$S/SO$_2$ and CS/SO$_2$ can be used to establish the role of photodissociation in the chemistry of NGC\\,253. Particularly instructive is the SO/SO$_2$ abundance ratio. SO/SO$_2$ $\\sim100$ towards PDRs is about two orders of magnitude larger than that in NGC\\,253 which suggests that photodissociation does not play a major role in the chemistry of its nuclear region.  A comparison with dark clouds shows a clearly lower abundance of OCS than that observed towards NGC\\,253. The OCS/SO ratio is 0.1 and 0.4 in L134N and TMC-1 respectively, while we measure $\\sim 4$ in NGC\\,253. The abundance ratios of SO$_2$/SO and H$_2$S/SO of $\\sim 0.2$ and $\\sim 0.05$ respectively measured towards these dark clouds are also much smaller than the corresponding ratios of $\\sim 1$ observed towards NGC\\,253. Therefore, dark molecular clouds like L134N and TMC-1 do not seem to dominate the molecular composition near the nucleus of NGC\\,253.  Chemistry driven by high velocity shocks in molecular outflows also differs significantly from the chemistry found in NGC\\,253. The OCS/SO ratio measured towards these sources is two orders of magnitude smaller than that in NGC\\,253, while CS/H$_2$S is three orders of magnitude lower. The large abundance of H$_2$S, as well as the low abundance of molecules such as CS, OCS, and H$_2$CS in local molecular outflows, is not observed in NGC\\,253.  As far as the hot cores are concerned, the molecular abundances vary significantly between Galactic sources. Different evolutionary states may account for these observed differences. As seen in Table~\\ref{tab:comparison} the abundance ratios in the Orion hot core mostly disagree with those in NGC\\,253. The relative abundances of most of the species towards SGR\\,B2(N) are similar to what we find in NGC\\,253 except for the underabundance of H$_2$S and OCS with respect to the other molecules. The ratio OCS/SO is almost two orders of magnitude lower in both hot cores than in NGC\\,253. In addition, the low rotational temperatures ($T_{\\rm rot}<25$K) derived for all the sulfur-bearing species in NGC\\,253 suggest that the observed emission cannot be due to hot cores where much higher rotational temperatures are expected ($T_{\\rm kin}>70$K).  \\citet{Martin03} pointed out that the molecular abundances of SO$_2$, NO, and NS are similar to those found in the envelope of the SGR\\,B2 molecular complex. Table~\\ref{tab:comparison} also includes the comparison with two different clouds in the envelope of SGR\\,B2. The first position, labeled SGR\\,B2(OH), is located $30''$ south of SGR\\,B2\\,M, and the second one $100''$ north-east of SGR\\,B2\\,M. SO$_2$, SO, and H$_2$CS abundances in the Galactic center clouds agree well with those observed towards NGC\\,253. OCS/SO abundances ratios of 5 and 9 are similar to the ratio of $\\sim 4$ found in NGC\\,253.  Our study of the sulfur chemistry towards the nuclear region of NGC\\,253 seems to indicate a chemistry similar to that of the envelope of SGR\\,B2 where low velocity shocks are thought to be the main heating mechanism of the molecular material \\citep{Flower,Pintado97,Pintado01}. The low rotational temperatures derived for NGC\\,253 are also in agreement with those derived by \\citet{Cummins} in the SGR\\,B2 envelope. These results support the idea of large scale low velocity shocks driving the chemistry of the inner molecular material in the nucleus of NGC\\,253 which is consistent with the high rotational temperatures derived from NH$_3$ and the large observed SiO abundances \\citep{Burillo00,Mauers03}.   \\subsection{Chemical models} In order to get an insight into the dominant chemistry around the central region of NGC\\,253, we compare our results with those available from chemical models \\citep{Millar97,Charnley97,Hatchell98,Buckle03}. For this we will assume that the whole nuclear region of NGC\\,253 can be described as a single molecular cloud. This is not likely as we are observing a 200\\,pc region where a number of unresolved molecular complexes with different physical and chemical properties coexist. Nevertheless, the dominant observed chemistry in the nuclear region of NGC\\,253 could be characterized by a single giant cloud as observed for the envelope of SGR\\,B2 which has a size of $\\sim30$\\,pc. The SiO mapping of the Galactic center region show molecular cloud complexes extended over even larger scales \\citep{Pintado97}.  We have compared the observed molecular abundances in NGC\\,253 with the time-dependent models used by \\citet{Hatchell98} and  \\citet{Buckle03} to describe the chemistry of hot cores and low mass protostars respectively. These models assume the release of H$_2$S from grain mantles as the main precursor of the sulfur chemistry. Models with $T_{\\rm kin}\\leq 50\\,\\rm K$ and n(H$_2$)$<10^5{\\rm cm}^{-3}$ seem to best reproduce the abundances measured for SO$_2$, SO, H$_2$S, and H$_2$CS shown in Table~\\ref{tab:comparison}. This result is in agreement with the idea that the emission from the nuclear region of NGC\\,253 is dominated by low excitation temperature and moderately dense molecular clouds similar to those found in the Galactic center clouds.  However current chemical models fall short of accounting for the observed abundances of OCS by more than one order of magnitude. It has been pointed out by \\citet{Hatchell98} and \\citet{Vandertak03} that OCS may play an important role in the chemistry of sulfur-bearing molecules. Near-IR observations of massive protostars \\citep{Geballe95,Tielens89,Palumbo95} show that sulfur may freeze out onto grains not only in the form of H$_2$S but also in the form of OCS. In fact OCS is the only sulfur bearing molecule detected so far in icy grain mantles. This would imply that OCS is also sputtered from grains by shocks. \\citet{Hatchell98} revised the chemical model by including an initial large abundance of OCS as well as reducing that of H$_2$S, assuming these species are initially injected into the gas phase. These models, depending on the initially injected H$_2$S and OCS \\citep{Wakelam04}, may account for the high OCS abundance observed in the nuclear region of NGC\\,253.  \\subsection{Sulfur isotopic ratios} The $^{32}$S/$^{34}$S abundance ratios of $8\\pm2$ and $13.5\\pm2.5$ observed towards the nuclear starburst environment of NGC\\,253 (Sect.~\\ref{sect:results}) and NGC\\,4945 \\citep{Wang04} respectively are similar to those found within the Galactic inner 3\\,kpc region, but considerably lower than the value of $\\sim 24$ measured in the Galactic disk \\citep{Chin}. The estimated lower limit of the $^{34}$S/$^{33}$S ratio towards NGC\\,253 is higher than the abundance of $\\sim6$ measured towards most of the Galactic sources studied by \\citet{Chin}.  Massive stars, as well as type Ib/c and II supernovae, appear to slightly overproduce  $^{34}$S and underproduce $^{33}$S compared to $^{32}$S \\citep[see][and references therein]{Chin}. Thus, the compact stellar clusters observed near the center of NGC\\,253 \\citep{Watson96} and the estimated overall star formation rate of 3.6\\,M$_\\odot$yr$^{-1}$ \\citep{Strickland04} and supernova rate (SNr) of 0.05--0.3\\,yr$^{-1}$ \\citep{Mattila01,Ulvestad97} might account for the low $^{32}$S/$^{34}$S as well as the enhanced $^{34}$S/$^{33}$S ratios in NGC\\,253. This fact is supported by the high massive star formation rate of $\\sim 0.1$\\,M$_\\odot$yr$^{-1}$ estimated by \\citet{Forbes93} for the inner $6\\arcsec$ of this starburst.  Uncertainties are large enough that the $^{32}$S/$^{34}$S ratio observed towards NGC\\,253 may not significantly differ from that in NGC\\,4945, where slightly higher star formation and supernova rates \\citep{Strickland04} might also have induced an overproduction of $^{34}$S in its nuclear region."
    },
    "0410/astro-ph0410670_arXiv.txt": {
        "abstract": "The evolution of the cluster mass function and the cluster correlation function from $z=0$ to $z\\approx 3$ are determined using $\\sim$$10^6$ clusters obtained from high-resolution simulations of the current best-fit $\\Lambda$CDM cosmology ($\\Omega_m=0.27$, $\\sigma_8=0.84$, $h$ $=0.7$). The results provide predictions for comparisons with future observations of high redshift clusters. A comparison of the predicted mass function of low redshift clusters with observations from early Sloan Digital Sky Survey data, and the predicted abundance of massive distant clusters with observational results, favor a slightly larger amplitude of mass fluctuations ($\\sigma_8 \\sim 0.9$) and lower density parameter ($\\Omega_m \\sim 0.2$); these values are consistent within 1-$\\sigma$ with the current observational and model uncertainties. The cluster correlation function strength increases with redshift for a given mass limit; the clusters were more strongly correlated in the past, due to their increasing bias with redshift--- the bias reaches $b\\sim 100$ at $z$=2 for $M>5\\times 10^{13} h^{-1}M_{\\odot}$ clusters. The richness-dependent cluster correlation function, represented by the correlation scale versus cluster mean separation relation, $R_0 - d$, is generally consistent with observations. This relation can be approximated as $R_0=1.7d^{0.6} h^{-1}$ Mpc for $d \\sim 20 - 60$ $h^{-1}$ Mpc. The $R_0 - d$ relation exhibits surprisingly little evolution with redshift for $z < 2$; this can provide a new test of the current $\\Lambda$CDM model when compared with future observations of high redshift clusters. ",
        "introduction": "Clusters of galaxies, the most massive virialized structures in the universe, provide vital information about large-scale structure of the universe and place powerful constraints on cosmology (Bahcall 1988; Peebles 1993; Carlberg et al. 1997; Rosati, Borgani \\& Norman 2002; Henry 2004; and references therein). The abundance of clusters as a function of mass (i.e. the cluster mass function) and the evolution of this abundance with redshift are sensitive probes of both the present day density parameter ($\\Omega_m$) and the amplitude of mass fluctuations ($\\sigma_8$); this provides a powerful test of the cosmological model (Peebles, Daly, \\& Juszkiewicz 1989; Henry \\& Arnaud 1991; Bahcall \\& Cen 1992; Oukbir \\& Blanchard 1992; Barlett \\& Silk 1993; White \\textit{et al.} 1993; Viana \\& Liddle 1996; Eke, Cole, \\& Frenk 1996; Pen 1996; Cen 1998; Ikebe \\textit{et al.} 2002; Seljak 2002; Bahcall \\textit{et al.} 2003b; and references therein).  The spatial distribution of clusters of galaxies serves as a complementary test of the cosmological model;  this cluster distribution is often described by the two-point cluster correlation function.  The amplitude of the correlation function offers a strong test of the cosmology (Bardeen, Bond, \\& Efstathiou 1987; Bahcall \\& Cen 1992; Mann, Heavens \\& Peacock 1993; Holtzman \\& Primack 1993; Croft \\& Efstathiou 1994; Borgani \\textit{et al.} 1995).  In fact, it was the unexpectedly strong observed cluster correlations (Bahcall \\& Soneira 1983) that provided the first evidence against the then standard $\\Omega_m = 1$ SCDM models (Bahcall \\& Cen 1992; Croft \\textit{et al.} 1997; Borgani, Plionis, \\& Kolokotronis 1999; Governato \\textit{et al.} 1999; Colberg \\textit{et al.} 2000; and references therein).  In addition, the evolution of the cluster correlation function with redshift, which has received comparatively less attention in the literature, is sensitive to the cosmology.  Taken together, the cluster mass and correlation functions provide two of the most powerful constraints on cosmological models.  Predictions have become increasingly robust with larger and higher resolution cosmological simulations, made possible by recent growth in computing power combined with more sophisticated algorithms. Considerable progress has also been made on the observational front to determine cosmological parameters.  Combining the recent WMAP data with finer-scale CMB experiments plus galaxy and Ly-$\\alpha$ forest data, Spergel \\textit{et al.} (2003) determined a best-fit power law $\\Lambda$CDM model.  While further data will refine this model, the differences are likely to be small. As a further check of this model, we present in this work the simulated mass and correlation functions of clusters of galaxies, and their evolution with redshift, determined from mock sky survey cluster catalogs generated from a $\\Lambda$CDM simulation of the current best-fit cosmological model (Spergel \\textit{et al.} 2003).  We compare our results to the most recent observations, and lay the groundwork for comparison with future cluster observations at both low and high redshift.  Such comparisons will provide important tests of the current cosmology, and will enable further improvements in the determination of cosmological parameters. ",
        "conclusions": "We use light cone outputs from a large-scale simulation of the currently best-fit $\\Lambda$CDM cosmological model to generate mock sky survey cluster catalogs that can be readily compared with observations.  The catalogs were used to determine the present day cluster mass and correlation functions, which together constitute a sensitive test of the cosmological model.  We determine the evolution of the cluster mass function from $z=0$ to $z=3$, and the evolution of the cluster correlation function from $z=0$ to $z=2$.  These results provide predictions of the current cosmological model for comparison with future observations of high redshift clusters.  Such comparisons can be used to test the current model and provide new and independent constraints on both the cosmological density parameter, $\\Omega_m$, and the amplitude of mass fluctuations, $\\sigma_8$.  The simulated cluster mass function at low redshift ($z$=0.1--0.2) is compared with the early Sloan Digital Sky Survey cluster mass function (SDSS: Bahcall \\textit{et al.} 2003b).  The $\\Lambda$CDM mass function predicts somewhat higher abundances than are observed.   This suggests that either a small bias ($\\sim$20\\%) exists in the observed cluster mass calibration, or a somewhat lower value of the cosmological density parameter is needed ($\\Omega_m \\approx 0.2$).   The results, however, are consistent with the current cosmology within the combined observational uncertainties. The evolution of the most massive clusters is exponentially dependent on $\\sigma_8^2$ (Bahcall \\& Fan 1998), and therefore can be used to break the $\\Omega_m-\\sigma_8$ degeneracy that exists in the low-redshift mass function.  We find that the $\\Lambda$CDM model predicts a considerably lower abundance of distant ($z > 0.5$) massive clusters than observed.  This suggests a normalization of $\\sigma_8\\approx 0.9$.  We determine the cluster correlation function in the $\\Lambda$CDM cosmology as a function of cluster mass threshold, summarized in the $R_0-d$ relation (Figure 7--8).  We find good agreement with observations of optically selected clusters (Figure 9); X-ray selected clusters appear to suggest somewhat increased cluster correlations.  The wide scatter in the observational data at high-$d$, makes it difficult at this time to provide precise constraints on the cosmology.  The data are, however, in good agreement with the current $\\Lambda$CDM model.  We determine the evolution of the richness dependent cluster correlation function in the $\\Lambda$CDM cosmology (Figures $4-8$). For a given mass limit, the correlation strength increases with redshift. Surprisingly, the $R_0-d$ relation shows no significant evolution out to $z < 2$.  This surprising result provides a new, independent test of the current cosmological model when compared with future observations of high redshift clusters."
    },
    "0410/astro-ph0410393_arXiv.txt": {
        "abstract": "I present light curves for two detached eclipsing binary stars in the region of the LMC cluster NGC 1850, which is possibly a young globular cluster still in formation. One, a likely spectral type O star, is a newly detected eclipsing binary in the region of the very young subcluster NGC 1850A. This binary is among a small number of highly massive O-type stars in binary systems found in LMC clusters. These two eclipsing binaries are the first discovered in the well studied NGC 1850, and the O-type star is the first eclisping binary found in NGC 1850A. Light curves for two NGC 1850 region Cepheid variables are also shown.  Discovering two eclipsing binaries in the young globlular-like cluster NGC 1850 is discussed in terms of the importance of the binary fraction to globular cluster evolution. ",
        "introduction": "Binary stars influence globular cluster evolution to a far greater degree than had long been believed \\citep{mey97,hut92}. Once thought not to exist in numbers significant enough to influence the evolution of globular clusters (GCs), binary stars are now seen as providing energy to GC single stars and halting core collapse. Not only has the importance of binary stars increased to theoretical studies, observers are now detecting globular cluster binaries, overturning earlier observational reports that GCs do not contain signficant numbers of primoridial binary stars \\citep{mey97}.  The question that arises is whether the binary fraction in GCs changes with age.  While young globular clusters have been studied theoretically, it is well known that Milky Way GCs were all formed in the first few Gyr of the history of the universe, but such is apparently not the case in the Magellanic Clouds, with LMC globular-like clusters (GLCs, also called ``young, globular-like stellar associations'') such as NGC 1850 and NGC 1818, which have ages on the orders of $10^7$ or even $10^6$ years.  Preliminary observational work has led to the discovery of some binaries in the field near these clusters \\citep{seb95}, but surveys are needed in order to compare the binary fraction of these ``young'' GLCs with that of the better studied but older Milky Way GCs.  Observations are needed to address the central question of how binaries affect the dynamical evolution of {\\it young} globular-like clusters. The detection reported here of two eclipsing binary stars in the region of the young GLC NGC 1850 does not yet answer this question, but shows that the use of moderate sized (0.9 m) telescopes can be brought to bear on constraining the answer.   I present new time series imaging of the LMC cluster NGC 1850 and an analysis of photometric variables using image subtraction methods. I present four light curves for NGC 1850 region systems. First, I discuss the finding of a second detached eclipsing binary in the region of NGC 1850 (refered to here as ``Variable 5'' or ``V5''); which is the first eclipsing binary found in the region of the subcluster NGC 1850A.  Next, I present an eclipsing binary (``Variable 1'' or ``V1'') in the larger region of NGC 1850 which was found by the Optical Gravitational Lens Experiment (OGLE) collaboration \\citep{uda03}, and its light curve is presented here. The OGLE group has graciously confirmed that their data show V5 to be an eclipsing binary that was originally missed in their large, comprehensive catalog of MC eclipsing binaries \\citep{uda99,uda03} (publicly available from the OGLE project website\\footnote{http://sirius.astrouw.edu.pl/\\textasciitilde ogle/ }), and data from the much longer-term OGLE project provided a better determination of the period than the current dataset alone (A. Udalski, private communication). ",
        "conclusions": ""
    },
    "0410/astro-ph0410500_arXiv.txt": {
        "abstract": "We analyse the projected radial distribution of satellites around bright primary galaxies in the 2dFGRS. We have considered several primary-satellite subsamples to search for dependences of the satellite number density profiles, $\\rho(r_p)$, on properties of satellites and primaries. We find significant differences of the behaviour of $\\rho(r_p)$ depending on primary characteristics. In star-forming primaries, the satellite number density profile is consistent with power laws within projected distance $20<r_p<500$ kpc. On the other hand, passively star-forming primaries show flat $\\rho(r_p)$ for $20<r_p\\le 70$ kpc, well fitted by generalized King models with a large core radius parameter ($r_c\\sim 68$ kpc). In the external regions of the haloes ($r_p>100$ kpc), the density profiles of all primaries are well described by power laws $\\rho(r_p)\\propto r^\\alpha$, although we notice that for red, early spectral type primaries, the outer slope obtained is steeper ($\\alpha^{red} \\sim 1.12$) than that corresponding to blue, late spectral type primaries ($\\alpha^{blue} \\sim -0.79$). We have tested our results by control samples of galaxies identical to the samples of satellites in apparent magnitude and projected distance to the primary, but with a large relative velocity. This sample of unphysical primary-galaxy pairs shows a flat radial density profile beyond $r_p=20$ kpc indicating that our results are not biased toward a decrease of the true number of objects due to catalogue selection effects. Our results can be understood in terms of dynamical friction and tidal stripping on satellites in the primary haloes. These processes can effectively transfer energy to the dark matter, flattening the central steep profiles of the satellite distribution in evolved systems. ",
        "introduction": "In hierarchical models of structure formation, galaxies are mainly the result of the assembling of several small subclumps of mass, that merge to each other into larger structures. Nevertheless, there are also several low mass subclumps which have survived this process, and at present time, they have turned into satellite galaxies of the main host, the primary galaxy. Thus, the radial distribution of the surrounding mass, $\\rho(r)$, can provide useful information to test current models of formation and evolution of structure.  Several analytic expressions for mass density profiles have been proposed in the literature. Earlier on 1972, Gunn \\& Gott proposed the secondary infall model to explain the formation of objects by a continuous accretion process. The model assumes an Einstein-de Sitter universe, where pressureless and non-interacting dust particles are accreted onto an isolated massive center. This model predicts density profiles that follow a power-law shape $\\rho(r) \\propto r^{-9/4}$. Numerical simulations of the formation and non linear evolution of galaxies, follow mass particles from initial conditions up to their present day configuration allowing for a detail understanding of halo structures. The Cold Dark Matter (CDM) hierarchical scenario have proved to be in agreement with many cosmological observations, and has become the most widely adopted model for structure formation. Within this CDM scenario, a universal law was proposed by Navarro, Frenk \\& White (1996 and 1997, NFW profile hereafter) from their studies of dark matter haloes. The authors find a simple law that does not depend on either the halo mass nor the spectrum of initial density fluctuations and cosmological model parameters. The NFW profile can be written as $\\rho(r) \\propto c/(x(1+x)^2)$, where $x$ is the distance r normalized to the radius that contains 200 times the critical density of the Universe ($r_{200}$), $x=r/r_{200}$. All the variables in this formula are related to the concentration parameter, $c$, which depends on the mass of the haloes. Massive objects are characterised by lower $c$ values than the low mass systems, so that small haloes are usually more concentrated. Therefore, NFW dark matter haloes have a cuspy inner density profile with $\\rho(r) \\propto r^{-1}$, while in the outskirts, the density falls as $\\rho(r)\\propto r^{-3}$. This steep behaviour of $\\rho(r)$ at small distances has been confirmed by other authors (Moore et al. 1998, Klypin et al. 2001) who find inner slopes of up to $\\rho \\propto r^{-1.5}$.  Observationally, at galactic scales the flat behaviour of external HI rotational curves indicates a linear growth of mass with galactocentric distance, suggesting an isothermal density profile, $\\rho(r) \\propto r^{-2}$. On the other hand, several authors found cores instead of cuspy behaviour at the central regions of galaxy haloes (Salucci \\& Bukert 2000, de Blok et al. 2001, Gentile et al. 2004). CDM models are unlikely to reproduce such constant density profiles obtained from observational data in the center of haloes, so alternative dark matter candidates have been proposed to reconcile theoretical predictions with observations (see for instance Spergel \\& Steinhardt 2000). Nevertheless, there are some attempts to reproduce central galaxy cores in numerical simulations within universe models dominated by cold dark matter. El-Zant et al. (2001) focus their study on the effects of dynamical friction on density profiles of galaxies, showing that it is possible to obtain a final core feature in haloes starting from a NFW cuspy initial density profile. Within the CDM scenario, the authors explore the energy transport from baryonic clumps to the dark matter component, which heats up the dark matter particles and flattens the density profiles of haloes in their internal regions.  The available observational data used to determine density profiles of galactic haloes does not extend much beyond than the optical radius, since HI rotation curves, are well measured up to $r_p\\sim60-100$ kpc in brightest galaxies. At $r_p>100$ kpc, satellite galaxies represent an interesting alternative to constrain density profiles of primary haloes starting from observations. It is not obvious that satellites follow the mass distribution within the parent halos and actually, many simulation results indicates that the density profiles traced by mass particles are steeper than that obtained using subhaloes. Nevertheless, it is important to note that although dark subhaloes are poor mass tracers, galaxies may follow the mass distribution more accurately. Astrophysical processes such as feedback, cooling or tidal disruption may help to explain the differences between the density profiles obtained by subhaloes and particles. Dekel et al. 2002 have proposed that feedback processes are suitable mechanisms for lowering satellite densities. A satellite with smaller density could be easily disrupted before being accreted onto the inner regions of the haloes, preventing the steepening of $\\rho$ at small $r$. Also, Ghigna et al. 2001 using dissipationless high resolution numerical simulations showed that, as a result of tidal disruption and mergers in cluster cores, subhaloes have an anti-biased number density profile with respect to the dark matter in the central regions of the host haloes. Similar results were obtained by Gao et al. (2004). The authors implemented a semi-analytical model in high resolution simulations of cluster haloes finding, particularly in the inner regions, a biased distribution of dark matter subhaloes compared to the dark matter particle distribution. On the other hand, the authors obtain that the radial distribution of galaxies within simulated clusters is closely related to that of the dark matter. In a recent work, Nagai \\& Kravtsov (2004) also conclude that galaxies in clusters are expected to give a less biased tracer of the dark matter compared to dark subhaloes. Using high-resolution numerical simulations of galaxy clusters in $\\Lambda$CDM cosmology and an implementation of a hydrodynamical model for galaxy formation, the authors find a distribution of subhaloes shallower than that of the dark matter particles, in agreement with previous results. However, since tidal disruption is stronger at the outer regions of subhaloes, their baryonic contents are expected to be less affected by tidal effects implying that galaxies are better tracers of the dark matter profile.  Satellite galaxies are used here in an attempt to put constrains on the radial density profiles of primary galaxies. Our work has the aim of studying the projected number density profiles of satellites around isolated primary galaxies up to 500 kpc in the 2dFGRS. We have searched for dependences of density profiles on several primary and satellite properties and we have tested the results using suitable control samples. In section 2 we describe the selection criteria adopted to construct the samples studied, and in section 3 we show the main results obtained. Finally, in Section 4 we discuss and summarise our principal conclusions. ",
        "conclusions": "In the hierarchical scenario, the growth of structures is expected to happen from low mass objects to larger structures. In such a picture, the accretion of satellites onto primaries is a common process that may determine many properties of the large galaxies (Klypin et al. 1999, Moore et al. 1999, Abadi et al. 2003, Moore et al. 2003). With the aim of studying the density profile of isolated galaxies up to 500 kpc of projected distance, we have analysed the projected number density profiles of satellites around bright primaries in the 2dFGRS. Although van den Bosch et al. (2004) using mock galaxy redshift surveys claim that 2dFGRS data is unable to put constrains on the radial distribution of satellites around primary galaxies, in our analysis, we have obtained flat density profiles for control subsamples of objects as close in projection to primaries than real satellites and with the same apparent magnitude distribution but with large relative velocity. This test gives a clear indication that neither image overlap and merging of APM galaxies nor fiber collisions have a significant effect on our analysis of the radial distribution of objects near an isolated primary, and so our results are essentially free of such possible systematics in the range $20<r_p<500$ kpc.  For the total sample we find a weak, but still significant, departure from a power law at small galactocentric projected distances ($20<r_p<50$ kpc). However, the subsample of star-forming primaries is consistent with a pure power law fit in the range $20<r_p<500$ kpc. On the other hand, the satellite radial distribution of poor star-forming primaries is characterized by a flat $\\rho(r_p)$ at small $r_p$ consistent with a power law with a constant density core. Accordingly, generalized King profiles with a large core radius parameter ($r_c \\sim 70$ kpc) provide good fits to the data. We notice, however, that $\\rho(r_p)$ is significantly steeper in the outer regions of these primaries than in blue, late spectral type primaries (we find $\\alpha\\sim -1.12$ for poor star-forming primaries and $\\alpha \\sim -0.79$ for galaxies with higher levels of star formation) as it was earlier suggested by Lorrimer et al. 1994. We have also explored dependences of $\\rho(r_p)$ on satellite properties finding that the number density profiles of faint, blue, or $\\eta>1.1$ satellite subsamples have smaller core radii and are well fitted by power law fits. By contrast, bright, red, or poor star-forming satellite subsamples require large core radii. Nevertheless, we argue that these dependencies of the number density profiles on satellite properties are probably caused by correlations between primary and satellite properties, since the differences between density profiles of satellites of different types are not present when the subsamples are restricted to belong to primaries of a given colour index and spectral type.  Results from numerical simulations can help to interpret our findings. As pointed out in the Introduction, several authors concluded that the cluster density profiles obtained using the radial distribution of subhaloes are flattened due to tidal disruption effects. Since primary-satellites systems can be considered as a small scale version of clusters, tidal forces can destroy some of the satellites and bias the observed $\\rho(r_p)$. As shown by Nagai \\& Kravstov (2004) results, this tidal disruption is expected to be stronger for subhaloes near the center of the host halo, and as consequence, a lower number of substructures is found at small cluster-centric distances. Even taking into account that tidal mass loss mainly affects outer regions of subhaloes, a percentage of baryonic mass could be loss for each satellite orbiting in the potential well of the primary. After several orbits , a satellite may lose a significant fraction of its stars, falling down to the detection limit of the catalogue, or in a more extreme case, it can be completely disrupted by primary drag forces. In such scenario tidal mechanisms could explain the flattening of the inner number density profile. Nevertheless, our results indicate that only galaxies dominated by old stellar populations (large colour index values and low $\\eta$-spectral type parameters) have a core feature in their profiles, while for active star-forming primaries $\\rho(r_p)$ is consistent with a pure power-law. We can interpret this behaviour as an evolutionary effect, since primaries with a low rate of star formation are probably old systems so that satellites orbiting around these galaxies may have been exposed to tidal forces during long time periods.  We notice that it remains unclear the validity of considering satellites galaxies as suitable mass tracers. As Nagai \\& Kravtsov pointed out, there are selection effects that could affect the subhalo and the galaxy radial distribution. Therefore, we can not make conclusive statements on the dark matter distribution around primary galaxies starting from their satellite radial distribution. High resolution numerical simulations with reliable star formation schemes including feedback process may provide consistent models of the distributions of mass and satellites around bright primaries and give a detailed explanation of the observational results."
    },
    "0410/astro-ph0410736_arXiv.txt": {
        "abstract": "\\vskip -4mm  We have used RR Lyrae and Blue HB stars as tracers of the old Galactic halo, in order to study the halo structure and the galactic rotation as a function of height above the plane. Our sample includes 40 RR Lyrae and 80 BHB stars that are about 2 to 15 kpc above the plane, in a roughly 250~deg$^2$ area around the North Galactic Pole (NGP). We use proper motions (derived from the GSC-II database) and radial velocities to determine the rotation of the halo. From the whole sample the motion appears to be significantly more retrograde than the samples in the solar neighborhood, confirming \\citet{maj92} results and our own preliminary results based on 1/3 the present sample \\citep[][]{kin03,spa03}. However, the better statistics has now revealed the likely existence of two components, whose characteristics need an accurate analysis of systematic errors on the proper motions in order to be assessed in detail. ",
        "introduction": "\\vskip -4mm  The kinematics of halo stars in the solar neighborhood ($Z \\le 2$~kpc) is well established: mean heliocentric rotation velocity $\\langle V \\rangle \\sim -220 \\pm 10$~km\\,s$^{-1}$ (from RR Lyrae and K giant stars, cf.\\ \\citealp{mar98,lay96,chi98,dam01}).\\\\ From halo stars observed at larger distances ($Z \\ge 5$~kpc), the kinematics of the high halo is less clear: the mean rotation was found retrograde ($\\langle V\\rangle = -275 \\pm 16$~km\\,s$^{-1}$) by \\citet{maj92} and \\citet{maj96} in SA 57 near the NGP, and confirmed by \\citet{gil02} in other directions.\\\\ However, other estimates calculated from orbital parameters \\citep[][]{car99,chi00}, as well as \\citet{sir04} recent analysis of very distant BHB stars from the SDSS, do not support any significant halo rotation. The present work attempts to resolve this discrepancy; it grew from earlier studies of halo stars in the North Galactic Cap \\citep[][]{kin96} which confirmed the streaming motion (in the W vector) that Majewski found for his subdwarf sample in SA 57.\\\\ So the questions are: {\\it i) Is the high halo in retrograde rotation? ii) Is there any halo substructure that may mimic  retrograde rotation in some directions?} ",
        "conclusions": "\\\\ \\vskip -5mm $\\bullet$  For $Z < 4$~kpc, the mean rotation V is close to that found for the halo in the solar neighborhood (see corresponding panel in Fig.~\\ref{fig:hist}).\\\\ \\vskip -5mm $\\bullet$ For $Z > 4$~kpc, the {\\em mean rotation is significantly more retrograde} - in agreement with Majewski's earlier finding. This seems to be due to a {\\em group of very retrograde stars}, whose relative importance  increases with $Z$ with respect to the dissipative halo component, and is particularly significant  in the range $Z = 4-10$~kpc (see Fig.~\\ref{fig:hist}). Therefore, this is likely to be a {\\em local substructure, maybe associated with an accretion event}, since  \\citet{sir04} analysis of a large number of distant BHB stars spread over the sky  shows little deviation from the solar neighborhood (dissipative halo) value.\\\\ \\vskip -5mm It will be interesting to extend the work to halo stars in Anticentre fields. Possibly this will allow us to detect gradients in the V motion and discover whether this  is a local effect limited to the NGP or part of a larger systematic effect.\\\\ \\vskip -5mm The errors in the GSC-II are small enough to make our results worth consideration. On the other hand, they are large enough to make it clear that {\\em  GAIA proper motions are essential for a detailed knowledge of the Galactic structure}."
    },
    "0410/astro-ph0410071_arXiv.txt": {
        "abstract": "We raise the question whether in the past a disk could have existed in our Galactic Center which has disappeared now. Our model for the interaction of a cool disk and a hot corona above (Liu et al. 2004) allows to estimate an upper limit for the mass that might have been present in a putative accretion disk after a last star forming event, but would now have evaporated by coronal action. ",
        "introduction": "We study the possibility of a cool disk existing in the Galactic Center in the framework of the disk-corona evaporation. Interaction between the hot corona and the cool disk underneath occurs via energy and mass exchange. The hot corona conducts heat downward by electrons. At the bottom the heat is radiated away. An equilibrium is established: If the density in the corona is too low, Bremsstrahlung is inefficient and the thermal conductive flux heats up some of the disk gas leading to mass evaporation from the disk into the corona. The resulting density increase in the corona raises the radiation loss and thereby counteracts further evaporation. If the coronal density is too high, radiative cooling is too strong and gas condenses into the disk. At the final equilibrium density one of the processes works, either evaporation of disk mass to the corona or condensation of mass from the corona into the cool disk. The outcome depends on the mass flow in the corona from outside. ",
        "conclusions": ""
    },
    "0410/astro-ph0410301_arXiv.txt": {
        "abstract": "We determine the possible masses and radii of the progenitors of white dwarfs in binaries from fits to detailed stellar evolution models and use these to reconstruct the mass-transfer phase in which the white dwarf was formed.  We confirm the earlier finding that in the first phase of mass transfer in the binary evolution leading to a close pair of white dwarfs, the standard common-envelope formalism \\rev{(the $\\alpha$-formalism)} equating the energy balance in the system \\rev{(implicitly assuming angular momentum conservation)}, does not work.  \\rev{An algorithm equating} the angular momentum balance \\rev{(implicitly assuming energy conservation)} can explain the observations.  This conclusion is now based on ten observed systems rather than three.  With the \\rev{latter algorithm (the $\\gamma$-algorithm)} the separation does not change much for approximately equal mass binaries. \\rev{Assuming constant efficiency in the standard $\\alpha$-formalism and a constant value of $\\gamma$,} we investigate the effect of both \\rev{methods} on the change in separation in general and conclude that when there is observational evidence for strong shrinkage of the orbit, the \\rev{$\\gamma$-algorithm} also leads to this.  We \\rev{then} extend our analysis to \\emph{all} close binaries with at least one white dwarf component and reconstruct the mass transfer phases that lead to these binaries. \\rev{In this way we find all possible values of the efficiency of the standard $\\alpha$-formalism and of $\\gamma$ that can explain the observed binaries for different progenitor and companion masses. We find that \\emph{all} observations can be explained with a single value of $\\gamma$, making the $\\gamma$-algorithm a useful tool to predict the outcome of common-envelope evolution.} We discuss the consequences of our findings for different binary populations in the Galaxy, including massive binaries, for which the reconstruction method cannot be used. ",
        "introduction": "Now that double white dwarfs are discovered regularly \\citep[e.g.][]{mar95,mar00,ncd+01} it has become more and more clear that most of them have a mass ratio close to unity \\citep*[e.g.][]{mm99,mmm02}. This is contrary to what is expected from standard population synthesis calculations \\citep*[e.g.][]{ity97,han98}. A possible resolution of this issue was investigated by \\citet{nvy+00,nyp+00}. In the first paper the observed masses of three double white dwarfs and the well known core-mass -- radius relation were used to reconstruct the evolution of the binary back to two main-sequence stars. It followed that the first phase of mass transfer could not be described by the standard common-envelope formalism \\rev{(based explicitly on energy balance, assuming angular momentum conservation implicitly)}, nor by stable Roche-lobe overflow. \\rev{Recent calculations using a detailed stellar evolution code have confirmed this conclusion (Van der Sluys et al., in preparation). Stable Roche-lobe overflow leads to final double white dwarfs with a mass ratio larger than one \\citep[e.g.][]{ity97,han98} and the observed masses can only be reached by stars with initial masses between about 2.3 and 3.5 \\msun that fill their Roche-lobes within a very small initial separation interval (in order to start mass transfer in the Hertzsprung gap). For standard population synthesis assumptions \\citep[e.g.][]{nyp03} this interval only accounts for 0.3 per cent of objects forming white dwarfs, so is inconsistent with the observation that about 10 per cent of white dwarfs are close pairs \\citep{mm99}.}  Instead \\citet{nvy+00} proposed an \\rev{empirical} algorithm based \\rev{explicitly} on angular momentum balance \\rev{(implicitly assuming energy conservation) with a single free parameter and concluded that all the observed systems could be explained with the same value of the free parameter.}  The second paper showed that \\rev{using this algorithm} a satisfactory model for the Galactic population of double white dwarfs can be obtained.  Since then, quite a few more double white dwarfs have been discovered. In particular the SPY project \\citep{ncd+01}, a large survey on the ESO Very Large Telescope, to measure radial velocity variations of some thousand white dwarfs in order to detect duplicity has, and will, enlarge the known double white dwarf sample. We therefore repeat the analysis of \\citet{nvy+00}, including the new discoveries (Section~\\ref{reconstruct}). Furthermore, we study the difference between the \\rev{standard and alternative method} in some detail (Section~\\ref{formalisms}).  We then extend the analysis to \\emph{all} binaries with at least one white dwarf component (Section~\\ref{WD_binaries}) and sdB binaries (Section~\\ref{sdB_binaries}) \\rev{in order to determine what the free parameter in the alternative method must be to explain the observations.}  We then continue with a discussion of the consequences of the our results for the different binary populations (Section~\\ref{consequences}) and round off with our conclusions. ",
        "conclusions": "We have used the masses of observed white dwarfs in binaries to estimate the radii of their progenitors on the assumption that the white dwarf masses are good approximations to the core masses of their progenitor giants.  Using these progenitor masses we have reconstructed the parameters of the progenitor binary systems of the observed white dwarf binaries.  These, together with the observed binary parameters, were used to reconstruct the change in orbital separation during the mass transfer phase in which the white dwarf was formed. By comparing this change to the expected change for the \\rev{standard $\\alpha$-formalism}, \\rev{explicitly} based on the energy balance, and an the \\rev{$\\gamma$-algorithm}, \\rev{explicitly} based on the angular momentum balance, we derived the values of the free parameters in these \\rev{methods}.  The main result is that, as was found earlier, for the first phase of mass transfer in the evolution leading to the currently observed double white dwarf systems, the \\rev{standard $\\alpha$-formalism} cannot explain the observations, while the \\rev{$\\gamma$-algorithm} can. For all the other reconstructed phases either \\rev{method} can explain the observations.  However, the reconstructed values for the \\rev{$\\gamma$-algorithm} strongly cluster around 1.5, while the values of the free parameter in the \\rev{standard $\\alpha$-formalism} (the efficiency parameter), show a wide range of values, skewed towards low ($\\la 0.5$) values. \\rev{Thus the predictive power, at least in statistical sense, of the $\\gamma$-algorithm seems to be greater than the standard $\\alpha$-formalism}."
    },
    "0410/astro-ph0410137_arXiv.txt": {
        "abstract": "The galaxy CFC97 Cen 05 has in the past been considered an H{\\small I}-rich dwarf galaxy in the nearby Centaurus~A group.  We have used Australia Telescope Compact Array observations to show that the H{\\small I} associated with CFC97 Cen 05 by \\citet{CO97} is most likely a Galactic High Velocity Cloud that is centered $\\sim$17$\\arcmin$ from the optical image of the galaxy. At the optical location of the galaxy, which is not that tabulated by \\citet{BA99}, there is no indication of the presence of H{\\small I} for velocities less than the upper limit of the H{\\small I}PASS survey at $\\sim$12,500 kms$^{-1}$.  In addition, WFPC2 images of CFC97 Cen 05 obtained from the {\\it  HST} Archive reveal that this galaxy is in fact not a dwarf at all, but rather is a distant background spiral not associated with the Centaurus~A group. ",
        "introduction": "The current paradigm for the formation of structure in the Universe ($\\Lambda$CDM) envisages that relatively low-mass \\citep[$\\sim10^7\\/ M_\\sun$, e.g.][]{LM02} dark matter halos form first from density fluctuations in the early Universe.  These small-scale clumps of dark matter then merge in a hierarchical process to form larger structures.  This model has revived interest in dwarf galaxies since such systems may well represent the closest present-day analogues to the products of the evolution of the initial dark matter halos.  Dwarf galaxies are ubiquitous, being found in environments as diverse as rich clusters through to the general field.  In many instances, the study of dwarf galaxies and the influence of environment on their evolution is based on observations of dwarfs in relatively nearby groups.  The proximity of such systems allows the investigation of the properties of the dwarfs over a substantial range in luminosity, including systems approaching the absolute magnitude of all but the faintest of the Local Group dwarf population.  Frequently, the initial classification of a system as potential dwarf galaxy member of a group is done via visual inspection of Sky Survey films \\citep[e.g.][]{CO97, HB00}. Such classification as a dwarf candidate is necessarily a subjective process: morphology of the image, surface brightness, extent of resolvable structure, etc, all play a role in the categorization process.  Subsequent follow-up analyses necessarily require the census of dwarfs in a group to be as complete as possible.  Consequently, confirmation of group membership needs to be carried forward for the dwarf candidates.  This can come about via measurement of a radial velocity consistent with group membership, either through optical spectra or H{\\small I} observations, or via a determination of a distance compatible with group membership obtained either directly from, for example, the tip of the red giant branch, or through surface brightness fluctuation techniques \\citep[e.g.][]{HJ00}.  In this paper we present a discussion of the galaxy CFC97 Cen 05.  This galaxy was discovered and classified visually by \\citet{CO97} as one of a number of potential dwarf irregular members of the Centaurus~A group.  This group lies at a distance of $\\sim$3.5 Mpc, and as the name implies, has the unusual E galaxy Cen~A (NGC~5128) as the dominant group member. \\citet{CO97} sought to confirm group membership of their candidates through the measurement of radial velocities obtained with either the Parkes radio-telescope at 21 cm, and/or with the Siding Spring Observatory 2.3 m telescope at H$\\alpha$\\@. Although the heliocentric velocity for CFC97 Cen 05 given by \\citet{CO97}, +122 kms$^{-1}$ obtained from H{\\small I} observations, is more than 200 kms$^{-1}$ lower than the next lowest velocity of the candidates with H{\\small I} detections (ESO~383-G087)\\footnote{Cen 05 is not plotted in the lower panel of Fig.\\ 5 of \\citet{CO97}, as its V$_{LG}$, calculated using the definition employed by \\citet{CO97}, is --166 kms$^{-1}$.}, \\citet{CO97} classified CFC97~Cen 05 as a confirmed dwarf member of the Cen~A group, and listed the total H{\\small I} content as 1.6 $\\times$ 10$^{7}\\/M_{\\sun}$.  Assuming membership in the Cen~A group and using the photometry of \\citet{SC95} corrected for interstellar extinction but not absorption internal to the galaxy, the absolute blue magnitude of CFC97 Cen 05 is M$_{B}$ $\\approx$ --12.0, yielding M$_{H\\small I}$/L$_{B}$ $\\approx$ 1.7 in solar units, a value typical of most dwarf irregulars \\citep{BW04}.  CFC97 Cen 05 is also mentioned in the blind H{\\small I} survey of the Cen A group carried out by \\citet{BA99} with the Parkes telescope and the multibeam focal plane array system \\citep{BD01}.  \\citet{BA99} listed an H{\\small I} detection with a heliocentric velocity of 130 kms$^{-1}$ that corresponded to an  H{\\small I} mass of 1.7 $\\times$ 10$^{7}\\/M_{\\sun}$ assuming membership in the Cen~A group.  These values agreed with those of \\citet{CO97} for CFC97 Cen 05 and consequently, \\citet{BA99} assigned their H{\\small I} detection to this object.  However, \\citet{BA99} apparently failed to notice that the position for the object they claimed as the detection of H{\\small I} in CFC97 Cen 05 is in fact $\\sim$17$\\arcmin$ from the location of the optical image of the galaxy.  As part of a larger project investigating the influence of environment on the evolution of dwarf galaxies, we are carrying out a study of the distribution of the H{\\small I} in a number of dwarf galaxies in nearby groups, including the Cen~A group.  Local Group observations show that in low luminosity dwarfs, the gas can be centered on the optical component, offset from the center of the optical component, or perhaps even completely detached \\citep[e.g.][]{YK97,SG99,RS02}. While the Parkes telescope is adequate in most cases for detecting the presence of H{\\small I}, given the small angular size of the target dwarfs, and the large beam size of that telescope ($\\sim$15$\\arcmin$), higher spatial resolution observations are required to investigate the location of the neutral gas relative to the optical component.  Consequently, we have employed the Australia Telescope Compact Array (ATCA) for our study. One of our targets was CFC97 Cen 05, which lies $\\sim$7$\\arcmin$ ($\\sim$7 kpc in projection) from the luminous spiral NGC~4945 that has a Seyfert~2 nucleus.  Our ATCA observations and subsequent follow-up using HST Archive data have revealed that not only is the H{\\small I} detected by \\citet{CO97} and \\citet{BA99} not associated with CFC97 Cen 05, it is instead most likely a foreground Galactic High Velocity Cloud (HVC), but also the galaxy itself is background spiral and not a dwarf in the Cen~A group.  The remainder of the paper is arranged as follows.  In the following section we present and discuss our ATCA observations of CFC97 Cen 05, while in Sect.\\ 3 we present an optical image of the galaxy obtained from HST Archive WFPC2 images.  Our results are summarized in the final section. ",
        "conclusions": "We have used H{\\small I} observations with the Australia Telescope Compact Array and HST Archive data to clarify the nature of the galaxy CFC97 Cen 05, reputedly a relatively gas-rich dwarf in the nearby Cen~A group.  The H{\\small I} observations reveal that the gas previously associated with this galaxy is instead a Galactic High Velocity Cloud centered some $\\sim$17$\\arcmin$ from the optical image of CFC97 Cen 05.  No gas is detected at the optical location of the galaxy for redshifts less than $\\sim$12,500 kms$^{-1}$. HST WFPC2 observations then reveal that CFC97 Cen 05 is not a dwarf member of the Cen~A group at all, rather it is a distant almost face-on spiral galaxy.  While mistaken identifications of this type are presumably not common, it is a salutory lesson that group membership of apparently low luminosity small galaxies generally needs confirmation via a number of techniques in order to be certain of group membership."
    },
    "0410/astro-ph0410698_arXiv.txt": {
        "abstract": "95\\% of all stars end their lives as white dwarfs. About 20\\% of the hot post-AGB stars are hydrogen deficient. Most of these are the result of a late helium-shell flash, but the evolutionary status of a fraction of about 10--20\\% of the hottest hydrogen-deficient stars, namely four O(He) stars, is as yet unexplained. They could be the long-searched hot successors of RCrB stars, which have not been identified up to now. If this turns out to be true, then a third post-AGB evolutionary sequence is revealed, which is probably the result of a double degenerate merging process. More generally, understanding details of merging double degenerate stars is of interest in the context of SN\\,Ia events and hence cosmology. ",
        "introduction": "Quantitative spectral analyses of hot hydrogen-deficient post-AGB stars performed by our groups during the last decade have revealed two distinct evolutionary sequences. Besides the well-known ``usual'' hydrogen-rich sequence which is established by central stars of planetary nebulae from early post-AGB until the hot white dwarf stages, a hydrogen-deficient sequence has been discovered. It is composed of Wolf-Rayet type central stars which evolve into PG1159-stars and finally might evolve into non-DA white dwarfs. Our analyses reveal that the atmospheres of the hydrogen-deficient stars are dominated by carbon, helium, and oxygen (see reviews by Werner 2001, and Koesterke 2001). A typical abundance pattern is He:C:O = 33:50:17 (by mass), which is for example found for the prototype PG1159-035. It has long been argued that these stars are the result of a late helium-shell flash (late thermal pulse, LTP). The occurrence of such a flash (the re-ignition of helium-shell burning in a post-AGB star or white dwarf), has been predicted in earlier investigations by Iben et al\\@. (1983). But only recently we were able to quantitatively explain the observed surface chemistry with evolutionary models (Herwig et al\\@. 1999). Flash induced envelope mixing causes ingestion and burning of hydrogen. At the same time, the star evolves back onto the AGB (``born-again AGB star'') and retraces post-AGB evolution for a second time, but now as a hydrogen-deficient star. Prominent examples for ongoing ``born-again'' events are FG~Sge and Sakurai's object.  During the course of our studies of hydrogen-deficient post-AGB stars we have identified a small group of four extremely hot objects which have almost pure helium line spectra in the optical. These are classified as O(He) stars following M\\'endez (1991). Our analyses indeed find helium dominated atmospheres with trace amounts of CNO elements, if detectable at all (Tab.\\,1). While our evolutionary models can explain the rich diversity of different He/C/O patterns in Wolf-Rayet and PG1159-stars, they never result in such helium-dominated surface abundances. It is therefore natural to speculate on the existence of a third post-AGB evolutionary sequence and its origin. Our extremely hot helium-rich stars could be the long-searched progeny of the RCrB stars, which are relatively cool ($T_\\mathrm{eff}$ around 10\\,000\\,K) stars with helium-dominated atmospheres, too.  \\begin{table}[ht] \\caption[]{Parameters of the four known O(He) stars, determined by our analyses of optical spectra (Rauch et al\\@. 1998). Typical uncertainties are: $T_\\mathrm{eff}$ $\\raisebox{0.20em}{{\\tiny \\hspace{0.2mm}\\mbox{$\\pm$}\\hspace{0.2mm}}}$10\\,\\%, $\\log g$ $\\raisebox{0.20em}{{\\tiny \\hspace{0.2mm}\\mbox{$\\pm$}\\hspace{0.2mm}}}$0.5\\,dex, abundance ratios $\\raisebox{0.20em}{{\\tiny \\hspace{0.2mm}\\mbox{$\\pm$}\\hspace{0.2mm}}}$0.3\\,dex. For comparison, the last two lines give the mean element abundances of the majority RCrB stars and the peculiar RCrB star V854~Cen, respectively (Rao \\& Lambert 1996). The scatter around the mean C, N, O, and Si abundances is $\\raisebox{0.20em}{{\\tiny \\hspace{0.2mm}\\mbox{$\\pm$}\\hspace{0.2mm}}}$0.15, $\\raisebox{0.20em}{{\\tiny \\hspace{0.2mm}\\mbox{$\\pm$}\\hspace{0.2mm}}}$0.21, $\\raisebox{0.20em}{{\\tiny \\hspace{0.2mm}\\mbox{$\\pm$}\\hspace{0.2mm}}}$0.46, $\\raisebox{0.20em}{{\\tiny \\hspace{0.2mm}\\mbox{$\\pm$}\\hspace{0.2mm}}}$0.18\\,dex, respectively.}\\vspace{-2mm} \\label{php} \\begin{center} \\begin{tabular}{lcr@{.}lr@{.}lr@{.}lr@{.}lr@{.}l} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{$T_\\mathrm{eff}$} & \\multicolumn{2}{c}{$\\log g$} & \\multicolumn{2}{c}{H\\,/\\,He} & \\multicolumn{2}{c}{C\\,/\\,He} & \\multicolumn{2}{c}{N\\,/\\,He} & \\multicolumn{2}{c}{O\\,/\\,He}  \\\\ \\cline{5-12} \\noalign{\\smallskip} \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{kK} & \\multicolumn{2}{c}{cgs} & \\multicolumn{8}{c}{number ratio} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} K\\,1-27  & 105 & ~6&5 & \\hbox{}\\hspace{2mm}\\raisebox{0.20em}{{\\tiny \\mbox{$<$}}} 0&2 & \\raisebox{0.20em}{{\\tiny \\mbox{$<$}}} 0&005 & 0&005 & \\multicolumn{2}{c}{} \\\\ LoTr\\,4  & 120 &  5&5 &      0&5 & \\raisebox{0.20em}{{\\tiny \\mbox{$<$}}} 0&004 & 0&001 & \\raisebox{0.20em}{{\\tiny \\mbox{$<$}}} 0&008       \\\\ HS\\,1522+6615 & 140 &  5&5 &      0&1 &      0&003 & \\multicolumn{4}{c}{}           \\\\ HS\\,2209+8229 & 100 &  6&0 & \\hspace{2mm}\\raisebox{0.20em}{{\\tiny \\mbox{$<$}}} 0&2 & \\multicolumn{6}{c}{} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} majority RCrB &          &   \\multicolumn{2}{c}{}  & \\hspace{2mm}\\raisebox{0.20em}{{\\tiny \\mbox{$<$}}} 0&0001&    0&010 & 0&004 & 0&005 \\\\ V854 Cen      &          &   \\multicolumn{2}{c}{}  &      0&5   &    0&030 & 0&0003& 0&003\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table}  For RCrB stars two origins for their hydrogen-deficiency are discussed in the literature since many years. Either they are the result of a late thermal pulse or of a merging process of two white dwarfs (double degenerate scenario). Since we now can rule out the LTP scenario for helium-dominated post-AGB stars, it seems inevitable that both, the four hot helium-rich stars and the RCrB stars are mergers and form another evolutionary post-AGB sequence. We therefore expect similar heavy element abundances. In order to investigate this possible evolutionary link it is necessary to compare CNO abundances of both groups, because detailed theoretical predictions for abundances resulting from a merger event are not available. The element abundances in RCrB stars are well studied (Tab.\\,1). ",
        "conclusions": "FUSE spectra of all O(He) stars have been taken during Cycle 2. These have been analyzed using wind models.  No significant wind features are detected neither in the observation nor in the model spectra (Fig.\\,1). Thus, the mass-loss rates of O(He) stars are not higher than predicted by radiative-driven wind theory and a change of the surface composition due to the stellar wind is unlikely. Unfortunately, the FUSE spectra do not show isolated metal lines and thus, allow to determine upper limits for abundances only (ongoing study).  \\begin{figure}[ht] \\plotone{rauch2_fig1.eps} \\caption{Section of the FUSE spectra (around the O\\,{\\sc vi} $\\lambda\\lambda$ 1031.9, 1037.6\\,\\AA\\ resonance doublet) of HS\\,1522+6615, the central stars of LoTr\\,4 and K\\,1-27, and HS\\,2209+8229 (top to bottom) compared with synthetic spectra which consider mass loss as predicted by Pauldrach et al\\@. (1988) with mass-loss rates of $\\log \\dot{M} / \\mathrm{M_\\odot/yr} = -7.6, -7.7, -9.1, -9.7$, respectively. The emission features found in the spectra are due to airglow, and the absorption features not found in the models are interstellar, predominantly H$_2$ } \\end{figure}"
    },
    "0410/astro-ph0410651_arXiv.txt": {
        "abstract": "In order to resolve the overproduction problem of $^{11}$B in supernova explosions during Galactic chemical evolution, the dependence of the ejected masses of the light elements produced through the $\\nu$-process in supernova explosions on supernova neutrino parameters is investigated and constraints on the supernova neutrinos are evaluated. Detailed nucleosynthesis in a supernova explosion model corresponding to SN 1987A is calculated by postprocessing. The ejected masses of $^{11}$B and $^7$Li depend strongly on the temperature of $\\nu_{\\mu,\\tau}$ and $\\bar{\\nu}_{\\mu,\\tau}$ and are roughly proportional to the total neutrino energy. The range of temperature of $\\nu_{\\mu,\\tau}$ and $\\bar{\\nu}_{\\mu,\\tau}$ appropriate for the amount of $^{11}$B necessary for Galactic chemical evolution and the total neutrino energy deduced from the gravitational energy of a typical neutron star is between 4.8 MeV and 6.6 MeV. In the case of neutrino energy spectra with non-zero chemical potential, this range decreases by about 10 \\%. ",
        "introduction": "Supernova explosions are one of the sites promoted for the production of light elements, Li, Be, and B, in Galactic chemical evolution (GCE), in addition to the Galactic Cosmic Rays (GCRs), AGB stars,and novae \\cite{fo00}. They continuously provide $^7$Li and $^{11}$B through the interactions of nuclei such as $^4$He and $^{12}$C with neutrinos emitted from a proto-neutron star. This synthesis process is called the $\\nu$-process \\cite{wh90}.  Overproduction of $^{11}$B in supernovae is one of the standing problems in supernova nucleosynthesis and the GCE (e.g., \\cite{fo00}). Previous works have indicated that the supernova contribution of $^{11}$B production evaluated from supernova nucleosynthesis models \\cite{ww95} is too large by a factor of 2.5 $\\sim$ 5.6 to reproduce the GCE of the light elements; the factor depends on models of the GCE \\cite{fo00,rl00,al02}. Most supernova nucleosynthesis calculations were carried out with a total neutrino energy of $3 \\times 10^{53}$ ergs and temperature of 8 MeV for $\\nu_{\\mu,\\tau}$ and $\\bar{\\nu}_{\\mu,\\tau}$ and 4 MeV for $\\nu_{\\rm e}$ and $\\bar{\\nu}_{\\rm e}$ \\cite{wh90,ww95}. The supernova explosion mechanism has not been resolved, so the characteristics of the supernova neutrinos, especially their energy spectra, have not been fully determined \\cite{kr03}. Therefore, the light element amounts have not been obtained from a specific neutrino model precisely determined by supernova explosion models. Therefore, we systematically investigate the dependence of the light element synthesis in supernova explosions on the supernova neutrinos. Furthermore, on the basis of the evaluated dependence and the results of the GCE of the light elements, we restrict the characteristics of the supernova neutrinos, such as their energy spectra. ",
        "conclusions": "\\begin{figure}[htb] \\begin{minipage}[t]{76mm} \\includegraphics[width=76mm]{f1.eps} \\vspace*{-8mm} \\caption{ Mass fraction distribution of light elements (Li, Be, and B) in the case of $E_\\nu = 3 \\times 10^{53}$ ergs and $T_\\nu = 6$ MeV. The mass fraction of $^7$Li is the sum of the mass fractions of $^7$Li and $^7$Be. The mass fraction of $^{11}$B is the sum of the mass fractions of $^{11}$B and $^{11}$C. } \\label{fig:f1} \\end{minipage} \\hspace{\\fill} \\begin{minipage}[t]{76mm} \\includegraphics[width=76mm]{f2.eps} \\vspace*{-8mm} \\caption{ The ejected mass of $^{11}$B as a function of $T_\\nu$. Each attached number denotes $E_\\nu$ in units of $1 \\times 10^{53}$ ergs. Two horizontal lines indicate the proper $^{11}$B mass range evaluated from the GCE models \\cite{fo00,rl00,al02}. The shaded region shows $T_\\nu$ satisfying the proper ranges of the $^{11}$B mass and $E_\\nu$. A point denoted by WW95 is the $^{11}$B ejected mass of S20A model \\cite{ww95}. } \\end{minipage} \\end{figure}  Figure 1 shows mass fraction distribution of light elements in the supernova ejecta. The light elements are mainly produced in the He/C layer and smaller amounts are in the inner O rich layer. We have shown that $^7$Li and $^{11}$B are produced through the $\\nu$-process reactions: $^4$He($\\nu,\\nu'p)^3$H, $^4$He($\\nu,\\nu'n)^3$He and $\\alpha$-captures in the He layer and $^{12}$C($\\nu,\\nu'p)^{11}$B in the inner O rich layer \\cite{yt04}. They are partly destroyed by the capture of $\\alpha$-particles.  The ejected mass of $^{11}$B as a function of the temperature of $\\nu_{\\mu,\\tau}$ and $\\bar{\\nu}_{\\mu,\\tau}$, $T_\\nu$, is shown in Fig. 2. The dependence of $T_\\nu$ on the ejected mass is stronger than linear. Since the ejected mass is roughly proportional to $E_\\nu$ \\cite{yt04}, the dependence of $T_\\nu$ is stronger than the $E_\\nu$ dependence. This is because the dependence on the cross sections of the $\\nu$-process is larger than a linear dependence. The $^{11}$B ejected mass in the case of $E_\\nu=3 \\times 10^{53}$ ergs and $T_\\nu$ = 8 MeV is consistent with that in \\cite{ww95} in spite of different values of $T_{\\nu_{\\rm e}}$ and $T_{\\bar{\\nu}_{\\rm e}}$. Since $T_{\\nu_{\\rm e}}$ and $T_{\\bar{\\nu}_{\\rm e}}$ are smaller than $T_\\nu$ and neutral current interactions are important for the $\\nu$-process of the light elements, the difference of $T_{\\nu_{\\rm e}}$ and $T_{\\bar{\\nu}_{\\rm e}}$ from \\cite{ww95} scarcely affects the ejected masses of the light elements \\cite{yt04}.  We will apply the above results to GCE of the light elements. Recent studies of the GCE have indicated that both GCRs and supernovae contribute to the B production: $^{10}$B is produced through the GCRs and $^{11}$B is produced through the GCRs and SNe. In order to reproduce meteoritic $^{11}$B/$^{10}$B ratio (=4.05) at solar metallicity, the supernova contribution of $^{11}$B is important since the $^{11}$B/$^{10}$B ratio of the GCRs is 2.5 \\cite{fo00,wm85}. However, $^{11}$B is overproduced in SN nucleosynthesis models \\cite{ww95} compared to the evaluation of the GCE. Prior evaluations of the GCE of the light elements give the overproduction factor $f_\\nu$ to be 0.18 \\cite{rl00}, 0.28 \\cite{al02}, and 0.40 \\cite{fo00}. Thus, we set the range of the factor appropriate for the $^{11}$B amount in GCE to be $0.18 \\le f_\\nu \\le 0.40$ as in \\cite{yt04}. The corresponding range of the $^{11}$B ejected mass is $3.3 \\times 10^{-7} M_\\odot \\le M({\\rm ^{11}B}) \\le 7.4 \\times 10^{-7} M_\\odot$ (two horizontal lines of Fig. 2). We also restrict $E_\\nu$ to be the gravitational energy of a typical neutron star (see \\S 2) \\cite{lp01}. Finally, we obtain a shaded region satisfying both of the conditions of the $^{11}$B mass and $E_\\nu$ (see Fig. 2). The obtained range of $T_\\nu$ is $4.8 \\, {\\rm MeV} \\, \\le \\, T_\\nu \\, \\le \\, 6.6 \\, {\\rm MeV}$, which is smaller than 8 MeV in \\cite{ww95}.   Our evaluation is carried out only in the case of SN 1987A, which corresponds to about 20 $M_\\odot$ star in ZAMS stage. For $^{11}$B production, the $\\nu$-process reactions of $^4$He in the He layer and $^{12}$C in the O layer are important. It depends on the progenitor mass models which layer is more important. Fortunately, however, the neutrino temperature dependence does not change strongly in either $\\nu$-process reactions. Thus, the appropriate range of the neutrino temperature would be a good approximation for supernovae with different progenitor masses. The investigation of the dependence of the $^{11}$B production on the neutrinos spectra for different progenitor mass is a future subject.   We briefly discuss the effect of non-zero chemical potential of the supernova neutrinos. Detailed discussion is written in \\cite{yk04}. Recent studies on the neutrino transfer during supernova explosions have indicated that the energy spectra of the supernova neutrinos approximately obey ^^ ^^ pinched'' Fermi-Dirac distribution rather than that with zero chemical potential \\cite{kr03}. In such a case, it is expected that the ejected masses of the light elements change since reaction rates of the $\\nu$-process change. Hartmann et al. \\cite{hm99} discussed the effect of non-zero chemical potential assuming that cross sections of the $\\nu$-process are proportional to the square of the neutrino energy. We generalize their discussion: we approximate the cross sections of the $\\nu$-process to the $\\alpha$th power low of the neutrino energy. Then, we evaluate the effect on the reaction rates of the $\\nu$-process. The reaction rates of the $\\nu$-process can be written as a function of the degeneration factor $\\eta_\\nu=\\mu_\\nu/kT_\\nu$ for given $\\alpha$ and the neutrino temperature $T_\\nu$, and increase by a factor of 1.4 $\\sim$ 1.5 in the case of $\\eta_\\nu=3$ and of the $\\alpha$ range $4 \\le \\alpha \\le 7$ compared to that with $\\eta_\\nu=0$. Since the ejected masses of $^{11}$B and $^7$Li are roughly proportional to the $\\nu$-process reaction rates, they increase with the same factor. In this effect, the neutrino temperature range appropriate for the $^{11}$B amount in GCE should decrease in the case of non-zero chemical potential. In the case of $\\eta_\\nu=3$, this range decreases by 10\\% at most."
    },
    "0410/astro-ph0410467_arXiv.txt": {
        "abstract": "Imaging data from the Sloan Digital Sky Survey is used to measure the empirical size-richness relation for a large sample of galaxy clusters.  Using population subtraction methods, we determine the radius at which the cluster galaxy number density is \\twoOM\\ times the mean galaxy density, without assuming a model for the radial distribution of galaxies in clusters. If these galaxies are unbiased on Mpc scales, this galaxy-density-based \\rtwo\\ reflects the characteristic radii of clusters. We measure the scaling of this characteristic radius with richness over an order of magnitude in cluster richness, from rich clusters to poor groups.  We use this information to examine the radial profiles of galaxies in clusters as a function of cluster richness, finding that the concentration of the galaxy distribution decreases with richness and is systematically lower than the concentrations measured for dark matter profiles in N-body simulations.  Using these scaled radii, we investigate the behavior of the cluster luminosity function, and find that it is well matched by a Schechter function for galaxies brighter than $M_r=-18$ only after the central galaxy has been removed.  We find that the luminosity function varies with richness and with distance from the cluster center, underscoring the importance of using an aperture that scales with cluster mass to compare physically equivalent regions of these different systems. We note that the lowest richness systems in our catalog have properties consistent with those expected of the earliest-forming halos; our cluster-finding algorithm, in addition to reliably finding clusters, may be efficient at finding fossil groups. ",
        "introduction": "Galaxy groups and clusters appear in many guises. Observationally, they can be identified as pools of X-ray emitting gas, collections of galaxies, Sunyaev-Zel'dovich decrements of the CMB, or strong features in the gravitational shear field.  Theoretically, they are identified as the largest over-dense `halos' of dark matter. Mass is the key defining attribute of a galaxy cluster. The evolution of the cluster mass function and its variance plays an important role in constraining cosmological parameters describing dark energy, such as $w$\\ and $\\Omega_{\\Lambda}$\\ \\citep[e.g.][]{Lim05}, and large-scale structure, such as $\\Omega_m$\\ and $\\sigma_8$ \\citep[e.g.][]{Bah03b}. Mass estimates and studies of cluster members both rely upon knowing the size of a cluster of given mass. To define the mass of a cluster, typically a cluster radius is specified through some prescription, and the total mass taken to be the mass contained within that radius. Various operational definitions are used to determine cluster size and mass, but all of these show a regular increase in cluster size with mass. In order to study objects in a wide range of masses in a meaningful way, we must be able to determine an appropriate size scale for these clusters and groups; we may then use this characteristic scale as an aperture within which to make comparisons.  In numerical simulations, the precision of cluster mass and size measurements is limited only by resolution.  Still, there are a variety of definitions for both size and mass in use, as discussed in some detail by \\citet{Whi01}. One class of estimates is based on top-hat filtered spherical over-densities. In this model, clusters are expected to be virialized within regions where the enclosed mean mass density exceeds the critical density by a factor $\\Delta \\sim$ 200 \\citep{Pee93, Pea99}. The radius at which this over-density is reached, $R_{\\Delta=200}$, is used as the characteristic radius of the cluster. The total mass within this radius, $M_{\\Delta=200}$, is used as the characteristic mass. A number of choices of $\\Delta$ are in use in the literature, from an overdensity of 180 times the mean background \\citep[e.g.][]{Jen01,Kra04}, 200 times the critical density, i.e. 200$\\Omega_{M}^{-1}$ times the mean background \\citep[e.g.][]{Dia01, Evr02, Koc03}, to the ``virial mass'' \\citep{ECF96, Bul01}. Alternative definitions identify halos by `friends-of-friends' (FOF) methods \\citep{Dav85, Jen01}. In these methods, particles are associated with halos to which they are linked by sequences of neighboring particles. Masses for FOF halos are often given by the sum of member particle masses, but as the halos are not required to be spherical, halo size is less clearly defined.  Observationally, cluster mass and size are difficult to measure directly; typically some mass estimator is adopted as a proxy, and typically a mass model is assumed to calculate the virial size. With deep observations, lensing may be used to make detailed mass maps of rich clusters, but this is not yet a practical technique for a large sample of systems spanning a wide range of masses. With large, shallower surveys, the lensing signals from a set of lens systems can be stacked to determine a composite mass profile \\citep{She01}. However, the large number of lens systems needed for high signal to noise makes such a measurement difficult for examining narrow mass ranges.  While examining the lensing signal from an individual cluster recovers a specific mass, that estimate is affected by the errors, such as those due to projection effects, in modeling the mass. Combining the lensing signal from many clusters is advantageous because projection effects are unimportant, and the stacking simplifies modeling of neighboring structures. While stacking does limit the resolution in mass for a large set of clusters, this limitation is due to the amount of data available rather than the technique used. Other observational techniques for estimating mass, whether using X-ray temperature, SZ flux, or velocity dispersion as a mass estimator, also rely upon either expensive spectroscopy to gain a detailed understanding of a few rich systems, or upon models for the mass distribution used to infer the total system mass and size. To compare observational data to theoretical models, and to compare the observed properties of clusters of different masses, it would be preferable to avoid using a model-dependent mass/radius scaling.  With a large photometric optical survey it is now possible to use cluster richness to characterize galaxy systems without suffering from the projection effects that have plagued such a mass estimator in the past. Since the cluster catalog of \\citet{Abe58}, systems of galaxies have been sorted and compared using a variety of richness parameters, many of which have been based on the number of galaxies within a certain luminosity range and distance from the estimated cluster center.  The richness parameters of \\citet{Abe89, Cou91, Dal92, Lum92, Lid96, Pos96, Ost98, Ols99, Gla02, Gotcat, Pos02, Gal03} are of this type; see \\cite{Bah81, Yee99, Bah03} for further discussion and comparison of some of these richness estimators. We can stack systems in narrow richness bins, and measure directly the distribution of galaxies in clusters over a wide range of masses. This galaxy distribution is used to estimate the virial size of these systems. Ideally, cluster members would be identified spectroscopically.  Such data is not feasible to obtain for a very large sample of clusters, so we rely upon projected photometric data taken in multiple bandpasses, and correct for the foreground and background galaxies that contaminate our line of sight to each cluster.  The Sloan Digital Sky Survey \\citep[SDSS,][]{Yor00,Sto02} data offer thousands of clusters and groups for study, and can be used to measure excellent photometric redshifts for those objects.  We use data from the SDSS to directly determine a size-richness relation for groups and clusters with a model-independent method. Since we cannot directly measure the radius at which the cluster has a mass over-density of ${\\Delta M}$, we instead determine the radius $R_{\\Delta}^{\\cal N}$ at which the space density of cluster galaxies, \\Ndense, is over-dense by ${\\Delta}$\\Ndense. We present the scaling of \\rN\\ with richness, which can be employed to further study galaxy clusters and their members. For example, to study the relationship between different mass estimators, such as the mass from lensing measurements and the total luminosity of the cluster, it is essential to know this scaling of cluster size with richness.  If the distribution of galaxies in a halo traces the overall dark matter distribution, our galaxy-density based \\rN\\ will reflect the characteristic radius of clusters.  In detail, the relation between the dark matter density profile and the radial distribution of a population of galaxies depends on a number of physical processes including dynamical friction and tidal striping, and depends on the properties of the galaxy sample, but both simulations and previous observational work suggest that the distribution of galaxies in a halo at least roughly traces the overall dark matter distribution \\citep[e.g.,][]{NK04, lin04}.  This hypothesis is supported by recent lensing and galaxy clustering measurements \\citep{she04,Wei04}, which suggest that the bias of typical SDSS galaxies is approximately one and is roughly scale-independent on scales larger than a few hundred $h^{-1}$ kpc.  In any case, without assuming a model for the radial distribution of galaxies in clusters, we can directly measure the radius at which the galaxy density in clusters is $\\Delta$ times more dense than the average background.  Our investigation does not distinguish between galaxy groups and clusters.  Systems of galaxies come in a range of masses from single galaxies in larger halos up to the most massive of clusters.  There is a clear boundary between galaxies and systems of galaxies; there is not any clear dividing point between poor and rich systems of galaxies as observed in the optical. We use the full range of our cluster finder to develop our catalog of systems; our mass estimator spans the range from highly populated clusters down to very sparse systems, and we refer to any system with two or more galaxies as a ``cluster''. There is no strong break in the properties of the systems found as a function of richness.  In \\S \\ref{data} we describe the SDSS data used and discuss the cluster finding technique and richness measurement; in \\S \\ref{method} we present and test our method of background subtraction through examination of the radial profile and luminosity function. Our calculation of \\rN\\ as a function of richness is presented in \\S \\ref{r200}. We further examine the radial density profile within \\rN\\ in \\S \\ref{rpr200}; we examine the luminosity function within \\rN\\ and how the luminosity function changes as a function of $r$/\\rN\\ for clusters of different richnesses in \\S \\ref{lfr200}. Throughout, we assume a flat, LCDM cosmology with H$_{0} = 100h$ km s$^{-1}$ Mpc$^{-1}$, $h = 0.7$, and matter density $\\Omega_m = 0.3$. ",
        "conclusions": "\\label{disc} By using statistical background subtraction, we have observationally determined the changes in the radial distribution of cluster galaxies as a function of cluster richness, and used that information to calculate a characteristic radius for clusters of each richness. This model-independent radius, \\rN, is based on the number density of galaxies analogously to the way \\rtwo\\ in use in simulations is based on mass density. We find that \\rN\\ exhibits a power law scaling relationship with cluster richness. This result is in good agreement with the expected scaling from simulations.  Our result is also in general agreement with results obtained by other groups using different methods for measuring richness and \\rtwo. \\citet{Gir95} examined spectroscopically confirmed members of 90 rich clusters, and determined virial and core radii by fitting the observed galaxy distribution with an isothermal sphere. They find virial radii of 0.5 - 1 $h^{-1}$\\ Mpc, the same range we find for \\rN\\ for rich clusters. \\citet{Yee03} examined \\rtwo\\ as a function of cluster richness $B_{gc}$\\ (the amplitude of the galaxy cluster center correlation function measured for each cluster, scaled by a luminosity function and spatial distribution) for 16 rich clusters with spectroscopically confirmed members. They measured \\rtwo\\ by applying a singular isothermal sphere model to the velocity dispersion data \\citep{Car97_r200}, and find \\rtwo\\ $\\sim B_{gc}^{0.54 \\pm 0.18}$. This scaling matches well with what we find, although with greater scatter than our result. With our techniques, we have been able to use photometric data to confirm this relationship for rich clusters and extend the relationship to much less rich systems.  To compare equivalent regions of clusters of different richness, \\rN\\ may be used as an aperture within which to compare the properties of cluster members. The space density distribution of galaxies within \\rN\\ is well described by an NFW profile, and the derived concentration parameter varies with cluster richness. We examined the population distribution function of cluster galaxies to determine how the luminosity function of cluster members changes both radially and with cluster richness, using our determination of \\rN\\ to compare clusters in a wide range of richnesses.  The radial variation of the luminosity function of cluster galaxies is similar in clusters of all richnesses, but does depend on the cluster richness. It is important to note that we can still detect a signal from the clusters at 2$h^{-1}$ Mpc from the cluster, even for poor groups. That there is still a significant over-density at these distances suggests caution when measuring the contribution of the background in an annulus centered on the cluster.  We find that the central galaxies of clusters are distinguishable from the rest of the cluster galaxy population, as has been noted by others in both observational studies and in theoretical models. The BCG population is clearly evidenced in the luminosity function within \\rN, and more dramatically in the LF of the central region of clusters, as a ``bump'' at bright magnitudes that rises above the LF of the other cluster members. As the richness of the cluster increases, the BCG population becomes brighter but contributes less to the overall cluster light, in agreement with the results of $K$-band observations of BCGs in 93 X-ray selected systems studied by \\citet{lin04BCG}. We find that the luminosity function of low- and intermediate-richness systems cannot be well described by a Schechter function when the BCGs are included; this function provides an acceptable fit only for the richer systems where the fractional contribution of the central galaxy to the LF is small.  The total cluster luminosity function within \\rN\\ is well modeled by the sum of a Gaussian for the central galaxy and a Schechter function for the satellites over the whole range of \\Ngals. The trends of mean BCG luminosity and fractional contribution of the BCGs are in good agreement with the models of \\citet{Zhe04}. We find that the mean BCG luminosity scales with mean log \\Ngals\\ as \\Ngals$^{1/2}$, a similar though slightly steeper result to what is seen in other studies.  The low-richness objects identified with our cluster finding algorithm have properties that are quite different from rich clusters. Systems with \\Ngals\\ $<$\\ 5 and \\Ngals\\ $>$\\ 7 have clearly different radial profiles and luminosity functions. The transition region $5 \\leq $\\ \\Ngals\\ $\\leq 7$\\ has intermediate properties. One possible contribution to this difference is that the lowest-\\Ngals\\ objects are also the most likely to suffer from the effects of misidentifying the cluster center and/or the the brightest cluster member. In addition, the \\Ngals\\ $<$ 8 objects have a strong selection function: maxBCG is finding systems where specifically only a few red galaxies are within 1$h^{-1}$ Mpc of the galaxy identified as the BCG. It is interesting that demanding so few galaxies of that color be in the neighborhood has the effect of finding objects that have very concentrated radial profiles, are located in isolated regions, and are dominated by bright galaxies.  A possible explanation for these very low \\Ngals\\ systems is that at least some of them are fossil groups. Observationally, a fossil group has the X-ray luminosity of a group or poor cluster, but in the optical, only a highly luminous early-type galaxy without bright neighbors is observed. The first such object was detected by \\citet{Pon94}; subsequent study has shown that there are a few other galaxies in this group, but all are significantly fainter ($\\Delta \\rm m \\ge 2.5mag$) than the first-ranked galaxy \\citep{Jon00}. Other fossil groups have been identified as well \\citep{Mul99,Vik99,Rom00,Mat01}. Such systems are thought to be the end-product of groups in which most of the galaxies have merged, producing a highly luminous central galaxy mostly alone in a halo of hot X-ray gas. This scenario is supported by the observations that the central galaxies emit a very high fraction of the total optical light of the group, and that there is a dearth of $\\rm L_{*}$\\ galaxies in the central region of these systems \\citep{Jon03}.  At fixed halo mass, there is a theoretical expectation that there should be a correlation between halo formation time and the total number of galaxies above a given luminosity threshold \\citep{Zent05}. This correlation may select a special population of of clusters at low \\Ngals.  We speculate that in high mass halos, all systems have bright red galaxies regardless of their formation time, and are equally well found by the maxBCG algorithm, while in smaller, group-mass halos the requirement to select red galaxies may select only the earliest forming halos.  These low \\Ngals\\ systems would be expected to be early-forming halos with higher than average masses for their \\Ngals, whose outer satellites have already merged with the central object. This theory is consistent with the LF shape and the extremely steep radial profiles seen here, as well as with the higher masses that are indicated by cluster-mass correlation function examined by \\citet{she05}.  Many of these systems may also be classified as compact groups (CGs), such as those studied by \\citealt{Hic82,Hic93} and others. Compact groups have been identified in SDSS data by \\citet{Lee04}. They investigate the local environment of CGs, and find that while on average the number density of surrounding galaxies is comparable to the local environment around field galaxies, there is considerable scatter to both more and less dense environments (see their Figure 9). By demanding few other bright, red galaxies nearby, the maxBCG algorithm may be preferentially selecting fossil groups, which would also look like those CGs in environments with a low number density of neighbors. Further study is needed to understand this interesting set of systems. More investigation can be done using SDSS data, but it also would be interesting to look for an X-ray signal from a stacked set of these low-richness objects.  We draw attention to the range of luminosity functions for galaxies within \\rN\\ of clusters of different richness, and also to the differences in the LF seen radially. Such variation makes it difficult to compare between different cluster catalogs, where different proxies for mass and aperture are in use. The LF of clusters in a mass range are different when different fractions of \\rN\\ are sampled, and different again when a fixed aperture is used for each cluster. Nonetheless, we do see similar results to those of other authors, including an upturn at the faint end for rich clusters. This effect comes primarily from galaxies located in the outer regions of the clusters.  We are currently working to compare our results with simulations to show how our space density-based \\rN\\ relates to the mass density-based \\rtwo\\ commonly used in N-body models. Also, we will examine how our richness parameter, \\Ngals, relates to mass $M_{200}$ from the simulations. Further work will be done to use the rich SDSS dataset to explore the distribution of many properties of cluster galaxies in addition to those considered here.  Using a small subset of the SDSS data is sufficient to determine the scaling of \\rN\\ with cluster richness, and to robustly measure the luminosity functions and radial profiles in small \\Ngals\\ bins. Using the full SDSS dataset will allow us to study $\\sim$ 500 of the richest clusters and over 25,000 groups, and will allow cluster detection to higher redshifts. We will check for evolution in our scaling relationship by examining the \\rN -\\Ngals\\ trend in different redshift samples.  Preliminary results indicate that the scaling is the same, but with higher normalization for lower redshift objects. To compare the properties of any stacked set of clusters, the relationship of \\rN\\ to \\Ngals\\ is essential. We can now compare different mass estimates, including those measured from total luminosity, velocity dispersion, and from lensing information. We will also use lensing studies to compare the profiles of the luminous and dark components of the clusters, more closely investigate the possible differences in concentration of the two profiles, and examine the bias between luminous and dark matter in these dense environments. Extending this work to include color and morphology indicators, we will probe many characteristics of galaxies in clusters, exploring the history of galaxy formation as a function of environment, and studying the BCG population.   For large surveys it is typically not practical to determine cluster mass or size using spectroscopy. Without assuming a model for the distribution of galaxies, we provide a way to determine cluster size from photometric data alone for $z \\le$ 0.3.  With knowledge of how the size-mass scaling evolves and a better understanding of the scaling of \\rN\\ with total mass, our method provides a feasible way to measure the characteristic radii and masses of clusters that will be found in future large, high-redshift surveys."
    },
    "0410/astro-ph0410184_arXiv.txt": {
        "abstract": "\\noindent The determination of the primary cosmic ray mass composition from the characteristics of extensive air showers (EAS), obtained at an observation level in the lower half of the atmosphere, is still an open problem. In this work we propose a new method of Multiparametric Topological Analysis  and show its applicability for the determination of the mass composition of the primary cosmic rays at the PeV energy region. ",
        "introduction": "It is clear that due to large fluctuations in the longitudinal development of extensive air showers (EAS) and relatively weak sensitivity of EAS characteristics to the mass of the parent particle, the observed parameters of showers initiated by particles with different primary masses largely overlap. In order to estimate the primary mass in the case of the individual shower or to estimate the mean mass composition at a certain energy we have to use as many observables as possible. To process the big amount of information various multiparametric and non-parametric methods are employed. Among them are multivariate fitting, KNN-method, Bayesian approach, neural net analysis and others. In \\cite{Ambr1} we proposed a new method of the multiparametric topological analysis (MTA) for the study of the {\\em mean} mass composition on the basis of the EAS longitudinal development data. We tested this method using the data simulated in the EeV energy region for the Pierre Auger experiment and were encouraged by its accuracy and the easy applicability. In this contribution we test this method at the lower, PeV energy region close to the well known 'knee' in the primary cosmic ray energy spectrum ($\\sim$3 PeV). Experimental arrays for the EAS study in this energy region are usually more compact and are able to measure the total size of electron $N_e$, muon $N_\\mu$, hadron $N_h$ and Cherenkov light $N_{ph}$ components of the shower. The typical example of such a complex EAS array is KASCADE \\cite{Anton}. Here we use only two parameters: $N_e$ and $N_\\mu$, although the method is easily generalized for the larger number of observables. ",
        "conclusions": "We proposed and tested the Multiparametric Topological Analysis method for the determination of the {\\em mean} primary cosmic ray mass composition on the basis of measurements of just two parameters of the observed EAS - the total number of electrons and muons. Definitely the method needs further development before being used for the processing of real experimental cascades, but the first results are very encouraging. The next efforts should include the increase of simulation statistics, the study of the effect of the shower energy spectrum and angular distribution, the accurate estimate of the reconstruction errors. The great advantage of the proposed MTA method is that it is extremely easy to use and generalize for the larger number of observables.  \\vspace{5mm}  \\noindent{\\large{\\bf Acknowledgements}} \\medskip  \\noindent  Authors thank K.H.Kampert and L.Perrone for useful discussions."
    },
    "0410/astro-ph0410521_arXiv.txt": {
        "abstract": "We present results of a large scale, multi-epoch optical survey of the Ori OB1 association, carried out with the QuEST camera at the Venezuela National Astronomical Observatory. We identify for the first time the widely spread low-mass, young population in the Orion OB1a and OB1b sub-associations. Candidate members were picked up by their variability in the $V$-band and position in color-magnitude diagrams. We obtained spectra to confirm membership. In a region spanning $\\rm \\sim 68 \\Box^\\circ$ we found 197 new young stars; of these, 56 are located in the Ori OB1a subassociation and 142 in Ori OB1b. The spatial distribution of the the low mass young stars is spatially coincident with that of the high mass members, but suggests a much sharper edge to the association. Comparison with the spatial extent of molecular gas and extinction maps indicates that the subassociation Ori 1b is concentrated within a ring-like structure of radius $\\sim$ 2$^\\circ (\\sim$ 15 pc at 440 pc), centered roughly on the star $\\epsilon$ Ori in the Orion belt. The ring is apparent in $^{13}$CO and corresponds to a region with an extinction $A_V \\ge 1$. The stars exhibiting strong H$\\alpha$ emission, an indicator of active accretion, are found along this ring, while the center is populated with weak H$\\alpha$ emitting stars. In contrast, Ori OB1a is located in a region devoid of gas and dust. We identify a grouping of stars within a $\\rm \\sim 3\\> deg^2$ area located in 1a, roughly clustered around the B2 star 25 Ori. The Herbig Ae/Be star V346 Ori is also associated with this grouping, which could be an older analog of $\\sigma$ Ori. Using using several sets of evolutionary tracks we find an age of 7 - 10 Myr for Ori 1a and of $\\sim 4 - 6$ Myr for Ori OB1b, consistent with previous estimates from OB stars. Indicators such as the equivalent width of H$\\alpha$ and near-IR excesses show that the number of accreting low-mass stars decreases sharply between  Ori 1b and Ori 1a. These results indicate that while a substantial fraction of accreting disks remain at ages $\\sim 5$ Myr, inner disks are essentially dissipated by 10 Myr. ",
        "introduction": "The Orion OB1 Association, located at roughly 400 pc (see review by \\citealt{ges89}) and spanning over $\\rm 200 \\> deg^2$ on the sky, is one of the largest and nearest regions with active star formation. With a wide range of ages and environmental conditions, Orion is an ideal laboratory for investigating fundamental questions related to the birth of stars and planetary systems. \\citet{bla64} identified four major subgroups which differ in age and their content of gas and dust: Ori OB1a, 1b (that roughly encompasses the belt stars $\\delta$, $\\epsilon$ and $\\zeta$ Ori), 1c and 1d (where the Orion Nebula Cluster [ONC] is located). Within this vast region there are extremely young groupings of stars still embedded in their parent clouds, such as NGC 2068, 2071, 2023 and 2024 in the Orion B molecular cloud, observed mostly at IR and sub-mm wavelengths (\\citealt{dal99}; \\citealt{hlp01}; \\citealt{maw01}); the young, dense, and partially embedded ONC \\citep[hereafter H97]{het86,ali95,hil97}; clusterings of intermediate age ($\\sim 3-5$ Myr) stars including the $\\sigma$ Ori \\citep{gar67,wwf97} and $\\lambda$ Ori groups \\citep{mup77,dom99,dom01,dom02} and older populations like the $\\sim 10$ Myr old OB1a sub-association, located in an area essentially devoid of gas and dust.  While the earliest stages of stellar evolution must be probed with infrared and radio techniques, many fundamental questions including lifetimes of molecular clouds, cluster dispersal, protoplanetary disk evolution and triggered star formation, can only be addressed by optical surveys of older populations with ages $\\sim$ 3-20 Myr. To make progress on these issues, wide-field studies of stellar populations in and near star-forming regions are required.  In the past, studies of OB associations have been used to investigate sequential star formation and triggering on large scales (e.g., \\citealt{bla91} and references therein). However, OB stars are formed essentially on the main sequence (e.g. \\citealt{pas92,pas93}) and evolve off the main sequence on timescales of order 10 Myr (depending upon mass and amount of convective overshoot); thus, they are not useful tracers of star-forming histories on timescales of 1-3 Myr. Moreover, we cannot investigate cluster structure and dispersal or disk evolution without studying low-mass stars. Wide-field studies of the low mass populations have been carried for some portions of the Upper Scorpius OB association \\citep{prz99,pbb02}. In Orion, many young individual clusters have been studied at both optical and infrared wavelengths (e.g. \\citealt{lad92,psh94,hil97,smm99}), but these only represent the highest-density, and often youngest, regions. Little is known about the more widely spread low mass stellar population in these OB associations.  A few large scale surveys have tried to map the widely distributed low-mass pre-main sequence population in Orion. The Kiso objective-prism H$\\alpha$ survey \\citep{wky89,wky91,wky93,kyw89,nwk95}, covered 225$\\Box^\\circ$ around the clouds but was biased towards the strongest H$\\alpha$ emitting stars. The ROSAT All-Sky Survey \\citep{atw96} detected many X-ray strong young stars widely distributed across Orion; however, this selection was strongly affected by older foreground field stars unrelated to the young population of Orion \\citep{bhs97} and because of its shallow limit it did not reach the fainter K and M stars. \\citep{wws98} carried out a detailed photometric and spectroscopic study of the region around the O9.5 type star $\\sigma$ Ori, belonging to 1b, and found a large concentration of pre-main sequence low mass stars down to V $\\sim$ 19. Similar high densities of young low mass stars have been found in other selected small regions in 1a and 1b by \\citet{sww99} and \\citet{she03}, indicating that many low mass stars remain undiscovered.  Until recently, it was not possible to conduct unbiased, large-scale optical surveys but technological advances have now made it possible, building on the availability of cameras with multiple CCDs and multiobject spectrographs on telescopes with wide fields of view. One such survey is the optical variability study we have conducted over $\\rm \\sim 173\\> deg^2$ in the Orion OB1 association with the purpose of  finding, mapping, and studying large numbers of widely-dispersed, low mass ($\\la 1 \\msun$) stars with ages from 1 to $\\sim$ 20 Myr.  Optical photometric variability is one of the defining characteristics of pre-main sequence stars \\citep{joy45,her62}. However, due to limitations in the size of CCD detectors, past work on photometric variability of young stars have concentrated on follow up studies of selected samples, that had been identified as pre-main sequence objects by some other means. Only very recently has the potential of variability as a technique to pick out young stars amongst the field population started to be realized \\citep{bvc01,lbm04}.  In this contribution we present results from our variability survey in an area of nearly $\\rm 68\\> deg^2$, that includes the Ori OB1b sub-association, a large fraction of Ori OB1a, and part of the Orion B molecular cloud \\citep{mmm87}, with the young clusters NGC 2023, 2024, 2068 and 2071. In \\S 2 we describe the optical variability survey and follow up spectroscopy. In \\S 3 we describe our results, and in \\S 4 we present a discussion and conclusions. ",
        "conclusions": "\\label{discussion}  A large fraction of the stars in our galaxy were formed in OB associations \\citep{wan00}. \\citet{adm01} suggest that many, if not most stars form in stellar groups with $10 < N < 100$ stars. These small groups are not strongly bound like more massive entities destined to become open clusters, and they evolve and disperse much more rapidly than clusters do. Sparse, low density regions like Taurus and Chamaeleon are more representative of what \\citep{adm01} call isolated star formation, while the $N\\sim 100$ groups are expected to form in OB associations. These groups are also expected to consist mostly of solar type and less massive stars; thus, they are not necessarily picked up by the existing studies of massive stars in these regions. Few studies exist that have been capable of mapping the low mass, young population over wide enough areas in nearby OB associations (e.g. \\citet{mml02}; \\citet{pzg97}; \\citet{dom02}).  We show here that multi-epoch surveys over very large areas of the sky are powerful tools for identifying young, low-mass stars in the older, extended regions of nearby OB associations were the molecular gas has dissipated and no longer serves as a marker of recent star formation. These new type of surveys build on the availability of CCD mosaic cameras on wide field telescopes of small to moderate size.  We have identified about 200 low mass ($\\rm M \\sim 0.12 - 0.9 \\msun$) pre-main sequence stars widely spread over nearly $\\rm 70\\> deg^2$ in the Orion OB1 association. The spatial distribution of the newly identified young stars indicate two large star forming events within this large area, supporting previous studies that recognized two distinct regions, the younger Ori OB1b and the older Ori OB1a. However, our sample is much more numerous, providing a better estimate of the properties of each sub-association than studies of the more massive stars. We find that 1b is almost completely confined to a large bubble of gas of $\\rm \\sim 2^\\circ$ radius roughly centered on the BOIae star $\\epsilon$ Ori. We also note that there seems to be a rather well defined edge to the OB association in the E-W direction, with essentially no T Tauri stars found at $\\rm RA \\la 5^h 10^m$, though many of our scans started further west than this. Another striking result is the discovery of a concentration of stars in 1a, located around an early type star, which also includes a possible member of the Herbig Ae/Be class. This is similar, but more numerous, than other small clusters like $\\eta$ Chamaleontis in the Sco-Cen association \\citep{lcm01}, possibly an example of the grouped type of star formation proposed by \\citep{adm01}.  Within present uncertainties (distances and depth of each association, differences between theoretical evolutionary models, etc.) we find an age of $\\sim 4-6$ Myr for 1b and $\\sim 7-10$ Myr for 1a. The fact that almost all accreting T Tauri stars (those with strong H$\\alpha$ emission), which are also the ones exhibiting near-IR excesses, are concentrated in 1b, together with the environmental characteristics of each region, supports the idea of 1b being much younger than 1a, consistent with previous estimates from the massive stars.  In Figure \\ref{diskfr} we plot the inner disk frequency as a function of age for a number of clusters and association, including our new data. The disk frequency has been determined by different methods. Haish, Lada, \\& Lada (2001) measured the excess in the near-IR bands JHKL in a number of young clusters, spanning the age range 0.3 - 30 Myr. The disk frequency in Ori OB1a and 1b has been determined by the fraction of CTTS; these fractions are  $\\sim$ 11\\% and 23\\%, respectively. We estimate the fraction of accreting stars in the TW Hydra association, using the breadth of the H$\\alpha$ profile from Muzerolle et al. (2000, 2001) as an indicator of accretion, and adopting 24 as the total number of members from Song et al. (2003). In general, there is a close correlation between the presence of inner disks, as indicated by near-infrared excesses, and the existence of accretion onto the central star (Hartmann 1998).  While not all T Tauri inner disks exhibit K-band excesses, depending upon inclination and magnetospheric hole size (Meyer et al. 1997; Muzerolle et al. 2003), the use of L-band data by Haisch \\etal (2001) should yield much more complete detections of disk excesses (Hartigan \\etal 1990).  The new data of Ori OB1a and 1b are consistent with the previously determined decrease of the disk fraction with age (Hillenbrand \\& Meyer 1999; Haisch \\etal 2001) and with the much more elaborate determination by Hillenbrand, Meyer, \\& Carpenter (2004, in preparation), supporting the idea of a timescale of $\\sim$ 10 Myr for the dissipation of inner disks around low mass, young stars.  Finally, the discovery of a significant population of low mass, pre-main sequence stars in a region devoid of gas like 1a favors the growing general idea that molecular clouds disperse over times of less than 10 Myr (Hartmann et al. 2001).   Acknowledgments. This work has been supported by NSF grant AST-9987367 and NASA grant NAG5-10545. C. Brice\\~no acknowledges support from grant S1-2001001144 of FONACIT, Venezuela. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This research has made use of the NASA/ IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. Grant. We are grateful to Susan Tokarz at CfA, who is in charge of the reduction and processing of FAST spectra. We also thank Dr. Robert W. Wilson of the CfA Radio \\& GeoAstronomy Division for making available to us the $\\rm ^{13}CO$ map shown in Figure 6. We thank the invaluable assistance of the observers and night assistants at the Venezuela Schmidt telescope that made possible obtaining the data over these past years: A. Bongiovanni, C. Castillo, O. Contreras, D. Herrera, M. Mart{\\'\\i}nez, F. Molina, F. Moreno, L. Nieves, G. Rojas, R. Rojas, L. Romero, U. S\\'anchez, L. Torres and K. Vieira. We also acknowledge the support from the CIDA technical staff, and in particular of Gerardo S\\'anchez. Finally, we acknowledge the comments from an anonymous referee, which helped improved this article."
    },
    "0410/astro-ph0410535_arXiv.txt": {
        "abstract": "We present the results from an analysis of surface photometry of $B$, $R$, and H$\\alpha$ images of a total of 114 nearby galaxies ($v_{\\mathrm{helio}}$$<$4000\\,km\\,s$^{-1}$) drawn from the Palomar/Las Campanas Imaging Atlas of Blue Compact Dwarf galaxies. Surface brightness and color profiles for the complete sample have been obtained. We determine the exponential and S\\'ersic profiles that best fit the surface brightness distribution of the underlying stellar population detected in these galaxies. We also compute the ($B-R$) color and total absolute magnitude of the underlying stellar population and compared them to the integrated properties of the galaxies in the sample. Our analysis shows that the ($B-R$) color of the underlying population is systematically redder than the integrated color, except in those galaxies where the integrated colors are strongly contaminated by line and nebular-continuum emission. We also find that galaxies with relatively red underlying stellar populations (typically ($B-R$)$\\geq$1\\,mag) show structural properties compatible with those of dwarf elliptical galaxies (i.e$.$ a smooth light distribution, fainter extrapolated central surface brightness and larger scale lengths than BCD galaxies with blue underlying stellar populations). At least $\\sim$15\\% of the galaxies in the sample are compatible with being dwarf elliptical (dE) galaxies experiencing a burst of star formation. For the remaining BCD galaxies in the sample we do not find any correlation between the recent star formation activity and their structural differences with respect to other types of dwarf galaxies. ",
        "introduction": "\\label{intro}  The vast majority of the Blue Compact Dwarf (BCD) galaxies are currently forming stars at a very high rate, as suggested by the large fraction of them ($>$93\\%; Gil de Paz, Madore, \\& Pevunova 2003) showing significant H$\\alpha$ emission (EW$>$20\\,\\AA). The star formation rate (SFR) derived for these galaxies ranges up $\\sim$10\\,M$_{\\odot}$\\,yr$^{-1}$ (Fanelli et al$.$ 1988; Gil de Paz et al$.$ 2003). This relatively high SFR, combined with a moderate gas content, implies gas-consumption time-scales of $\\sim$10$^{9}$\\,yr, which are much shorter than the age of the Universe.  This high SFR, combined with the low present-day metal abundances (see e.g$.$ Hunter \\& Hoffman 1999), led Searle et al$.$ (1973) to suggest that either these objects are intrinsically young galaxies or that they have had an episodic star formation history involving very short bursts of star formation followed by long quiescent periods. It is now widely accepted that most of these objects possess a relatively evolved underlying stellar population (USP hereafter) associated with their low-surface-brightness (LSB) envelopes, i.e$.$ they are not exclusively young galaxies (Schulte-Ladbeck et al$.$ 1999, 2001; Drozdovsky et al$.$ 2001; Crone et al$.$ 2000, 2002; Gil de Paz et al$.$ 2000a,b; Papaderos et al$.$ 1996a,b). Despite considerable recent progress, the properties of this USP are still poorly known. Although the presence of an evolved population supports a relatively episodic star formation history it has been recently shown that these episodes of star formation may, in fact, last as long as 10$^{8}$\\,yr (Crone et al$.$ 2002; Papaderos et al$.$ 2002). However, it is still necessary that the SFR in BCDs had been much lower in the past (and for extended periods of time) as compared to today. It has been suggested that, even during the so-called quiescent phases, star formation could have taken place at a very low level, at a rather continuous rate (Legrand 2000).  The following questions naturally arise: What did BCD galaxies look like during these long periods of time of low (or null) star-formation activity? Were they similar to local dwarf elliptical galaxies today? and, is there a single evolutionary model that can include all phases and types of dwarf galaxies?  Current knowledge bearing on these questions comes mainly from the analysis of the surface brightness profiles of the USP in BCDs and a comparison of the profiles of dwarf irregular and dwarf elliptical galaxies (Papaderos et al$.$ 1996; Noeske et al$.$ 2000; Cair\\'os et al$.$ 2003). These results seem to suggest that the LSB envelopes associated with the USP of BCD galaxies are more compact and have a higher central surface brightness than those seen in dI and dE galaxies. Some additional clues have been recently provided by the study of the dynamics of BCD neutral hydrogen halos (van Zee, Salzer, \\& Skillman 2001; Tajiri \\& Kamaya 2002). These authors have shown that the suggested evolution between BCDs and dE is difficult because of the relatively low efficiency of stellar feedback in potentially removing the HI envelopes from these galaxies. It is worth emphasizing here that in all of these previous studies the number of objects under consideration was quite limited and no definitive conclusions could be drawn concerning existence or not of a unified evolutionary model for dwarf galaxies.  In Paper~I (Gil de Paz et al$.$ 2003) we described the integrated properties of a sample of 114 nearby galaxies in the Palomar/Las Campanas Imaging Atlas of Blue Compact Dwarfs (BCD hereafter). In this, the second paper of the series, we analyze the surface brightness profiles of the sample in $B$ and $R$ bands and in the light of H$\\alpha$. The morphological information derived, along with the properties of the USP detected in these galaxies, is now compared with the integrated properties measured in Paper~I.  In Section~\\ref{sampleobs} we briefly describe the sample and the observations carried out within the Atlas. The procedures used to derive the surface brightness and color profiles and the corresponding best-fitting exponential and S\\'ersic laws are described in Sections~\\ref{analysis.profiles} and \\ref{analysis.fitting}, respectively. The analysis of the surface brightness profiles is presented in Section~\\ref{results.sbp}. We describe the structural properties, color, and luminosity of the USP of these galaxies in Section~\\ref{results.lsb}. Finally, we discuss the implications of this study regarding the existence of possible evolutionary links between BCDs and dwarf elliptical galaxies (Section~\\ref{results.dwarfs}), and the possible impact of the recent star formation on the structural properties of BCDs (Section~\\ref{results.sfh}). Our conclusions are summarized in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In summary,  \\begin{itemize}  \\item We have presented the surface brightness profiles in $B$, $R$, and H$\\alpha$, for a total of 114 galaxies taken from the Palomar/Las Campanas Imaging Atlas of BCD galaxies. A total 104 of the galaxies are classified as BCDs (see Paper~I). The profiles in the continuum bands are characterized by the presence of a HSB component on top of the nearly exponential low-surface-brightness component associated with the galaxy's underlying stellar population (USP). At large galactocentric radii the color profiles of 70\\% of the galaxies flatten. This flattening occurs approximately at the position where the USP starts to dominate the galaxy surface brightness profile.  \\item The color of the USP (corrected for line and nebular-continuum emission) is systematically redder than the observed integrated color. The color of the USP is bluer than the integrated one only in those objects with the highest equivalent widths of H$\\alpha$ (where the line and nebular-continuum emission is significant; i.e$.$ EW(H$\\alpha$) larger that a few hundred angstroms).  \\item We find that galaxies with relatively red USP (($B-R$)$\\geq$1\\,mag) show structural properties compatible with those of dwarf elliptical galaxies. They show smoother (continuum) light distributions, fainter extrapolated central surface brightness, and larger scale lengths than BCD galaxies with blue envelopes. This result indicates that a non-negligible fraction of the BCD galaxies ($\\sim$15\\%) could be dwarf ellipticals that are now experiencing (or have recently experienced) an episode of star formation.  \\item We do not find any correlation between the equivalent width of H$\\alpha$ and the {\\it degree of expansion} of the USP, the latter being defined as the difference in scale length between BCD and dE galaxies of identical luminosity. The difference measured between the scale length of BCDs and dE galaxies is much larger than that expected from changes in the gravitational potential due to the collective effect of supernova-driven winds, especially considering the relatively small differences in ($B-R$) color between the galaxies and their USP. This suggests that the level of recent star formation in BCDs does not have a significant impact on the structural properties of these galaxies. A detailed study of a large number of individual objects is needed to confirm this.  \\end{itemize}"
    },
    "0410/astro-ph0410229_arXiv.txt": {
        "abstract": "We present high-resolution VLA observations of the molecular gas in the host galaxy of the highest redshift quasar currently known, SDSS\\,J1148+5251 (z=6.42). Our VLA data of the CO(3--2) emission have a maximum resolution of 0.17$''\\times0.13''$ ($\\leq$1\\,kpc), and enable us to resolve the molecular gas emission both spatially and in velocity. The molecular gas in J1148+5251 is extended to a radius of 2.5\\,kpc, and the central region shows 2 peaks, separated by 0.3$''$ (1.7\\,kpc). These peaks account for about half of the total emission, while the remainder is more extended. Each of these unresolved peaks contains a molecular gas mass of $\\sim5\\times10^9$\\,M$_{\\odot}$ (similar to the total mass found in nearby ULIRGS) and has an intrinsic brightness temperature of $\\sim35$\\,K (averaged over the 1\\,kpc-sized beam), comparable to what is found in nearby starburst centers.  Assuming that the molecular gas is gravitationally bound, we estimate a dynamical mass of $\\sim4.5\\times10^{10}$\\,M$_{\\odot}$ within a radius of 2.5\\,kpc ($\\sim5.5\\times10^{10}$\\,M$_{\\odot}$ if corrected for a derived inclination of $i\\sim65^{\\circ}$). This dynamical mass estimate leaves little room for matter other than the detected molecular gas, and in particular the data are inconsistent with a $\\sim10^{12}$\\,M$_\\sun$ stellar bulge which would be predicted based on the M$_{\\rm BH}-\\sigma_{\\rm bulge}$ relation. This finding may indicate that black holes form prior to the assembly of the stellar bulges and that the dark matter halos are less massive than predicted based on the black hole/bulge mass relationship. ",
        "introduction": "\\label{par:intro}  More and more objects have been found at the highest redshifts in recent years, back to the 'Dark Ages' (the epoch of formation of the first luminous structures), at z$>$6 when the universe was less than a gigayear old (e.g., Fan et al.\\ 2002, 2003, 2004, Hu et al.\\ 2002, Stanway et al.\\ 2003). Optical spectra of the brightest quasars at z$>$6 show a clear Gunn \\& Peterson (1965) effect, demonstrating that these objects are located at the end of cosmic reionization (Becker et al.\\ 2001, Fan et al.\\ 2003, White et al.\\ 2003). It is of paramount importance to study these sources in detail to measure the stellar and gaseous constituents, chemical abundances and dynamical masses of the host galaxies, which in turn constrain galaxy evolution models in the very early universe. Studies of the host galaxies of bright quasars are extremely challenging in the optical and/or near-infrared as the central active galactic nucleus (AGN) greatly overshines its surrounding.  Radio observations of the molecular gas phase, however, are a promising method to measure such important properties of distant quasar host galaxies as the total quantity of molecular gas and the dynamical masses.  The highest redshift quasar currently known is SDSS\\,J114816.64+525150.3 (hereafter: J1148+5251) at z=6.42 (Fan et al.\\ 2003). J1148+5251 is a highly luminous object which is thought to be powered by mass accretion onto a supermassive black hole of mass 1--5$\\times10^9$ M$_{\\odot}$ (which accrets at the Eddington limit, Willot et al.\\ 2003).  Thermal emission from warm dust has been detected at millimetre wavelengths, implying a far infrared (FIR) luminosity of $1.3\\times 10^{13}$\\,L$_\\odot$ (Bertoldi et al.\\ 2003a), corresponding to about 10$\\%$ of the bolometric luminosity of the system. J1148+5251 falls on the radio-FIR correlation found for nearby starburst galaxies (Carilli et al.\\ 2004); if the heating of the dust is dominated by young stars, then the implied star formation rate is an impressive $\\sim3000$\\,M$_{\\odot}$\\,yr$^{-1}$. High--resolution HST imaging does not reveal multiple imaging of the central AGN (White et al.\\ 2004), rendering strong magnification unlikely (cf. the multiple optical imaging in the gravitationally lensed system PSS2322+1944, Carilli et al.\\ 2003).  Using the Very Large Array (VLA) and the Plateau de Bure interferometer (PdBI), we have detected various rotational lines of carbon monoxide (CO), the most common tracer of molecular gas, in J1148+5251 (Walter et al. 2003, Bertoldi et al. 2003b). This discovery demonstrated that abundant, dense and warm molecular gas (the requisite fuel for star formation) is already present at these large lookback times (a redshift of z=6.4 corresponds to an age of the universe of 870\\,Myr in a cosmology with H$_0\\!=\\!71$\\,km\\,s$^{-1}\\,$Mpc$^{-1}$, $\\Omega_{\\Lambda}$=0.73, $\\Omega_{m}$=0.27 which we adopt in the following).  Here we present follow-up high-resolution observations of the molecular gas distribution of J1148+5251 obtained with the NRAO VLA\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}. These observations resolve, for the first time, the host galaxy of this luminous quasar on kiloparsec scales (1\\,kpc = 0.18$''$ at z=6.42) and enable us to constrain its molecular gas properties and the dynamical mass. ",
        "conclusions": "We present the first resolved maps of a system located at the end of cosmic reionization at z=6.4. Our high-resolution VLA data enable us to resolve the molecular gas emission in the host galaxy both spatially and in velocity space.  The molecular gas distribution in J1148+5251 is extended out to radii of 2.5\\,kpc. The central region is resolved and shows 2 peaks, separated by 1.7\\,kpc; they account for about half of the total emission, with the other half present in the more extended molecular gas distribution. Each of these peaks harbors a molecular gas mass of $\\sim5\\times10^9$\\,M$_{\\odot}$ within a radius of 0.5\\,kpc, respectively; this mass is similar to the total mass found in nearby ULIRGS such as Mrk\\,273 or Arp\\,220 (Downes \\& Solomon 1998). The peaks have intrinsic brightnesses of $\\sim35$\\,K (averaged over the 1\\,kpc-sized beam), similar to what is found in the centres of nearby active galaxies (Downes \\& Solomon 1998, assuming constant surface brightness down to the CO(1--0) transition), albeit measured over a larger physical area.  We have assumed that the gas is gravitationally bound and is situated in a disk, although the data do not rule out the presence of an unbound system, such as an ongoing merger. Based on the extent of the molecular gas distribution and the line-width measured from the higher CO transitions we derive a dynamical mass of $\\sim4.5\\times10^{10}$\\,M$_{\\odot}$ ($\\sim5.5\\times10^{10}$\\,M$_{\\odot}$ if we correct for an inclination of $i\\sim65^{\\circ}$).  This dynamical mass estimate can account for the detected molecular gas mass within this radius but leaves little room for other matter. In particular, given a black hole mass of mass $\\sim$1-5$\\times10^9$ M$_{\\odot}$ (Willot et al.\\ 2003), this dynamical mass could not accomodate an order few$\\,\\times 10^{12}$\\,M$_{\\odot}$ stellar bulge which is predicted by the present--day M$_{\\rm BH}-\\sigma_{\\rm bulge}$ relation (Ferrarese \\& Merritt 2000, Gebhardt 2000), if this relation were to hold at these high redshifts. Even if we assume a scenario in which this bulge was 10 times the scale length of the molecular gas emission, we would still expect a bulge contribution of few$\\,\\times 10^{11}$\\,M$_{\\odot}$ within the central 2.5\\,kpc (assuming a central density--profile of $\\rho\\sim r^{-2}$, e.g., Jaffe 1983, Tremaine 1994) which can not be reconciled with our results. Our finding therefore suggests that black holes may assemble before the stellar bulges. This would be in conflict with the popular picture of co-eval evolution of the central black hole and the stellar bulge. The derived dynamical mass and the hypothetical stellar bulge mass (based on M$_{\\rm BH}-\\sigma_{\\rm bulge}$) can be only brought in agreement if (i) our derived inclination is significantly off (i.e., $i<10^\\circ$), (ii) the estimated black hole mass is over-estimated by orders of magnitudes (which would imply highly super-Eddington accretion), or (iii) if the assumption that the gas is gravitationally bound is seriously wrong. We note, however, that depending on the space density of similar sources at these high redshifts, the smaller dynamical masses suggested in this study may be in better agreement with the masses predicted by CDM simulations in the very early universe as a 10$^{12}$\\,M$_{\\odot}$ bulge would imply a rather massive dark matter halo of $>10^{14}$\\,M$_{\\odot}$.  This study shows the potential for future studies of the molecular gas content and dynamical masses in the highest redshift galaxies using ALMA, where resolutions of $<0.1''$ will be achieved routinely."
    },
    "0410/astro-ph0410159_arXiv.txt": {
        "abstract": "We report a study of the spectra and structure of the accretion disc of the V2051 Oph while the star was in an unusual, faint brightness state during 1996. ",
        "introduction": "V2051 Oph ($P_{orb}=90\\,min$) belongs to a small group of ultrashort period eclipsing dwarf nova together with Z Cha, OY Car and HT Cas. During 1996, it was seen at a low brightness state ($B\\simeq16,2$), in which mass accretion was reduced to a minimum (Baptista et al. 1998).  High-speed ultraviolet and optical spectroscopy was secured with the Faint Object Spectrograph onboard the Hubble Space Telescope on 1996 Jan 29 (UV: G400H grating, $\\lambda\\ 3226-4781$~\\AA,~$\\Delta\\lambda= 1.5$ \\AA\\ pixel$^{-1}$; optical: G160L grating, $\\lambda\\ 1150-2507$~\\AA,~$\\Delta\\lambda= 3.5$ \\AA\\ pixel$^{-1}$) at a time resolution of 3.4\\,s.  The runs cover the eclipse cycles 109\\,988 (UV) and 109\\,989 (optical), according to the ephemeris of Baptista et al. (2003).  The out-of-eclipse UV spectra show prominent Ly$\\alpha$ $\\lambda 1216$, C\\,II $\\lambda 1323$, Si\\,IV $\\lambda 1394,1403$ and C\\,IV $\\lambda 1549,1551$ emission lines as well as broad absorption bands possibly due to Fe\\,II. The optical spectra show the Balmer continuum in emission as well as He emission lines and strong Balmer emission lines.  The UV spectra were divided into 34 narrow passbands (19-60 \\AA\\ wide in the continuum and velocity-resolved bins of 2000~km/s or 3000~km/s for the lines) while the optical spectra were divided into 68 passbands (15-42 \\AA\\ wide in the continuum and velocity-resolved bins of 1000~km/s for the lines) and light curves were extracted for each one.  Maximum-entropy eclipse mapping techniques (Baptista \\& Steiner 1993) were used to solve for a map of the disc brightness distribution and for the flux of an additional uneclipsed component in each band. ",
        "conclusions": ""
    },
    "0410/astro-ph0410190_arXiv.txt": {
        "abstract": "The behavior of the quintessence field is studied during inflation. In order to have a satisfactory model of dark energy, the quintessence field value today should be as insensible to the initial conditions as possible. Usually, only the dependence on the initial conditions specified at the end of inflation or, equivalently, at the beginning of the radiation dominated era, is considered. Provided the quintessence field is initially within a large but, crucially, finite interval, its present value becomes independent of the initial value it started from. The question as to whether inflation naturally drives the quintessence field to the above-mentioned interval is addressed. Since the quantum effects turn out to be important, the formalism of stochastic inflation is used in order to calculate the evolution of the quintessence field. Moreover, the quantum effects originating from the inflaton field are also taken into account and are proved to be sub-dominant in most cases. Finally, the requirement that the quintessence field is on tracks today is shown to imply quite tight constraints on the initial values of the quintessence and inflaton fields at the beginning of inflation. In particular, the initial value of the inflaton field cannot be too large which indicates that the quintessential scenario seems to be compatible with inflation only if the total number of e-folds is quite small. ",
        "introduction": "The observations suggesting that our Universe is presently undergoing a phase of accelerated expansion have recently accumulated~\\cite{SNIa, wmap, sdss,ISW}. If really confirmed, this discovery is certainly a breakthrough for cosmology but, at the same time, represents a big challenge since finding a convincing explanation for such a phenomenon is clearly a difficult task.  \\par  {}From a theoretical point of view, the presence of a non-vanishing cosmological constant whose energy density would be of the order of the critical energy density today seems to be the most natural solution. In addition, the currently available data on the equation of state are, so far, compatible with this assumption. But it is well-known that the theoretical preferred value of the cosmological constant corresponds to an energy density much larger than the critical energy density and there exists, at the moment, no convincing arguments which would explain this difference~\\cite{Wein}.  \\par  This situation has led the physicists to seek for alternatives. Among the solutions proposed, the quintessence scenario has recently attracted a lot of attention~\\cite{RP, quint, ZWP,PB,BM1,BM2,BMR1,BMR2,highphysquint}. It consists in postulating that the acceleration of the expansion is caused by a scalar field, the quintessence field $Q$, evolving in a potential the typical shape of which is given by $W(Q)=M^{4+\\alpha }Q^{-\\alpha }$, where $M$ is an energy scale and $\\alpha >0$ a free parameter~\\cite{RP}. The main advantage of this scenario is that the coincidence problem can be solved because the equations of motion possess a solution which is an attractor. Therefore, the present evolution of the quintessence field is independent from the choice of the initial conditions. Moreover, when the field is on tracks, $Q$ is typically of the order of the Planck mass and hence, for not too small values of $\\alpha $, the scale $M$ can be large. As a result, the fine-tuning is less severe than in other scenarios because it is possible to explain the presence of a very small scale (the vacuum energy density today) by means of a theory which, on the contrary, is characterized by a large scale $M$. This is due to the inverse power-law shape of the potential and is reminiscent of the ``see-saw'' mechanism in particle physics. This has also the advantage that model building can be considered in the realm of high energy physics~\\cite{PB,BM1,BM2,BMR1,BMR2,highphysquint}.  \\par  So far most of the studies have been devoted to understanding how the quintessence field evolves from the beginning of the radiation era until now. Another important (related) question, in view of its observational implications, has been to estimate the value of the equation of state today. In this paper, we address a new question, namely that of the behavior of $Q$ during (chaotic) inflation assuming, for simplicity, that the quintessence field and the inflaton are not coupled. This is an important problem since it is crucial to check that $Q$ is, at the end of inflation (or at the beginning of the radiation dominated epoch), in the range of values which are such that the field is on tracks today.  \\par  However, the problem does not only boil down to solving the Klein-Gordon equation in an inflationary background. Indeed, in Ref.~\\cite{ML}, it has been suggested that the quantum effects could play an important role. In this case, the techniques of stochastic inflation~\\cite{fokkerplanck,stocha} can be used to describe the evolution of the quintessence field. This was done for the first time in Ref.~\\cite{ML}. The method utilized in that article was to solve the Fokker-Planck equation in order to follow the evolution of the probability distribution of the quintessence field. It was then shown that, typically, the attractor is joined at relatively small redshifts.  \\par  In the present paper, we consider the above-mentioned question again but from a different perspective. One of our main purposes is to calculate the probability distribution function of the quintessence field at the end of inflation (or at the beginning of the radiation dominated era) in order to estimate whether it is likely that the value of $Q$ is such that the attractor is joined today. Moreover, requiring that the corresponding probability is significant can be used to constrain the space of the initial conditions, \\ie the initial value of the inflaton (or, equivalently, the total number of e-folds) and quintessence fields. In addition, we demonstrate that this also puts constraints on the power index $\\alpha $ characterizing the shape of the quintessence potential, namely small values of $\\alpha $ are disfavored.  \\par  Another goal of the present work is to include the inflaton fluctuations, to study under which circumstances their effect can be important and, when it is the case (and when it is possible), to calculate the corresponding correction to the behavior of the quintessence field. Indeed, in Ref.~\\cite{ML}, since the inflaton was treated as a classical field, the quantum effects were only sourced by the quintessence noise. However, the inflaton itself is also influenced by the quantum effects and, therefore, a priori the inflaton noise also affects the evolution of the quintessence field. In fact, the variance of the quintessence field can be written as \\begin{equation} \\sigma ^2(t)=\\frac{1}{4\\pi ^2}\\int _{\\tin}^t H^3\\left[\\varphi \\left(\\tau \\right)\\right]{\\rm d}\\tau  \\, , \\end{equation} where $H$ is the Hubble parameter. Roughly speaking, taking into account the inflaton noise amounts to put the coarse-grained inflaton in the above equation rather that its classical counterpart. If the expression calculated in this way differs significantly from the expression obtained by inserting the classical inflaton, then the inflaton noise plays indeed a non negligible role.  \\par  Another difference from Ref.~\\cite{ML} is that we directly solve the Langevin equation rather than the Fokker-Planck equation. Obviously, this is only a technical difference since the two approaches are equivalent. In a first time, the Langevin equation is solved by means of a perturbative expansion. In this regime, we show that the influence of the inflaton noise is always negligible. Then, in a second time, we try to solve the Langevin equation in the non-perturbative regime (for the quintessence field) by modeling the effect of the classical force with a wall.  \\par  This article is organized as follows. In the next section, we quickly review the basic principles and equations of the stochastic approach. Then, in Sec.~III, we present the perturbative method used to solve the Langevin equation. We apply this method to inflation, compare the results obtained with those already known in the literature and demonstrate that they are equivalent. In Sec.~IV, we apply the perturbative approach to the Langevin equation for the quintessence field and explicitly calculate the quintessential quantum effects. In this regime, we show that the contribution coming from the inflaton noise is negligible. Then, we present a model with a reflecting wall which allows us to explore a region where the perturbative approach breaks down. We study the constraints on the initial conditions of the inflaton and quintessence fields that exist in order for the coincidence problem to still be solved. We prove that these constraints are quite stringent. Finally, in Sec.~V, we discuss the results obtained in this article and present our conclusions. ",
        "conclusions": "In this section, we briefly discuss other aspects of the question studied in this article and present our conclusions.  \\par  In the preceding sections, we have calculated the evolution of the quintessence field during inflaton taking into account the quantum effects. What about these quantum effects in the subsequent cosmological eras? After the reheating, during the radiation and matter dominated phases, it is easy to show that the quintessence field evolves classically starting from the value reached at the end of inflation and determined by its random walk during inflation. Indeed, since the noise in the Langevin equation is controlled by the quantum fluctuations of the modes leaving the horizon, it rapidly becomes negligible as the accelerated expansion stops and the modes start to reenter the horizon.  \\par  However, a new phase of accelerated expansion driven by the quintessence field is now taking place and, clearly, the above argument does not apply in this case.  Therefore, one may wonder whether the influence of the quantum effects should not be taken into account when one computes the evolution of $Q$ at present time. This could have important observational consequences, in particular if the stochastic behavior of $Q$ modifies the value of the equation of state. However, it is easy to demonstrate that this is not so. The problem is in fact very similar to calculating the evolution of the inflaton field during inflation since, at present time, the quintessence field is no longer a test field but actually determines the evolution of the background. Therefore, even though the slow roll approximation might not be so satisfactory in this case, we can at least estimate the relevance of these late stochastic fluctuations by simply setting $n=-\\alpha$ into~Eq.~\\eqref{varinflaton} (since, in some sense, the Ratra-Peebles potential is nothing but an inflationary chaotic potential with a negative index). This equation, provided one substitutes $\\varphi _\\cl$ with $Q_\\cl$ and $V(\\varphi _\\cl)$ with $W(Q_\\cl)$, should be a reasonable estimate of $\\mean{\\delta Q^2}/Q^2$. Since the quintessence field is now on the attractor, its value must be of the order of the Planck mass. In addition, it is slowly rolling down its potential toward higher values. As a consequence, from Eq.~\\eqref{varinflaton}, one gets the following rough estimate \\begin{equation} \\frac{\\mean{\\delta Q^2}}{Q^2}\\bigg\\vert_\\mathrm{now} \\lesssim \\frac{W(Q_{\\rm now})}{\\mP^4}\\simeq 10^{-123}\\, , \\end{equation} meaning that any stochastic deviation from the classical trajectory is completely negligible at present time.  \\par  Let us now end this work by reviewing what are the main conclusions of our study. The main result is that, taking into account the quantum effects during inflation is important since the stochastic diffusion term dominates over the classical drift term and that, as a consequence, the value of $\\mean{Q}$ can differ from $Q_\\cl$ by several orders of magnitude. For the first time to our knowledge, we have given an analytical estimate describing the evolution of $\\mean{Q}$ during inflation, see Eq.~(\\ref{meanwall}).  \\par  Another new result is the fact that, requiring the quintessence field to have a large probability to be on tracks today, allows us to put quite stringent constraints on the initial conditions $Q_\\ini$ and $\\varphi _\\ini$. Typically, the quintessence field must start from small values. We have also established that large values of $\\alpha $ are favored (we notice in passing that, if we consider the Ratra-Peebles potential only and not the SUGRA one, the same conclusion is reached from the constraints that exist on the equation of state today). Another interesting result is that the inflaton field must also start from quite small values. This implies that the total number of e-folds during inflation is also limited. On the other hand, we have remarked the existence of an attractor for $\\mean{Q}$, see Eq.~\\eqref{attra}, due to the fact that the final value of $\\mean{Q}$ is independent of $Q_\\ini$. However, a dependence in the initial conditions is reintroduced in the calculation of the probability which has allowed us to put the constraints mentioned just before.  \\par  One of the main purpose of our paper was also to study the influence of the inflaton noise on the evolution of the quintessence field. The approximation consisting in neglecting the inflaton fluctuations has been shown to be justified in most cases, basically because the corresponding contributions to the mean value and/or to the variance are proportional to $V_\\ini/\\mP^4$, see for instance Eq.~(\\ref{finalQ}). Even in the extreme case of Planckian initial conditions for the inflaton field (\\ie $V_\\ini\\sim\\mP^4$), the inflaton noise is unlikely to modify $\\mean{Q}/\\mP$ by more than one order of magnitude compared to what is obtained taking into account the quintessence noise only.  \\par  It is also interesting to compare these results to those obtained in the paper~\\cite{ML} which was the first to take into account the quantum effects in the calculation of the evolution of $\\mean{Q}$. Basically, our findings confirm and/or justify the results of Ref.~\\cite{ML} and somewhat extend their validity. We have recovered the same equation for the variance and our new equation~(\\ref{meanwall}) for the mean value of $Q$ confirms the conclusions that can be drawn from the figures of Ref.~\\cite{ML}, namely that the quantum effects can play an important role during inflation. In Ref.~\\cite{ML}, the inflaton noise has not been considered and, as mentioned above, we have demonstrated that this is, in most cases, a good approximation.  \\par  Finally, let us describe some questions that are left unanswered and some possible improvements to the present study. In order to be able to find analytical solutions, we have modeled the classical drift term with a reflecting wall. The price to pay is that we have lost the dependence in the parameter $\\alpha $. The drift term acts differently for different Ratra-Peebles potentials while the wall repels the field regardless to $\\alpha $. Although we do not expect a strong dependence, it would be interesting to quantify this effect. The problem is that, if one includes the exact classical term, then the Langevin equation is no longer analytically solvable. The only way out seems to numerically integrate this equation. However, even this solution could be difficult because a term like $Q^{-\\alpha }$ can rapidly become very large and, hence, problematic from the numerical point of view.  \\par  Another interesting question would be to study what happens when one considers the case of a colored noise since it is clear that a white noise is not, physically, the most relevant case. Concretely, this amounts to replace the Heaviside function in the expansion of the field by a smooth function and, in principle, this could affect the evolution of the quintessence field during inflation although, again, we do not expect a very important effect.  \\par  For the moment, we postpone the study of all the issues to future works.  \\vspace{0.5cm} \\centerline{\\bf Acknowledgments} \\vspace{0.2cm}  We wish to thank S.~Matarrese for many enlightening comments and discussions.  \\appendix"
    },
    "0410/astro-ph0410473_arXiv.txt": {
        "abstract": "Current empirical evidence on the star-formation processes in the extreme, high-pressure environments induced by galaxy encounters (mostly based on high-resolution {\\sl Hubble Space Telescope} observations) strongly suggests that star {\\it cluster} formation is an important and perhaps even the dominant mode of star formation in such starburst events. The sizes, luminosities, and mass estimates of the young massive star clusters (YMCs) are entirely consistent with what is expected for young Milky Way-type globular clusters (GCs). Recent evidence lends support to the scenario that GCs, which were once thought to be the oldest building blocks of galaxies, are still forming today. Here, I present a novel empirical approach to assess the shape of the initial-to-current YMC mass functions, and hence their possible survival chances for a Hubble time. ",
        "introduction": "The production of luminous, massive yet compact star clusters seems to be a hallmark of the most intense star-forming episodes in galaxies. Young massive star clusters (YMCs; with masses often significantly exceeding $M_{\\rm cl} = 10^5 {\\rm M}_\\odot$) are generally found in intense starburst regions, often in galaxies involved in gravitational interactions of some sort (e.g., de Grijs et al.  2001, 2003a,b,c,d,e and references therein).  An increasingly large body of observational evidence suggests that a large fraction of the star formation in starbursts actually takes place in the form of such concentrated clusters, rather than in small-scale star-forming ``pockets''. YMCs are therefore important as benchmarks of cluster formation and evolution. They are also important as tracers of the history of star formation of their host galaxies, their chemical evolution, the initial mass function (IMF), and other physical characteristics in starbursts.  Using optical observations of the ``Mice'' and ``Tadpole'' interacting galaxies (NGC 4676 and UGC 10214, respectively) -- based on a subset of the Early Release Observations obtained with the {\\sl Advanced Camera for Surveys} on board the {\\sl Hubble Space Telescope (HST)} -- and the novel technique of pixel-by-pixel analysis of their colour-colour and colour-magnitude diagrams, we deduced the systems' star and star cluster formation histories (de Grijs et al. 2003e). In both of these interacting systems we find several dozen YMCs (or, alternatively, compact star-forming regions), which overlap spatially with regions of active star formation in the galaxies' tidal tails and spiral arms (from a comparison with H$\\alpha$ observations that trace active star formation; Hibbard \\& van Gorkom 1996). The tidal tail of the Tadpole system is dominated by star-forming regions, which contribute $\\sim 70$\\% of the total flux in the {\\sl HST} F814W filter (decreasing to $\\sim 40$\\% in the F439W filter). If the encounter occurs between unevenly matched, gas-rich galaxies then, as expected, the effects of the gravitational interaction are much more pronounced in the smaller galaxy. For instance, when we compare the impact of the interaction as evidenced by star cluster formation between M82 (de Grijs et al. 2001, 2003b,c) and M81 (Chandar et al.  2001), or the star cluster formation history in M51 (Bik et al. 2003), which is currently in the process of merging with the smaller spiral galaxy NGC 5194, the evidence for enhanced cluster formation in the larger galaxy is minimal if at all detectable.  Nevertheless, we have shown that star cluster formation is a major mode of {\\it newly-induced} star formation in galactic interactions, with $\\ge 35$\\% of the active star formation in encounters occurring in YMCs (de Grijs et al. 2003e).  The question remains, however, whether or not at least a fraction of the numerous compact YMCs seen in extragalactic starbursts, may be the progenitors of GC-type objects. If we could settle this issue convincingly, one way or the other, the implications of such a result would have profound and far-reaching implications for a wide range of astrophysical questions, including (but not limited to) our understanding of the process of galaxy formation and assembly, and the process and conditions required for star (cluster) formation. Because of the lack of a statistically significant sample of similar nearby objects, however, we need to resort to either statistical arguments or to the painstaking approach of one-by-one studies of individual objects in more distant galaxies. ",
        "conclusions": ""
    },
    "0410/astro-ph0410645_arXiv.txt": {
        "abstract": "Magnetohydrodynamic simulations of fully convective, rotating spheres with volume heating near the center and cooling at the surface are presented. The dynamo-generated magnetic field saturates at equipartition field strength near the surface. In the interior, the field is dominated by small-scale structures, but outside the sphere by the global scale. Azimuthal averages of the field reveal a large-scale field of smaller amplitude also inside the star. The internal angular velocity shows some tendency to be constant along cylinders and is ``anti-solar'' (fastest at the poles and slowest at the equator). ",
        "introduction": "The Hayashi track in the Hertzsprung--Russell diagram characterizes young stars in hydrostatic equilibrium that are fully convective. Other fully convective stars are low-mass main sequence stars (M dwarfs), and some cool giants. These stars show strong magnetic activity as is evidenced by chromospheric emission in H$\\alpha$ \\citep[e.g.][]{Hawley:ActivityLowMass,HawleyEtal:ChromosphericActivity} and by Zeeman broadening of classical T Tauri stars \\citep[e.g.][]{JohnsKrullEtal:BFieldBPTauri}. In the latter case, the stars are generally rapidly rotating with rotation periods of just a few days, and it is known that the magnetic field shows strong departures from axisymmetry \\citep{JohnsKrullEtal:SpectrapolarimetryBPTauri}. However, for less massive stars (M9 dwarfs and beyond) there is a sharp decline in chromospheric magnetic activity \\citep[e.g.][]{GizisEtal:NewNeighbors}, which may be connected with dust formation and the almost fully neutral photospheres \\citep{MohantyBasri:RotationActivity}.  Despite some progress in low-resolution Doppler imaging \\cite[e.g.][]{JoncourtEtal:DopplerImgV410Tau}, not much is known about the surface differential rotation of these stars, and even less is known about their internal angular and meridional velocities. Theory suggests that the absolute differential rotation in fully convective stars decreases with increasing overall angular velocity due to rotational quenching of the turbulent transport effect that causes the differential rotation \\citep{KuekerEtal:AlphaOmegaRotation,% KitchatinovEtal:LambdaQuenching,% KuekerRuediger:TTauriDiffRot,% KuekerRuediger:TTauriAlpha2}. As in the solar case, the equator is still predicted to rotate more rapidly than the poles. However, some observations of rapidly rotating stars support what is sometimes referred to as ``anti-solar'' differential rotation, where the equator spins less rapidly than the poles \\citep{BarnesEtal:RotationSpotsHKAqr,% KitchatinovRuediger:AntiSolar,% WeberEtal:AntiSolar}. Since differential rotation enters as an important ingredient in dynamo theory, it is important to develop self-consistent models of the large-scale velocity field in fully convective stars.  Magnetic field generation in fully convective stars is also interesting from a dynamo theoretical point of view. With the realization that the magnetic field inside stars might be highly intermittent and concentrated in thin flux tubes, the question of storing such intermittent and strongly buoyant magnetic fields over the course of the 11 year cycle became a growing concern \\citep[e.g.][]{Moreno-Insertis:RiseTimes}. This led to the proposal that dynamos in convective shells (as in the case of the Sun) might operate at or below the bottom of the convection zone. This scenario would not be applicable to fully convective stars, because they lack the overshoot layer where strong flux tubes could be stored. However, it is known that the chromospheric activity does not disappear for later spectral types, i.e.\\ towards fully convective stars \\citep{Vilhu:Activity,VilhuEtal:Activity, BergerEtal:RadioXrayHalphaObservations}. It has therefore been claimed \\citep{DurneyEtal:Generation,% HawleyEtal:ActivityLowMass} that fully convective stars lack large-scale magnetic fields, but can still have small-scale fields generated by non-helical near-surface turbulent dynamo processes.  Attempts to model such small-scale dynamo action \\citep{DorchLudwig:SmallscaleMdwarf} have however led to the conclusion that the photospheric conductivities of M-dwarfs are most probably too low to allow for local small-scale dynamo action This would imply that the observed magnetic activity must be due to dynamo action in deeper layers.  From a kinematic mean-field $\\alpha^2$ dynamo model, \\cite{KuekerRuediger:TTauriAlpha2} predicted that rapidly rotating (Coriolis number of 3 or larger) fully convective stars generate a non-axisymmetric steady magnetic field of quadrupolar symmetry and azimuthal order $m=1$ that looks roughly like a dipole field with the dipole axis lying in the equatorial plane.  Global models of convective dynamos are still in their infancy, even though some tremendous progress was made some 20 years ago when \\cite{Gilman:DynamosII} and \\cite{Glatzmaier:ConvDynamosII} presented the first simulations of dynamos in a spherical shell representing a solar-like convection zone. These models predicted cyclic magnetic fields propagating toward the poles, in contrast to the solar case. The reason for this discrepancy remains a matter of debate even today, when much higher numerical resolution is available. Recent simulations still predict angular velocity to be roughly constant on cylinders, although some simulations show at least a tendency toward solar-like angular velocity contours \\citep{MieschEtal:3dSphericalI,% BrunToomre:TurbulentConvection}. Recent simulations of dynamo action in spherical shells now begin to produce useful models of global turbulent dynamos \\citep{Brun:Interaction,% BrunEtal:GlobalScale}. Meanwhile, such global models have also been applied to core convection \\citep{BrowningEtal:CoreConvection} and to dynamo action in these cases \\citep{BrunEtal:CoreConvDynamo}.  In this paper, we present global dynamo simulations in spheres using a Cartesian grid, i.e.~the sphere is embedded in a cubic box. This may seem to be an unnatural approach to spherical geometry, but it has distinct practical advantages. First, it avoids the coordinate singularity at the center when using spherical coordinates, without invoking expensive transformations from and to spherical harmonics. Second, this approach has proven useful in view of computational simplicity and numerical parallelization efficiency; it has recently been applied by a number of groups to purely hydrodynamic simulations \\citep{PorterEtal:SlabSpheroidal,% FreytagEtal:SpotsBetelgeuse,% WoodwardEtal:3dStellarConvection}, and attempts have already been made to model dynamo action in this approach \\citep{Dorch:BetelgeusePencil}. ",
        "conclusions": "The present work suggests that fully convective stars are capable of generating not only turbulent magnetic fields, but also strong large-scale fields that dominate the magnetic energy spectrum. In most of our models, the large-scale field has a strong quadrupolar component, in contrast to what is expected from mean-field theory for dynamo action in thick shells and in full spheres \\citep{Roberts:KinematicDynamos}. We have so far not seen evidence of magnetic cycles. The resolution of our models is still too low to be able to tell whether this type of magnetic field generation will continue to operate at much larger magnetic Reynolds numbers, but our results disprove the claim that a strong shear layer or a stably stratified core are necessary ingredients for the generation of large-scale magnetic fields. As one would expect in the absence of strong shear layers, the toroidal and poloidal components of the mean magnetic field are roughly comparable.  Another important result concerns the self-consis\\-tently produced differential rotation. In our simulations, the angular velocity shows some tendency to be constant along cylinders, which is plausible for rapidly rotating stars. Whether or not this is realistic is difficult to say. Asteroseismology may in the future be able to reveal the internal angular velocity of stars, but at present the time coverage is still too short and incomplete. There is at least some hope of observing the surface differential rotation, at least of sufficiently rapidly rotating fully convective stars such as T Tauri stars, using surface imaging \\citep{CollierCameronEtal:Astrotomography}. This would be particularly interesting, given that our simulations predict a more slowly rotating equator. This behavior is opposite to that in the solar case. Thus far, theory in terms of the $\\Lambda$ effect \\citep[e.g.][]{RuedigerHollerbach:MagneticUniverse} also tends to produce a faster equator, unless the turbulent motions possess a predominantly radial structure ($u_{r,{\\rm rms}} \\gg u_{\\varphi,{\\rm rms}}$). In our case, however, there is strong meridional circulation, which, due to conservation of angular momentum, causes the outer layers to rotate more slowly; see \\cite{KitchatinovRuediger:AntiSolar}. Again, this result may no longer hold in real stars, because in our model the degree of stratification is far too low and the luminosity too high, so the convective velocities and meridional circulation tend to be exaggerated.  Another reason for the slowly rotating equator could be connected with the outer boundary condition. In connection with geodynamo simulations there are indications that a no-slip outer boundary condition (with respect to a rigidly rotating sphere) tends to produce a more slowly rotating equator \\citep{ChristensenEtal:ParameterStudy}. In the limit of a short damping time, our effective outer boundary condition at $r=R$ should indeed be closer to a no-slip condition than to a free-slip condition. On the other hand, in strongly magnetized stars the coronal magnetic field may enforce a rigidly rotating exterior, and hence produce conditions close to what is represented by our model."
    },
    "0410/astro-ph0410123_arXiv.txt": {
        "abstract": "We performed extensive, multi-wavelength observations of the prototypical symbiotic star Z Andromedae between 2000 and 2003, during a large eruption.  The rise to optical maximum occurred in three distinct stages.  During the first stage, the rise was very similar to an earlier, small outburst which we determined was due to an accretion-disk instability. In the second stage, an optically thick shell of material was ejected, and in the third stage, the shell cleared to reveal a white dwarf whose luminosity was roughly $10^4 L_{\\sun}$.  We suggest that the outburst was powered by an increase in the rate of nuclear burning on the white-dwarf surface, triggered by a sudden burst of accretion.  This outburst thus combined elements of both dwarf novae and classical novae. ",
        "introduction": "Determining the cause of classical symbiotic-star outbursts has been a long-standing challenge.  These eruptions recur too frequently for nova-like thermonuclear runaways, and their peak luminosities appear to be too large for dwarf-nova-like disk instabilities. Although quasi-steady nuclear burning is probably present on the surface of the white dwarfs (WDs) in most symbiotics, classical symbiotic outbursts are also distinctly different from the long-term variability seen in supersoft X-ray sources.  Symbiotic stars are interacting binaries in which material is transferred from an evolved red-giant star to a more compact, hot star, usually a WD \\citep[see, e.g.,][]{kbook,lapalmabook}.  In most symbiotics, the red giant under-fills its Roche lobe, and the mass transfer proceeds via gravitational capture of the red giant's wind. An accretion disk may or may not form \\citep{livio88}.  Radiation from the accreting WD partially ionizes the nebula formed by the red-giant wind.  This ionized nebula gives rise to intense emission lines.  Because typical quiescent-state WD luminosities in symbiotic stars are roughly $10^3\\, L_{\\sun}$ \\citep[e.g.,][]{mur91}, quasi-steady nuclear shell burning is thought to be taking place on the WD surface in the majority of systems \\citep[see also][]{vdh92,sbh01}.   There are at least three types of symbiotic-star outbursts.  Slow novae and recurrent novae appear to be due to thermonuclear runaways on the WD surface.  The nature of the more common classical symbiotic outbursts, however, is not yet known. These eruptions can recur as frequently as every few years, and can brighten in the optical by one to a few magnitudes. \\cite{kbook} found that dwarf-nova-like accretion-disk instabilities in a symbiotic star can only produce outbursts of up to about one magnitude, so outburst modeling efforts have generally focused on expansion of the WD photosphere at constant bolometric luminosity (due to the accretion rate rising above the maximum value for steady-burning), or hydrogen shell flashes. The amount of mass lost during classical symbiotic outbursts has implications for whether symbiotic WDs can accrete enough material to explode as Type Ia supernovae.  Since the mass loss can take the form of a collimated jet, the nature of classical symbiotic outbursts is also linked to the issue of jet formation. ",
        "conclusions": ""
    },
    "0410/astro-ph0410409_arXiv.txt": {
        "abstract": "We have studied the variation of $e^+$ and $\\bar p$ top of the atmosphere spectra due to the parameters uncertainties of the Milky Way geometry, propagation models and cross sections. We have used the B/C data and Galprop code for propagation analysis. We have also derived the uncertainty bands for subFe/Fe ratio, H and He. Finally we have considered a neutralino induced component in the mSUGRA framework. We conclude that even in the case of diffusion convection model for the standard propagation, where secondary spectra gives a good fit of the data, SUSY contribution can not be excluded up to now due to the overall uncertainties. ",
        "introduction": "We have chosen Galprop \\cite{SM1} as a public code for the treatment of propagation of all cosmic rays (CR) together. Our scope has been to determine the total uncertainties in the calculation of $e^+$ and $\\bar p$ top of the atmosphere spectra due to the uncertainties of geometrical and propagation parameters and cross sections.  Here we give very short description of processes included in propagation equation  \\begin{eqnarray} \\frac{\\partial \\psi ({\\mathbf r},p,t) }{\\partial t} &=& q({\\mathbf r}, p)+ \\nabla \\cdot ( D_{xx}\\nabla\\psi - {\\mathbf V_c} \\psi ) + \\frac{d}{dp}\\, p^2 D_{pp} \\frac{d}{dp}\\, \\frac{1}{p^2}\\, \\psi \\nonumber\\\\ &-& \\frac{\\partial}{\\partial p} \\left[\\dot{p} \\psi - \\frac{p}{3} \\, (\\nabla \\cdot {\\mathbf V_c} )\\psi\\right] - \\frac{1}{\\tau_f}\\psi - \\frac{1}{\\tau_r}\\psi\\, , \\label{eq1} \\end{eqnarray}  where $\\psi ({\\mathbf r},p,t)$ is total phase space density. This equation is valid  for all the types of particles. Isotropic diffusion is defined by the coefficient that depends from rigidity (momentum per unit of charge, $\\rho = \\frac{p}{Z}$) $D_{xx} = \\beta D_0(\\rho/\\rho_0)^{\\delta}$, inspired by Kolmogorov spectrum ($\\delta = 1/3$) of weak magnetohydrodynamic turbulence \\cite{SH}. In some models we have used a break in the index $\\delta$ at some rigidity $\\rho_0$, with a value $\\delta_1 = 0$ below the reference rigidity $\\rho_0$. The convection velocity field ${\\mathbf V_c}$, that corresponds to the Galactic wind, has a cylindrical symmetry and its z-component is the only one different from zero. It increases linearly with the distance $z$ from the galactic plane, in agreement with magnetohydrodynamical models \\cite{Z}. Reacceleration is determined by the diffusion coefficient for the impulse space $D_{pp} $ that is a function of the corresponding configuration space diffusion coefficient and of the Alfven velocity $V_{A}$  in the framework of quasi-linear MHD theory \\cite{FSPB}. Of course, Alfven velocity and convection velocity gradient in the Milky Way for reacceleration and convection terms are unknown parameters of propagation (there are no other sources of information from which we could extract them, except the spectra of cosmic rays) and their possible range will be constrained by the analysis of fits of suitable data. The same procedure is valid for constraining the height of the Galactic halo and the other unknown parameters. This will be analyzed further in order to obtain all the possible spectra of antiprotons and positrons using the sets of the constrained parameters. This procedure was already used in \\cite{FN} for another propagation code, even for the antiprotons from neutralino annihilations. In the same reference can be found also the review of the propagation models in the Galaxy with an extended colection of historical and recent references. Injected spectra of all primary nuclei are power law in impulse $dq(p)/dp \\propto p^{-\\gamma}$. This power law approximation has been shown to be allowed in the framework of diffusive shock acceleration models, as well as a small break in the injection index $\\gamma$ \\cite{BO, SMx, E}. Source term $q({\\mathbf r}, p)$ for secondaries contains cross sections for their production from progenitors on H and He targets. The last two terms in equation \\ref{eq1} are loss terms with characteristic times for fragmentation and radioactive decay. Propagation equation is solved numerically, using the Crank-Nicholson algorithm.  The heliospheric modulation in the vicinity of the Earth has to be taken in account. We have used a model in which transport equation (that describes diffusion processes in the heliosphere and includes effects of heliospheric magnetic field and Solar wind) is solved in the force field approximation~\\cite{GAP}. In this case Solar modulation is a function of just a single parameter that describes the strength of the modulation. All the dynamical processes are simulated with relatively simple changing of the interstellar spectra during the propagation inside the heliosphere, described by the formula  \\begin{eqnarray} \\label{eq2} \\frac{ \\Phi^{toa} (E^{toa})}{\\Phi^{is}(E^{is})} =  (\\frac{p^{toa}}{p^{is}})^{2} , \\\\ E^{is} -  E^{toa} = |Ze| \\phi \\end{eqnarray}  where $E$ and $p$ are energies and impulses of the interstellar and top of the atmosphere fluxes and $\\phi$ is the unique parameter that determines the solar modulation. ",
        "conclusions": "For $e^+$ in DR model, even when the uncertainties are included, the curve of the minimal $e^+$ production still remains above the experimental results. Breaking the primary spectra gently improves just the low energy part of the spectra below the maximum, but even the minimal predictions still remain bigger than the experimental results around the maximum as well as below it. On the other side, this slightly changes the best B/C fit (fig.~\\ref{figBC}). This is also not sufficient to match protons and helium, that are still overestimated. In the end, $e^-$ remain largely overproduced at low energies, even more than without the break. Uncertainty bands of $\\bar p$ in DR models touch the experimental data from below. In this case the fit can be improved easily with any primary component that is coming from eventual neutralino annihilation or some other exotic contribution. For DC model all the results are in good agreement with the data, with only some problems with B/C data.  As a possibility for the future considerations we would like to emphasize that it is natural to expect models that include both of the processes -- convection and reacceleration, and this perhaps can be a key to solve simultaneously problems with positrons, electrons and B/C. But, we have not yet found such a kind of models. In any case, further measurements of $\\bar p$ and $e^+$ spectra, primary to secondary CR ratios and Solar effects on them, as well as the precise determination of important nuclear cross sections seem to be crucial for answering the question of the models quality.  In the framework of DC model exotic contributions remain possible at high energies (E $>$ 20 GeV), and also they are not excluded at lower energies due to the relatively large uncertainties (if the Standard Model productions are smaller than the best fits). In the case of DR model background it is possible to improve the $\\bar{p}$ data fit adding a neutralino induced component for a large part of the parameter space. Finally we have shown how a possible neutralino induced component in a clumpy halo scenario could give a contribution to the total $\\bar{p}$ flux, in the case of DC model background."
    },
    "0410/astro-ph0410586_arXiv.txt": {
        "abstract": "{} {The Strouhal number (St), which is a nondimensional measure of the correlation time, is determined from numerical models of convection. The Strouhal number arises in the mean-field theories of angular momentum transport and magnetic field generation, where its value determines the validity of certain widely used approximations, such as the first order smoothing (hereafter FOSA). More specifically, the relevant transport coefficients can be calculated by means of a cumulative series expansion if St $<$ St$_{\\rm crit} \\approx 1$.} {We define the Strouhal number as the ratio of the correlation and turnover times, which we determine separately, the former from the autocorrelation of velocity, and the latter by following test particles embedded in the flow.} {We find that the Strouhal numbers are, generally, of the order of 0.1 to 0.4 which is close to the critical value above which deviations from FOSA become significant. Increasing the rotational influence tends to shorten both timescales in such a manner that St decreases. However, we do not find a clear trend as function of the Rayleigh number for the parameter range explored in the present study.} {} ",
        "introduction": "The mean-field theories of angular momentum transport (e.g. R\\\"udiger \\cite{Rudiger80}, \\cite{Rudiger89}) and hydromagnetic dynamos (e.g. Steenbeck \\& Krause \\cite{SteenKrau69}; Krause \\& R\\\"adler \\cite{KraRad80}) require knowledge of the Reynolds stresses and the mean electromotive force, respectively. Direct analytical calculation of these quantities for astrophysical purposes is not possible at present due to the lack of an established theory of turbulence. Numerical calculations are not much better off since in most astrophysical systems the computational resources needed in order to resolve all physically relevant scales are several orders of magnitude larger than currently available.  In mean-field theories these problems are circumvented by relating the Reynolds stresses and electromotive force to the mean quantities (the rotation vector $\\vec{\\Omega}$ and the mean magnetic field $\\langle \\vec{B} \\rangle$, respectively) and their gradients by means of a cumulative series expansion (see van Kampen \\cite{vanKampen1974a},\\cite{vanKampen1974b}). In the best known and most often used approach, the first order smoothing approximation (FOSA), only the first terms of these expansions are taken into account. This approach can be shown to be valid if either the relevant Reynolds number, ${\\rm Re} = ul/\\nu\\;, {\\rm Rm} = ul/\\eta$, or the Strouhal number \\begin{equation} {\\rm St} = \\frac{|(\\vec{u} \\cdot \\nabla)\\vec{u}|}{|\\partial \\vec{u}/\\partial t|} \\approx u \\frac{\\tau_{\\rm c}}{l_{\\rm c}} \\propto \\frac{\\tau_{\\rm c}}{t_{\\rm to}}\\;, \\label{equ:st} \\end{equation} is sufficiently small. Above, $u$ and $l$ are the typical velocity and length scales, and $\\tau_{\\rm c}$ and $l_{\\rm c}$ are the correlation time and length of the turbulence, respectively. We show in Sect.\\,\\ref{subsec:turnovert} that $1/t_{\\rm to} \\propto u/l_{\\rm c}$. Only the case of small St is relevant in astrophysical circumstances, where typically Re, Rm $\\gg$ 1. Hence the condition St $\\ll$ 1 should be fulfilled if FOSA is to be valid in stellar convection zones. If this is not the case then FOSA is likely to be too crude a truncation but, as long as St is smaller than the critical value for convergence it is possible to construct higher-order mean-field theories by including more terms of the cumulant expansion. However, a universal critical value of St for convergence of the cumulant expansion does not exist because this depends on the geometry of the problem and on the flow itself. In general though, the critical value is thought to be less than unity. For simple turbulence models it is possible to compute it; Nicklaus \\& Stix (\\cite{NickStix88}, hereafter NS88) obtained $\\mbox{St}_{\\rm crit}=1$ for their model. Furthermore, the results of NS88 show that the higher order effects remain small in comparison to the FOSA result when St $\\lesssim 0.5$.  It is well-known that the requirement for a small Reynolds number is not satisfied in stellar environments. However, the question of the Strouhal number is not settled. On account of the observations of the solar surface granulation, one can estimate the correlation and turnover times to be roughly equal, indicating that St $\\approx 1$ at the solar surface (see e.g. Chapter 6 of Stix \\cite{Stix02}). Similar values can also be estimated for supergranulation for which typical numbers are $u \\approx 100$~m~s$^{-1}$, $l \\approx 10^7$m, and $\\tau \\approx 10^5$~s. However, even for granulation the precise value has, to our knowledge, not been established, and nothing is known about St in the deeper layers. Furthermore, recent results from forced turbulence calculations indicate that if the higher order correlations in the equations of the passive scalar flux (Brandenburg et al. \\cite{Brandea04}, hereafter BKM) and electromotive force (Brandenburg \\& Subramanian \\cite{BranSubr05}) are taken into account via the so-called minimal $\\tau$-approximation (Blackman \\& Field \\cite{BlackField2002}, \\cite{BlackField2003}), the Strouhal number can be seen to substantially exceed unity and roughly equal to unity, respectively.  Although forced turbulence is rather different in comparison to convection, the aforementioned studies still raise the question whether the results of the convection calculations can be interpreted within the framework of the standard mean-field theory as has been done in numerous studies during recent years (e.g. Brandenburg et al. \\cite{Brandea90}; Pulkkinen et al. \\cite{Pulkkiea93}; Ossendrijver et al. \\cite{Ossea01}, \\cite{Ossea02}; K\\\"apyl\\\"a et al. \\cite{Kapyea04}; R\\\"udiger et al. \\cite{Ruedigerea2005}). Motivated by the unknown status of the Strouhal number for convection and the previous studies on the subject in different contexts, we set out to calculate St from numerical calculations of convection. In order to do this, we calculate the correlation time from the autocorrelation of velocity and determine the turnover time by following test particles embedded into the flow and define the Strouhal number as the ratio of the two.  The remainder of the paper is organised as follows: Sect.\\,\\ref{sec:model} summarises briefly the numerical model used and in Sect.\\,\\ref{sec:results} the results of the study are discussed. Finally, Sect.\\,\\ref{sec:conclu} gives the conclusions. ",
        "conclusions": "\\label{sec:conclu} We estimate the nondimensional measure of the correlation time, the Strouhal number, from numerical models of convection. We calculate the correlation and turnover times separately from the velocity autocorrelation function and the trajectories of embedded test particles, respectively, and define the Strouhal number as the ratio of the two.  The Strouhal number arises in the mean-field theories of angular momentum transport and hydromagnetic dynamos where its value determines the validity of certain widely used approximations, such as the first order smoothing. These approximations are based on a cumulative series expansion of the relevant turbulent correlation, e.g. the electromotive force in the dynamo theory. Essentially, the higher order terms in this expansion are proportional to the Strouhal number. Thus the value of St determines whether or not the expansion converges.  The main results can be summarised as follows: \\begin{itemize} \\item We find that the correlation time does not depend on the Rayleigh number for the parameter range explored in the present study. Note that the results for $\\tau_{\\rm c}$ presented here differ significantly from those of K\\\"apyl\\\"a et al. (\\cite{Kapyea05}) due to an error in the analysis in the latter study.  \\item As function of rotation, $\\tau_{\\rm c}$ decreases by a factor of roughly three when the Coriolis number is increased from 0 to 10. At the same time, the turnover time decreases only by about $20\\%$, leading to an approximate relation ${\\rm St} \\propto {\\rm Co}^{-0.32}$ for moderate and rapid rotation.  \\item A noteworthy fact is that if one takes the usual estimate of dividing the characteristic length by the characteristic velocity, the turnover time increases as function of rotation due to the overall decreasing velocities. However, the test particle data gives an opposite result. The discrepancy is due to the fact that even though velocities decrease in general, the spatial scale of convection decreases even more. Thus it is imperative that the definition of St takes into account the spatial scale of the motions, a fact which was also noted by NS88. \\end{itemize}  Our inability to find any dependence on the Rayleigh number and the relatively high values ($0.1 \\ldots 0.4$) of the Strouhal number raise the question of the validity of FOSA. In the study of NS88, although not from a directly comparable context, the higher order effects became visible when St was of similar magnitude. At the moment, we cannot draw a firm conclusion whether or not it is safe to use FOSA or not. We think that the next step in the direction of determining this would be to develop mean-field expressions for the transport coefficients that are directly applicable to the present numerical calculations. Furthermore, these expressions should be derived using FOSA and preferably also with some higher order expansion in order to determine which of them gives a better description of the numerical result, and, more importantly, whether the whole concept of the cumulative expansion is applicable. This study, however, does not fit in the scope of the present paper."
    },
    "0410/astro-ph0410065_arXiv.txt": {
        "abstract": "Over the last few years, rapid progress has been made in distance measurements for nearby galaxies based on the magnitude of the tip of red giant branch stars. Current CCD surveys with HST and large ground- based telescopes bring  $\\sim$10\\%-accurate distances for roughly a hundred galaxies within 5 Mpc. The new data on distances to galaxies situated in (and around) the nearest groups: the Local Group, M81 group, CenA/M83 group, IC342/Maffei group, Sculptor filament, and Canes Venatici cloud allowed us to determine their total mass from the radius of the zero- velocity surface, $R_0$, which separates a group as bound against the homogeneous cosmic expansion. The values of $R_0$ for the virialized groups turn out to be close each other, in the range of 0.9 -- 1.3 Mpc. As a result, the total masses of the groups are close to each other, too, yielding total mass-to-blue luminosity ratios of 10 -- 40 $M_{\\sun}/L_{\\sun}$. The new total mass estimates are 3 -- 5 times lower than old virial mass estimates of these groups. Because about half of galaxies in the Local Volume belong to such loose groups, the revision of the amount of dark matter (DM) leads to a low local density of matter, $\\Omega_m \\simeq 0.04$, which is comparable with the global baryonic fraction $\\Omega_b$, but much lower than the global density of matter, $\\Omega_m = 0.27$.  To remove the discrepancy between the global and local quantities of $\\Omega_m$, we assume the existence of two different DM components: 1) compact dark halos around individual galaxies and 2) a non-baryonic dark matter ``ocean'' with $\\Omega_{dm1} \\simeq 0.07$ and $\\Omega_{dm2} \\simeq 0.20$, respectively. ",
        "introduction": "In rich clusters of galaxies, virial mass estimates agree well with independent determinations of the cluster mass made from the X-ray flux of hot intracluster gas, and from weak gravitational lensing effects. A typical ratio of the total mass-to-blue luminosity for rich clusters, $M_t/L_B \\sim 250 M_{\\sun}/L_{\\sun}$, extrapolated over the whole volume of the universe, yields the mean density of matter $\\Omega_m \\sim$0.25, in the excellent concordance with parameters of the standard $\\Lambda$CDM model: $\\Omega_m$ = 0.27, $\\Omega_{\\Lambda}$ = 0.73 (Spergel et al. 2003). However, about 85\\% of galaxies are situated outside the rich clusters. Roughly a half of them belong to groups of different size and population, while the remaining half are scattered in diffuse (unvirialized) ``clouds'' and ``filaments'' usually called the ``field''. Until recently, application of the virial theorem to galaxy groups remained the only way to trace the dark matter distribution on scales of 0.1 -- 1 Mpc. Measurements of the total mass of individual groups via their X-ray flux or weak gravitational lensing have not lead yet to distinct results.  In the case of our Local Group (LG), Lynden-Bell (1981) and Sandage (1986) proposed to determine its total mass using a method, which is based on a measurement of the radius of the ``zero-velocity surface'', $R_o$, where the overdensity of the group has halted expansion and infall is about to commence.  Under the assumption of spherical symmetry the total mass of a group can be expressed in terms of the radius $R_o$, the age of the Universe $T_o$ , and the Gravitational constant $G$ as $$M_t = (\\pi^2/8G)\\times R_o^3 \\times {T}_o^{-2}.$$ Application of this method to the LG as well to other groups requires accurate distances and radial velocities for all galaxies surrounding a group. Such data is now accessible for several nearby groups by using the luminosity of stars at the tip of the red giant branch (TRGB) as standard candles.  Over the last 5 years, searches for additional nearby dwarf galaxies, made on the POSS-II \\& ESO/SERC plates by Karachentseva \\& Karachentsev (1998; 2000), and also results of ``blind'' HI surveys of the southern sky, and the Zone of Avoidance by the HIPASS team (Kilborn et al. 2002), as well as other wide-field sky surveys, lead to a doubling of the population of known galaxies in the Local volume. Basic observational data on 450 Local volume galaxies with distances less than 10 Mpc are presented in the Catalog of Neighboring Galaxies (Karachentsev et al. 2004).  The new quantity and quality of observational data on nearby galaxy distances and radial velocities allowed us to determine virial and total masses of four of the nearest complexes: Milky Way+M31 (Karachentsev et al. 2002c), M81+NGC2403 (Karachentsev et al. 2002a), CenA+M83 (Karachentsev et al. 2002b), and IC342+Maffei (Karachentsev et al. 2003d), as well as of two other nearby scattered galaxy systems:  M96 = the Canes Venatici I cloud (Karachentsev et al. 2003a) and NGC253 = the Sculptor filament (Karachentsev et al. 2003c). Below we discuss the results of mass determination for all six nearest groups made via internal and external galaxy motions. ",
        "conclusions": ""
    },
    "0410/astro-ph0410253_arXiv.txt": {
        "abstract": "We perform a detailed analysis of the optical gravitational lens \\ER\\ imaged with WFPC2 on the Hubble Space Telescope.  Using software specifically designed for the analysis of resolved gravitational lens systems, we focus on how the image alone can constrain the mass distribution in the lens galaxy.  We find the data are of sufficient quality to strongly constrain the lens model with no \\emph{a priori} assumptions about the source.  Using a variety of mass models, we find statistically acceptable results for elliptical isothermal-like models with an Einstein radius of 1.17\\scnd\\ An elliptical power-law model ($\\Sigma \\propto R^{- \\beta}$) for the surface mass density favours a slope slightly steeper than isothermal with $\\beta = 1.08 \\pm 0.03$. Other models including a constant M/L, pure NFW halo and (surprisingly) an isothermal sphere with external shear are ruled out by the data. We find the galaxy light profile can only be fit with a S\\'ersic plus point source model.  The resulting total M/L$_B$ contained within the images is $4.7 h_{65} \\pm0.3$.  In addition, we find the luminous matter is aligned with the total mass distribution within a few degrees. This is the first time a resolved optical gravitational lens image has been quantitatively reproduced using a non-parametric source.  The source, reconstructed by the software, is revealed to have two bright regions, with an unresolved component inside the caustic and a resolved component straddling a fold caustic. The angular size of the entire source is $\\sim 0.1$\\scmd and its (unlensed) Lyman-$\\alpha$ flux is $3 \\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$. ",
        "introduction": "\\label{introduction}  The discovery \\citep{1988Natur.333..537H,Langston1989AJ.....97.1283L} and subsequent modelling \\citep{Kochanek1992ApJ...401..461K,1995ApJ...445..559K} of Einstein ring gravitational lenses at radio wavelengths raised the prospect of obtaining detailed quantitative information for both the lensed sources and the form of the mass distribution in the intervening lensing galaxies. \\citet{1995ApJ...445..559K} compared mass models for the lensed radio lobe MG1654+134 using \\textsc{LensCLEAN} \\citep{1992ApJ...401..461K} in image space. It was subsequently shown that \\textsc{LensCLEAN}-ing radio data in image space can introduce biases, hence visibility-based \\textsc{LensCLEAN} was introduced \\citep{1996ApJ...464..556E,2004MNRAS.349....1W}. \\citet{2004MNRAS.349...14W} recently used this method to measure $H_0$ in B0218+357.  Recent work at optical/near-IR wavelengths \\citep{2000ApJ...535..692K,2001ApJ...547...50K} has concentrated on what can be learnt using infrared images of Einstein rings. \\citet{1992MNRAS.259P..31M} pointed out that there should exist a much larger number of Einstein rings detectable at optical wavelengths, and systematic surveys for such systems are now underway \\citep{2000iutd.conf...94H}. \\ER, the first known example of a galaxy lensing another normal galaxy, was the first Einstein ring to be discovered at optical wavelengths \\citep{Warren1996MN}. The \\ER \\ system, together with the recently discovered lens 1RXS J113155.4-123155 \\citep{2003A&A...406L..43S}, are the only Einstein ring systems with good S/N Hubble Space Telescope (HST) images at optical wavelengths. Modelling of the system, exploiting the extended surface brightness distribution of the lensed source, would provide details of the source morphology and allow some discrimination between models for the mass distribution in the intervening lensing galaxy. Compilation of a large sample of optical Einstein rings could provide information on the luminosity function and morphologies of star--forming galaxies at redshifts $z=3-5$ whose (unlensed) magnitudes remain too faint for direct study using even 8m--class telescopes. Similarly, statistical properties derived from a large sample of such lensed systems, will provide new information on the still poorly constrained distribution of mass in massive early--type galaxies, which form the majority of the deflector population.  As a first step toward achieving these goals, this paper presents an analysis of the HST\\footnote{Based on observations made with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with program \\# 6560} imaging observations of \\ER.  This system has been studied by \\citet[hereafter `KT03']{2003ApJ...583..606K}, who combined dynamical information (the stellar velocity dispersion profile), with gravitational lens information (the angular size of the Einstein ring), to draw conclusions on the mass profiles of both the visible and dark matter. The HST image was used simply to determine the total mass within the ring, and to measure the lens galaxy light profile. In our complementary analysis, we concentrate on extracting as much information on the mass profile as possible from the image alone. We use software based on the \\textsc{LensMEM} method \\citep[hereafter `WKN96']{1996ApJ...465...64W} to perform a gravitational--lens inversion of the complete surface brightness distribution around the ring. This allows us to draw conclusions on the radial distribution of the (total) mass, and also produces new insight into the source morphology. An advantage of the optical CCD image over radio images is that the counts in the pixels are independent, so that an unambiguous goodness--of--fit statistic may be measured. To analyse a radio image properly requires the covariance matrix for the image pixels, or to work directly with the visibilities \\citep{1996ApJ...464..556E,2004MNRAS.349....1W}.  Section~\\ref{observations} describes the HST WFPC2 observations, the data reduction steps, the determination of the noise characteristics of the final image of \\ER, and the accurate measurement of the surface brightness profile of the lensing galaxy.  Section~\\ref{gravlens} presents details of the lens modelling algorithm.  Seven different mass models for the lens are considered. Four of the mass models provide satisfactory fits to the data, whereas three are rejected. A discussion of the results, and comparison with the analysis of KT03, is given in Section 4, before the principal conclusions are summarised in the final section.  Throughout this paper, we assume a $\\Lambda$CDM cosmology with $\\Omega_{\\Lambda}=0.7$, $\\Omega_{m}=0.3$ and $H_0 = 65 \\ \\hubunits$ (for consistency with KT03) unless otherwise stated. ",
        "conclusions": "We have performed a detailed analysis of \\ER\\ using data from WFPC2 on board HST. The data show that the lensed image has four distinct bright regions and two regions of extended emission. We have modelled the image with six different lens models using sophisticated software based on the \\textsc{LensMEM} algorithm. The software generates a best-fitting image and reconstructs the source brightness profile using a non-parametric source model.  We have tested the ability of the data to distinguish between a variety of lens models including isothermal, power-law, constant M/L and NFW mass distributions.  We find that the `canonical' SIE and PIEP lens models fit the data well and that a power-law model ($\\Sigma \\propto R^{-\\beta}$) favours $\\beta = 1.08 \\pm 0.03$, slightly steeper than isothermal.  In addition we find that we can fit the data with a mass profile based on the NFW profile, but only when the scale length is too small and central density too high for a dark matter halo expected in a large elliptical galaxy. Conversely, we find that a mass model based on constant M/L or a realistically-sized NFW halo cannot fit the data.  We also find that the simple SIS$+\\gamma$ lens model does not fit the data despite being an isothermal model.  The data are good enough to distinguish between elliptical lens models and the SIS$+\\gamma$ which is a first-order approximation of an elliptical model.  The difference appears to be in the shape of the critical curve between the elliptical models and external shear model.  We have discovered also that the source is actually a double. This has caused problems for models which assume the bright regions in the image come from a single source.  Our results are somewhat different to those of KT03 who found a larger Einstein Radius by $15\\%$. However, their model assumed a single source and included a substantial shear component.  The mass enclosed by KT03's model within a $1.17 \\scmd$ aperture is consistent with our model. Our models show that the image locations and brightnesses can be explained neatly with the simple SIE/PIEP mass models and the double source. We calculate the total M/L$_B$, inside 1.17\\scmd, to be $4.7 h_{65} \\pm0.3$ compared to $5.4\\pm0.8$, at a larger radius of 1.34\\scmd, in KT03.  We have qualitatively evaluated the behaviour of the lens models around the location of the lensed images by plotting contours of constant magnification. We found that the SIE and PIEP models are indistinguishable for these data and that any elliptical mass distribution with an isothermal-like density profile around the images is likely to be able to fit the data. We found that the lens models which could not fit the data produce images with incorrect magnification and/or location.  Finally, we note that the lens model requires no external shear to explain the data. The well constrained mass distribution, lack of external shear and isolation of the lens galaxy all suggest that \\ER\\ is an ideal candidate for further work such as monitoring for supernova events in the source, or making a detailed study of the properties of the galaxy's dark halo.  \\newcommand{\\nat}{Nat} \\newcommand{\\mnras}{MNRAS} \\newcommand{\\aj}{AJ} \\newcommand{\\pasp}{PASP} \\newcommand{\\aap}{A\\&A} \\newcommand{\\apjl}{ApJ}"
    },
    "0410/hep-ph0410136_arXiv.txt": {
        "abstract": "\\noindent Extraterrestrial neutrinos can initiate deeply developing air showers, and those that traverse the atmosphere unscathed may produce cascades in the ice or water. Up to now, no such events have been observed. This can be translated into upper limits on the diffuse neutrino flux. On the other hand, the observation of cosmic rays with primary energies~$> 10^{10}$~GeV suggests that there is a guaranteed flux of cosmogenic neutrinos, arising from the decay of charged pions (and their muon daughters) produced in proton interactions with the cosmic microwave background. In this work, armed with these cosmogenic neutrinos and the increased exposure of neutrino telescopes we bring up-to-date model-independent upper bounds on the neutrino-nucleon inelastic cross section. Uncertainties in the cosmogenic neutrino flux are discussed and taken into account in our analysis. The prospects for improving these bounds with the Pierre Auger Observatory are also estimated. The unprecedented statistics to be collected by this experiment in 6~yr of operation will probe the neutrino-nucleon inelastic cross section at the level of Standard Model predictions. ",
        "introduction": "Introduction}  Ultrahigh energy cosmic neutrinos may reveal aspects of nature hidden to us so far. They can point back to very distant sources, resolving the origin of the highest energy cosmic rays and the underlying acceleration mechanism. Upon their arrival at Earth, they interact with nucleons at centre-of-mass energies around several hundreds of TeV, probing the energy regime far beyond the reach of terrestrial colliders.  The detection of ultrahigh energy cosmic neutrinos is a very challenging task due to their weak interactions. Although the neutrino-nucleon inelastic  cross section increases with energy, the cosmic neutrino flux falls even more steeply with energy. Large-scale projects and novel techniques are being deployed and proposed. The Pierre Auger Observatory (PAO)~\\cite{Abraham}, followed by IceCube~\\cite{Ahrens:2002dv}, and possibly EUSO~\\cite{Catalano:2001mm} and OWL~\\cite{Stecker:2004wt}, will reach a sensitivity of the level of theoretical predictions for the cosmic neutrino fluxes. Moreover, before the next generation experiments lead us into the exciting era, the non-observation of neutrino-induced events reported by the Fly's Eye~\\cite{Baltrusaitis:mt}, the AGASA~\\cite{Inoue:cn,Yoshida:2001icrc} and the RICE~\\cite{Kravchenko:2002mm,Kravchenko:2003tc} collaborations can be used to set upper limits on the diffuse neutrino flux. Additionally, by exploiting a certain prediction for the neutrino flux, the search results can also be turned into upper bounds on the neutrino-nucleon inelastic cross section~\\cite{Berezinsky:1974kz,Anchordoqui:2002vb,Tu:2004ms}.  Ultrahigh energy cosmic neutrinos from diverse sources are predicted, and their existence is strongly supported by the observation of the ultrahigh energy cosmic rays (UHECRs). Among all, the so-called cosmogenic neutrinos \\cite{Berezinsky:1969} are almost guaranteed to exist. They originate from the decay of charged pions produced in the interactions of protons with the cosmic microwave background (CMB). There are still  some uncertainties in the estimate of the cosmogenic neutrino flux.  This is due to our poor knowledge of the nature and the origin of the UHECRs. Possible ranges for the size of the cosmogenic neutrino flux have been investigated elsewhere~\\cite{Fodor:2003ph} assuming that the cosmic ray spectrum beyond $10^{8}$~GeV is dominated by extragalactic protons with an isotropic distribution of sources. In this work, armed with these cosmogenic neutrinos and the increased exposure of neutrino detectors, we show that the neutrino-nucleon inelastic cross section is tighter constrained than previously noted. Uncertainties in the cosmogenic neutrino flux are discussed and taken into account in our analysis. The prospects for improving these bounds with the PAO are also estimated.   The structure of the paper is the following. In Sec.~\\ref{formulae} we examine acceptances for neutrino detection. In particular, we compute the effective apertures of AGASA and RICE as examples of ground arrays and under-ice neutrino telescopes, respectively. We also estimate the aperture of the PAO in order to investigate future sensitivity to physics beyond the Standard Model (SM). In Sec.~\\ref{flux} we present an overview of cosmogenic neutrino fluxes. In Sec.~\\ref{table} we derive model-independent upper bounds on the neutrino-nucleon inelastic cross section from the search results reported by the AGASA and the RICE collaborations. In Sec.~\\ref{PAO}, the prospects for improving these bounds with the PAO are estimated. Section~\\ref{conclusions} comprises our conclusions. ",
        "conclusions": "Conclusions}  The search for ultrahigh energy cosmic neutrinos is entering an exciting era. Recently, a wide array of projects have been initiated to detect neutrinos by searching for low-altitude quasi-horizontal showers, cascades in the Antarctic ice-cap, and radio emission from neutrino-induced showers. Some of these experiments, like PAO and RICE, have been taken data, and others, like the km$^3$ IceCube telescope, are under construction.  Although such a high energy cosmic neutrino flux has not been observed to date, it has long been known that there should be a guaranteed cosmogenic flux resulting from the interactions of cosmic protons with the CMB. Therefore, the non-observation of deeply developing showers or cascades in ice reported by neutrino-detection-experiments in conjunction with these cosmogenic neutrinos can be used to constrain the behaviour of the neutrino-nucleon inelastic cross section in the energy regime far beyond the reach of terrestrial colliders.  In this work we derived model-independent upper bounds on the neutrino-nucleon inelastic cross section from existing search results reported by the AGASA and the RICE collaborations. The bounds apply for neutrino-nucleon cross sections smaller than 0.5~mb (AGASA), or 1~mb (RICE). Interestingly, the search result reported by the RICE Collaboration improves the upper bounds derived from the non-observation of deeply developing showers at AGASA by more than one order of magnitude. Therefore, in the presence of additional neutrino-emitting-sources, there is only  little room for new physics contributions to the inelastic cross section.  We have also estimated the sensitivity of PAO to neutrino interactions. This hybrid detector will facilitate powerful air shower reconstruction methods and control of the systematics errors which have plagued cosmic ray experiments to date. Through a comparison  of deeply developing showers and Earth-skimming events, in 6 yr of operation the observatory will be able to probe neutrino interactions at the level of SM predictions, providing a final verdict on the rise of the neutrino-nucleon inelastic cross section at ultrahigh energies."
    },
    "0410/hep-th0410034_arXiv.txt": {
        "abstract": "We present significant evidence in favour of the holographic conjecture that ``4D black holes localized on the brane found by solving the classical bulk equations in $AdS_5$ are quantum corrected black holes and not classical ones''. The crucial test is the calculation of the quantum correction to the Newtonian potential based on a numerical computation of $\\langle T^a_{\\ b}\\rangle $ in Schwarzschild spacetime for matter fields in the zero temperature Boulware vacuum state. For the case of the conformally invariant scalar field the leading order term is found to be $M/45\\pi r^3$. This result is equivalent to the result which was previously obtained in the weak-field approximation using Feynman diagrams and which has been shown to be equivalent, via the AdS/CFT duality, to the analogous calculation in Randall-Sundrum braneworlds. This asymptotic behavior was not captured in the analytical approximations for $\\langle T^a_{\\ b}\\rangle $ proposed in the literature. The 4D backreaction equations are then used to make a prediction about the existence and the possible spacetime structure of macroscopic static braneworld black holes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410290_arXiv.txt": {
        "abstract": "We emphasize the importance of observations of young supernovae in wide radio band. We argue on the basis of observational results that only high or only low frequency data is not sufficient to get full physical picture of the shocked plasma. In the SN 1993J composite spectrum obtained with Very Large Array (VLA) and Giant Metrewave Radio Telescope (GMRT), around day 3200, we see observational evidence of synchrotron cooling, which leads us to the direct determination of the magnetic field independent of the equipartition assumption, as well as the relative strengths of the magnetic field and relativistic particles energy densities. The GMRT low frequency SN 1993J light curves suggests the modification in the radio emission models developed on the basis of VLA data alone. The composite radio spectrum of SN 2003bg on day 350 obtained with GMRT plus VLA strongly supports the internal synchrotron self absorption as the dominant absorption mechanism. ",
        "introduction": "Radio emission from supernovae is argued to be due to the synchrotron emission from the forward shocked shell due to the relativistic electrons in presence of magnetic field. The radio emission is initially absorbed from the external medium (free-free absorption i.e. FFA) or through synchrotron self absorption (SSA), i.e. an internal process.   The most critical parameter which affects the synchrotron emission is the magnetic field. The magnetic fields in a few supernovae have been estimated indirectly by assuming equipartition between relativistic energy density and the magnetic energy density. The magnetic field in the shocked plasma in supernovae is enhanced due to hydrodynamic instabilities in the plasma. While it is plausible, there is no convincing argument for the equipartition assumption. In many classical radio sources, such as supernova remnants (SNRs) like the Crab or Cassiopeia A, or in luminous radio galaxies, the radio spectral index is found to steepen at high frequencies (see e.g. Kardashev 1962). % This is due to the so called synchrotron aging of the source, as during the lifetime of the source, electrons with high enough energies in a homogeneous magnetic field will be depleted due to efficient synchrotron radiation compared with the ones with lower energies. An observation of a synchrotron break can yield a measurement of the magnetic field  {\\it independent of the equipartition argument} if the {\\it age} of the source {\\it is known}.  We emphasize here  the need for broad-band  radio observations of SNe to address the above issues. Combining radio data from a high frequency VLA and low frequency GMRT can offer such opportunities. We discuss two supernovae in this context - an eleven years old type IIb supernova SN 1993J and a one year old type Ic supernova SN 2003bg. ",
        "conclusions": "In right panel of Fig. 1 we show a comparison of SSA model (with a single optically-thin power-law index) fitted only to the low frequency data (0.22 GHz to 1.4 GHz) versus such a model fit obtained with only the higher frequency data (1.4 GHz to 22.5 GHz). This comparison shows that while the model fitted only to the low frequencies over-predicts the flux density at high frequencies, the model fitted only to high frequencies on the other hand fails to account for both synchrotron cooling break and seriously under-predicts the low frequency flux densities. The comparison underscores the importance of broad band observations for determining the physical processes taking place in the supernova.  The particle energy density is far below than the magnetic energy density (by a factor of 1/10000) for SN 1993J, although there are indications that in SNe like SN 1998bw/GRB980425, there may exist this equipartition (Kulkarni et al. 1998). Future studies of the equipartition factor for SN 1993J may indicate for the first time how particle acceleration efficiency in strongly turbulent magnetized plasma evolves with time in the large magnetic-Reynolds and Reynolds numbers limit. Light curves based on high frequency FFA models extrapolated to low frequencies overpredict the fluxes at low frequencies. Some extra opacity is needed to incorporate the difference. We added an extra opacity due to SSA which also could not account for the required absorption (Fig. 2). This suggests the low frequency opacity in SN 1993J is not the simple extrapolation of high frequency opacity and it is likely that hitherto unaccounted  absorption mechanisms are at work at low frequencies.  There could be a break in the radio spectrum of SN 2003bg. We are planning simultaneous observations with GMRT along with VLA and ATCA. This will provide us the spectrum from 0.2 GHz to 80 GHz and will establish whether the break is real or due to data artifact. We also could discard the homogeneous free-free absorption model over the SSA model from the low-frequency optically thin part of the spectrum.   We thank Dr. S. Bhatnagar for crucial help in the analysis of early GMRT data for SN 1993J. We thank K. Weiler, C. Stockdale and their collaboration for kindly providing the VLA data for SN 1993J at high frequencies. We also thank A. Soderberg and S. Kulkarni for observing SN 2003bg on our request and providing us the results of observations. We thank the staff of the GMRT, NCRA (TIFR) that made these observations possible. Poonam Chandra is a Sarojini Damodar International Fellow."
    },
    "0410/astro-ph0410259_arXiv.txt": {
        "abstract": "We  report on  a  300  ks BeppoSAX  (0.12--200  keV) observation  of Circinus X--1 (Cir X--1) at phases between 0.62 and 0.84 and on a 90 ks BeppoSAX observation of Cir  X--1 at phases 0.11--0.16. Using the canonical model adopted until now to fit the energy spectrum of this source large  residuals appear below  1 keV.  These are  well fitted using  an  equivalent  hydrogen  column  of  $0.66  \\times  10^{22}$ cm$^{-2}$,  adding absorption  edges of  \\ion{O}{7},  \\ion{O}{8} and \\ion{Ne}{9} in the spectra  extracted from the observation at phases 0.62--0.84  and adding absorption  edges of  \\ion{O}{7}, \\ion{O}{8}, \\ion{Mg}{11} and \\ion{Mg}{12} and absorption lines of \\ion{O}{8} and \\ion{Mg}{12} in the spectra extracted from the observation at phases 0.11--0.16.   During  the   observation  at  phases  0.62--0.84  the electron density associated to  the ionized matter is $\\sim 10^{13}$ cm$^{-3}$  remaining  quite constant  going  away  from the  compact object.  During  the observation  at phases 0.11--0.16  the electron density profile varies along the  distance going from $\\sim 6 \\times 10^{13}$ cm$^{-3}$ at  $\\sim 10^{11}$ cm to $\\sim  9 \\times 10^{10}$ cm$^{-3}$  at $\\sim  10^{13}$  cm.  The  equivalent hydrogen  column towards Cir X--1  is three times lower than  the value obtained from previous models.  This  low value would imply that Cir  X--1 is at a distance of 4.1 kpc. ",
        "introduction": "Cir X--1  is a peculiar Low  Mass X-ray Binary  (LMXB), which recently showed ultra-relativistic outflows similar to those produced by active galactic  nuclei   (Fender  et  al.   2004).  Because  of   its  rapid variability, a  black hole  was supposed to  be the compact  object in this  system (Toor  1977)  until  type I  X-ray  bursts were  detected (Tennant et al. 1986a, 1986b), suggesting that the compact object is a neutron star.  Until 2002 Cir X--1 was supposed to be a runaway binary because of its supposed association  with the supernova remnant  G321.9-0.3 (Clark et al.   1975)  at a  distance  between  5.5 kpc  and  6.5  kpc (Case  \\& Bhattacharya,  1998; Stewart  et  al.  1993).   However,  the work  of Mignani et  al.  (2002), using  HST data, excluded the  association of Cir  X--1 with  G321.9-0.3.  This  brings to  a new  debate  about the distance to Cir  X--1. From the observed 21-cm  absorption features in the radio  spectrum of  the source  a lower limit  to the  distance of 6.4--8~kpc  was  obtained  (Goss   \\&  Mebold,  1977;  Glass,  1994). Generally in LMXBs  the companion star has a  mass $\\la 1$ M$_{\\odot}$ and the  orbit is almost  circular; contrarily the  canonical scenario accepted for Cir  X--1 argues that the source has  a companion star of $\\sim  3-5$ M$_{\\odot}$  (probably a  subgiant,  see Johnston  et al. 1999), an orbital  period of $\\sim 16.6$ days,  deduced from the radio and x-ray lightcurve of the source, and an eccentric orbit with $e\\sim 0.7-0.9$ (Murdin et al.  1980; Tauris et al. 1999).  The correlated  spectral and  fast-timing properties of  accreting low magnetic neutron  stars allow distinguishing them into  two classes, Z and atoll sources, using the pattern each source describes in an X-ray color-color diagram (Hasinger \\& van der Klis 1989). The Z sources are very  bright and  could have  a moderate  magnetic field  ($< 10^{10}$ Gauss), while the atoll sources have a lower luminosity and probably a weaker magnetic  field.  The timing  analysis of RXTE data  (Shirey et al.  1996) evidenced  that Cir X--1 is more similar to  LMXBs of the Z class.    However  no   quasi  periodic   oscillations   at  kilohertz frequencies were observed in Cir X--1 (Shirey, 1996) contrarily to all the other Z-sources.  Cir X--1 also shows a  peculiar X-ray spectrum.  Brandt et al. (1996), using ASCA/GIS data taken in August 1994, found a sudden variation (in a timescale  of $\\sim  20$ minutes)  of the flux  at phase  just after zero, where the count rate increased from $\\sim 30$ counts s$^{-1}$ to $\\sim 300$ counts s$^{-1}$. In  the low-count rate state they obtained a good  fit to the  0.6--10 keV spectrum  of Cir X--1 using  a partial covering  component, in  addition  to a  two-blackbody  model, with  a corresponding  hydrogen  column of  $\\sim  10^{24}$  cm$^{-2}$ and  an intrinsic absorption  of $\\sim 1.7  \\times 10^{22}$ cm$^{-2}$;  in the high-count rate state  the partial covering of the  X-ray spectrum was not needed.  Because  of the similarity with the  spectra of Seyfert 2 galaxies with Compton-thin tori,  Brandt et al.  (1996) suggested that matter  at the  outer edge  of the  accretion disk,  together  with an edge-on disc  orientation, could explain  the partial covering  of the spectrum of Cir X--1 (see, however, Fender et al. 2004).  Iaria et al. (2001a, 2002), using BeppoSAX data, analysed the 0.1--100 keV spectrum  of Cir X--1  at phases 0.11--0.16 and  0.61--0.63.  They found that the  continuum could be well described  by a Comptonization model at phases 0.11--0.16 and  by the combination of a blackbody plus a Comptonization  model at phases 0.61--0.63;  the blackbody component is probably produced in the  inner region of the accretion disc, while the Comptonized component is produced by a (hotter) corona surrounding the neutron star.  A hard excess, dominanting the spectrum at energies higher than  15 keV,  was also observed  in both the  observations and fitted  by  a power-law  or  a  bremsstrahlung  component.  A  similar feature was  observed in  other Z-sources  (GX 17+2, Di  Salvo et  al. 2000; GX 349+2, Di  Salvo et al. 2001; GX 5--1, Asai  et al. 1994; Sco X--1, D'Amico et al. 2000; Cyg X--2, Di Salvo et al. 2002), suggesting that Cir  X--1 shows a  behavior similar to  that of other  Z-sources. Finally a  strong absorption edge  at $\\sim 8.4-8.7$ keV,  produced by highly ionized  iron, was needed  to fit the  spectrum of Cir  X--1 at phases 0.11--0.16.  The presence of  a hard power-law component was  recently confirmed by RXTE (3--200 keV)  observations of Cir X--1 (Ding  et al. 2003).  They extracted the source spectra at different positions in the color-color diagram  and concluded  that there  is no  clear relation  between the position of the source in  the color-color diagram and the presence of the hard tail in the energy spectrum.  An absorption edge, with energy ranging from 8.3  keV to 9.2 keV, was again  observed, together with a localized emission feature present around 10 keV in the energy spectra corresponding  to the  normal  branch and  flaring  branch; a  similar feature was observed by Shirey et al. (1999). They interpret this line as produced by heavy elements, such as nickel, or, less probably, as a blue-shifted iron line  due to the motion in a  relativistic jet or to the rotation  of the accretion disk; however  extreme conditions would be required to boost the line energy up to 10 keV in this case.  Iaria et al. (2001b), using  ASCA data, studied the energy spectrum of Cir  X--1 along its  orbit. They  distinguished three  different X-ray states of  the source as  a function of  its phase.  Therefore  it was possible to study the evolution of the absorption edge along the orbit of Cir  X--1. At  the phase  of the flaring  activity, near  the phase zero, the  hydrogen column  derived from the  edge was  $\\sim 10^{24}$ cm$^{-2}$;  at  larger  values   of  the  phase  the  hydrogen  column decreased;  the absorption  edge  (from ionized  iron)  was no  longer detected at  phases from 0.78  to 1, while  at these phases  a partial covering component was required  with an equivalent hydrogen column of $\\sim 10^{24}$ cm$^{-2}$.  In general  the X-ray spectrum of  Cir X--1 is quite  rich of discrete features.  Brandt  \\& Schulz (2000), using Chandra  data, observed Cir X--1 near the phase zero. They found P-Cygni profiles of lines emitted by highly ionized elements.  The outflow velocity necessary to explain the P-Cygni profiles was $\\sim  2000$ km/s.  Under the hypothesis that the source is  seen nearly edge-on, they supposed that  the model of a thermally  driven  wind (Begelman  et  al.  1983)  could explain  this outflow of  ionized matter from  the system. Also they  suggested that the  absorption  edge  observed  by  Iaria et  al.  (2001a)  could  be connected to the outflow.  The main  point that  should be  clarified is the  cycle of  16.6 days observed in the radio and X-ray lightcurve that was supposed to be the orbital period of  the source. Just before the  phase zero, associated to  the  radio  outbursts,  a  dip  is present  in  the  x-ray  folded lightcurve that was  associated to the eclipse of  the emitting source by  the  companion  star.   The  long-term variability  of  the  x-ray lightcurve  and the  large flaring  activity, present  near  the phase zero, carried to the conclusion  that Cir X--1 has an elliptical orbit with a large value of its  eccentricity (Murdin et al., 1980) and that the phase zero  corresponds to the periastron of  the orbit.  As noted by  Clarkson et  al.  (2004)  the  dips are  near-coincident with  the presumed periastron passage suggesting  that the semimajor axis of the binary orbit would have to be almost aligned with our line of sight to the binary system requiring  an improbable fine tuning.  Saz Parkinson et al.  (2003), studying the RXTE/ASM lightcurve of Cir X--1 noted that the periodicity  of 16.6 days  is not always observed.   They analysed the data from 6 Jan 1996 to 19 Sep 2002 finding that the cycle of 16.6 day is not present in a temporal interval lasting 270 days and that in this interval a  new weak periodicity of 40 days  is observed.  If the cycle of 16.6  days is the orbital  period of Cir X--1 is  not easy to explain why  during around  16 orbital periods  the periodicity  is no more observed  while a new periodicity  of 40 days  appears.  Then the results reported by Saz Parkinson et al.  (2003) and Clarkson et  al.  (2004) seem to indicate that  the cycle of 16.6 days in  not connected  to the  orbital period of  the binary  system. Recently, from three radio  observations Fender et al.  (2004) resolve the structure  of a jet  in Cir  X--1. The jet  is extended up  to 2.5 arcsec from the core.  The  authors, assuming a distance to the source of 6.5 kpc,  find a superluminal motion of the  jet having an apparent velocity of $v_{app}  \\simeq 15 c$; because of the  large value of the apparent  velocity  the  angle  between  the line  of  sight  and  the direction  of the  jet, $\\theta$,  is $  < 5^\\circ$  and  the intrinsic velocity  of  the  jet is  $\\beta  >  0.998$.   The detection  of  the superluminal  motion in the  jet of  Cir X--1  implies that  $\\theta < 5^\\circ$   and  since  the   jet  should   have  a   direction  almost perpendicular to the accretion disk it implies that Cir X--1 could not be an edge-on source contrarily to  what was  supposed until  now.  In this  work we assume  the ephemeris as  reported by Stewart et  al. (1991).   We present the  results of  the analysis  of the  broad band (0.12--200 keV) spectra of Cir  X--1 from a BeppoSAX observation taken at phases 0.62--0.84 in September 2001 and the reanalysis of the broad band spectrum of  Cir X--1 from a previous  BeppoSAX observation taken at phases  0.11--0.16 in  August 1998 (see  Iaria et al.,  2001a).  We note that, although the used model in that paper well fitted the data in the broad  energy band 0.1--200 keV, a small  bump in the residuals below 1 keV was present (see Fig.  3 in Iaria et al., 2001a).  We will show that to well  fit the data below 1 keV we  need to add absorption edges  of  H-like and  He-like  ions  of  oxygen, neon  and  magnesium confirming the  presence of highly  ionized matter around the  system. Moreover we discuss  the possibility that Cir X--1  is distant 4.1 kpc implying  that  it  does  not  show a  super-Eddington  luminosity  as supposed until now.  The  BeppoSAX observation of  Cir X--1  at phases  0.61--0.63, already published by Iaria et al. (2002)  and cited below in the text, will be not  reanalysed because the  available energy  range for  the spectral study is 1.8--200  keV, not allowing us to study  the spectrum below 1 keV. ",
        "conclusions": "In this work we present the analysis of the broad band (0.12--200 keV) spectrum of  Cir X--1 at phases  0.62--0.84 and the  reanalysis of the broad band (0.12--200 keV) spectrum  of Cir X--1 at phases 0.11--0.16, using   two  BeppoSAX  observations   taken  in   2001  and   1998  in respectively.  In the  spectra, a soft excess below  1 keV is observed using the model previously proposed.   We fitted the soft excess using the equivalent  hydrogen column of 0.66 $\\times  10^{22}$ cm$^{-2}$ and adding at low energies absorption edges  of \\ion{O}{7}, \\ion{O}{8}, \\ion{Ne}{9}, \\ion{Mg}{11} and \\ion{Mg}{12}. Moreover in the observation at phases 0.11-0.16  two absorption lines of \\ion{O}{8} and \\ion{Mg}{12} were added.   The  equivalent hydrogen column  of the  interstellar matter  of $0.66 \\times 10^{22}$  cm$^{-2}$ gives a  distance to Cir  X--1 of 4.1  kpc. This result is  discussed comparing it to the  values obtained by Goss \\& Mebold (1977)  and reanalyzing the HI spectrum  of the source.  We discuss the different continuum observed in the two observations as due to  the different  density of  the  ionized matter  around the  binary system in the two observations.  Finally the reader should  note that our results could be model dependent; for this reason further spectroscopic analysis of the Cir X--1 spectrum below 1 keV are need to confirm the scenario proposed in this work."
    },
    "0410/astro-ph0410545_arXiv.txt": {
        "abstract": "Recent observational and theoretical results on starburst galaxies related to the wavelength regime below 1200~\\AA\\ are discussed. The review covers stars, dust, as well as hot and cold gas. This wavelength region follows trends similar to those seen at longer wavelengths, with several notable exceptions. Even the youngest stellar populations show a turn-over in their spectral energy distributions, and line-blanketing is much more pronounced. Furthermore, the O~VI line allows one to probe gas at higher temperatures than possible with lines at longer wavelengths. Molecular hydrogen lines (if detected) provide a glimpse of the cold phase. I cover the crucial wavelength regime below 912~\\AA\\ and the implications of recent attempts to detect the escaping ionizing radiation. ",
        "introduction": "The astrophysically important wavelength region below $\\sim$1200~\\AA\\ is still relatively unexplored, at least at low redshift where restframe observations must be obtained from space. Prior to the launch of FUSE (Moos et al. 2000), far-ultraviolet (far-UV) studies were limited to bright objects. The earliest spectral data for bright stars were obtained by Copernicus (Rogerson et al. 1973) and ORFEUS (Grewing et al. 1991), and with the UV spectrometers on the Voyager~1 and 2 spacecraft (Longo et al. 1989). Voyager~2 also succeeded in recording a far-UV spectrum of M33 (Keel 1998). HUT (Davidsen 1993) was the first instrument sensitive enough to collect spectra of faint galaxies below Ly-$\\alpha$. The mission was flown on two missions and generated a rich database of far-UV spectra of actively star-forming and starburst galaxies. Subsequently, FUSE with its superior resolution and sensitivity fully opened the far-UV window to starburst galaxies. Most of this review deals with results obtained with FUSE and, to a smaller degree, with HUT. ",
        "conclusions": ""
    },
    "0410/astro-ph0410635_arXiv.txt": {
        "abstract": "{ Constraints on the density and thermal 3D structure of the dense molecular cloud core \\rod\\ are derived from a detailed 3D radiative transfer modeling. Two ISOCAM images at 7 and 15 $\\mu$m are fitted simultaneously by representing the dust distribution in the core with a series of 3D Gaussian density profiles. Size, total density, and position of the Gaussians are optimized by simulated annealing to obtain a 2D column density map. The projected core density has a complex elongated pattern with two peaks. We propose a new method to calculate an approximate temperature in an externally illuminated complex 3D structure from a mean optical depth. This \"$T_{\\overline\\tau}$\"-method is applied to a 1.3 mm map obtained with the IRAM 30m telescope to find the approximate 3D density and temperature distribution of the core \\rod. The spatial 3D distribution deviates strongly from spherical symmetry. The elongated structure is in general agreement with recent gravo-turbulent collapse calculations for molecular clouds. We discuss possible ambiguities of the background determination procedure, errors of the maps, the accuracy of the $T_{\\overline\\tau}$-method, and the influence of the assumed dust particle sizes and properties. ",
        "introduction": "Stars and planetary systems are believed to form within interstellar molecular clouds, from the collapse of condensations of matter often called pre-stellar cores. Despite recent observational and theoretical progress, the physical processes leading to the formation of these cores, and their collapse remain poorly known \\citep[e.g.][]{2000prpl.conf...59A}. The study of pre-stellar cores is essential to the understanding of star formation, since these objects represent the initial conditions of gravitational collapse which influence the subsequent stages of the star-forming process.  \\subsection{Verification of 1D models of early star formation by observations of molecular cloud cores} In the case of low-mass star-forming regions, it has traditionally been thought that stars form from the inside-out collapse of singular isothermal spheres that are initially in quasi-static equilibrium \\citep{1987ARA&A..25...23S}, supported against gravity by magnetic and thermal pressure and evolve only due to slow ambipolar diffusion processes \\citep{1991ApJ...373..169M}. In this case, pre-stellar cores at the onset of collapse are expected to follow a $1/r^2$ density profile at all radii $r$ \\citep{1977ApJ...214..488S}. Then they collapse dynamically under their own weight and form a protostar, which subsequently accretes matter from  its surrounding envelope.  In the last decade, much effort has been put in the understanding of these cold ($T\\sim$ 5-15 K) pre-stellar objects, which emit most of their radiation in the (sub)millimeter domain with observations that benefitted from the development of millimeter bolometers, sensitive line detector systems, and infrared cameras. Those observations made it possible to derive line-of-sight integrated properties of the pre-stellar cores. Several studies using (sub-)millimeter continuum \\citep{1994MNRAS.268..276W, 1996A&A...314..625A, 1998A&A...336..150M, 2004Laun} or mid-infrared (MIR) absorption \\citep{2000A&A...361..555B, 1996A&A...315L.329A} have in particular addressed the density structure. Determining the density structure of pre-stellar objects is fundamental to understand core and protostar evolution, since it has been previously shown that the density structure of {\\it pre-stellar} objects influences the proto-stellar accretion rate, i.e. after the formation of a central object. According to \\citet{1997A&A...323..549H} \\citep[see also][]{1993ApJ...416..303F}, non-singular density profiles in pre-stellar cores lead to accretion rates declining with time  in the proto-stellar stage, whereas the accretion rate is expected to be constant if the core has singular $r^{-2}$ profiles \\citep{1977ApJ...214..488S, 1987ARA&A..25...23S}.  It was found from the mm and MIR maps that column density profiles of pre-stellar objects are flattened out in their centers. Interpreted by 1D singular isothermal sphere models, the derived density distributions do not follow the $r^{-2}$-profile from the \\citet{1987ARA&A..25...23S} model, though the profiles do have an $r^{-2}$ part.  The fact that some cores even present sharp edges at large radii \\citep[$\\sim 0.1$\\,pc, see][]{2000A&A...361..555B} tends to show that the reservoir of mass available for subsequent accretion is limited. Therefore to some extent, the density structure of pre-stellar cores must govern stellar masses and influence the initial mass function \\citep{2003IAUS..221E.154U}.  Using the 1D approximation, several types of hydrostatic and quasi-static models have been proposed to account for core formation and evolution. Comparison of the density profiles derived from the observations with the models has been used in order to discriminate between various models \\citep[e.g.][]{1994MNRAS.268..276W, 1998Obs...118..346W, 1996A&A...314..625A, 2000A&A...361..555B}. In most of these studies, circular (or elliptical) averages of the dense core emission (either in the mm or in the MIR) have been performed according to the apparent (projected) geometry of the object. The azimuthal average was also used to increase the signal-to-noise ratio. The profiles thus derived were compared to various power law density models. It was found that the ambipolar diffusion models of magnetically-supported cores provide the best fits, requiring a background magnetic field ranging from 30 to 100 $\\mu$G \\citep{2003ApJ...592..233W}.  \\subsection{Limitations of 1D models of early star formation} Both from observational and from the theoretical side, there are many indications that star formation is an intrinsically three-dimensional process. The profile analysis was often performed in 1D or 2D because of the difficulty of obtaining information along the 3rd dimension, as well as the substantially enhanced numerical effort in 3D star-formation simulations and 3D radiative transfer modeling.  However, the increased sensitivity of bolometer arrays as well as the observations of cores in absorption in the mid-infrared \\citep[a technique which is more sensitive to small column densities - e.g.][]{2000A&A...361..555B, Stein04} reveals filamentary structures for the cores at large scales, which significantly differ from near-spherical approximations \\citep{2004PhDT.........1A}. Statistical analyses of core projected shapes \\citep{2001ApJ...551..387J, 2002ApJ...569..280J} also tend to show that cores are triaxial. If 1D interpretations represent in simple cases (i.e. for isolated cores) a good approximation of the cores, they can also lead to erroneous density structures, as illustrated by \\citet{2003ApJ...592..188B}. More specifically, \\citet{Stein04} used non-symmetric, filamentary density distributions from a smooth particle hydrodynamics simulation of a cloud core evolution to produce MIR and mm maps. By interpreting the absorption maxima in MIR maps and emission maxima mm maps with a 1D model, they found density profiles in ostensible agreement with standard 1D core formation models, but in contradiction to the true underlying non-symmetric distribution.  Moreover, the increased sensitivity and resolution of upcoming telescopes will make it possible to see small-scale structures in these objects and make more sophisticated models more and more necessary. The 3D density structure and eventually its evolution in time is also of crucial interest for chemical models \\citep[e.g.][]{2003ApJ...593..906A}.  \\subsection{Gravo-turbulent star formation} The hypothesis that star formation mostly occurs in magnetically subcritical cores and is regulated by slow ambipolar diffusion processes \\citep{1987IAUS..122....7S} has been challenged from several sides \\citep[see, e.g., the reviews by][] {Larsen2003, 2004RvMP...76..125M}. It is only applicable to isolated, single stars, while it is known that the majority of stars form in small aggregates or large clusters \\citep{2001ApJ...553..744A, 2003ARA&A..41...57L}. Furthermore, there is both observational evidence \\citep{2004mim..proc..123C, 2000prpl.conf...59A, 2001ApJ...554..916B, 2001ApJ...561..871H, 2003ApJ...592..233W} and theoretical reasoning \\citep[e.g.][]{1998ApJ...494..587N} that cloud cores do not have magnetic fields strong enough to support the core against gravitational collapse.  Also, the long lifetimes implied by the quasi-static phase of evolution in the model are currently discussed regarding, e.g., observational statistics of cloud cores \\citep{1996A&A...313..269T, 1999ApJS..123..233L, 2002AJ....124.2756V} and chemical age considerations \\citep{1998ARA&A..36..317V, 2000prpl.conf...29L}. \\citet{TasMou04}, however, have argued that the observational data are consistent with the predictions of the theory of self-initiated, ambipolar-diffusion--controlled star formation.  An alternative to mediation by magnetic fields as controlling agent for stellar birth is supersonic interstellar turbulence \\citep{2004RvMP...76..125M}. The presence of highly supersonic (and usually super-Alfv{\\'e}nic) turbulent motions in Galactic gas clouds is inferred from molecular line measurements.  The observed line widths typically exceed the values for thermal line broadening by far \\citep{1993prpl.conf..125B}. Only on very small scales and for the so-called quiescent cores does the total linewidth become comparable to the thermal width alone. It is important to note that turbulence usually carries sufficient energy to counterbalance gravity on global scales. The presence of supersonic turbulence therefore prevents contraction of the cloud as a whole. On small scales, however, it may actually provoke localized collapse \\citep{1982ApJ...256..505H, 1993ApJ...419L..29E, 1995MNRAS.277..377P, 1999ApJ...527..285B, 2000ApJ...535..887K, 1999ApJ...526..279P, 2002ApJ...576..870P, 2001ApJ...547..280H}. Supersonic turbulence establishes a complex network of interacting shocks, where converging flows generate regions of enhanced density.  The compression can be sufficient for local gravitational instability to set in.  The same random flow that creates density enhancements in the first place, however, can also disperse them again. For star formation to occur, the localized compression must be fast enough for the contraction region to `decouple' from the flow.  The theory of gravo-turbulent star formation predicts that molecular cloud cores form at the stagnation points of the complex turbulent flow pattern. Their density profiles are constant in the inner region, exhibit an approximate power-law fall-off further out, and usually appear truncated at some maximum radius \\citep[see][]{2003ApJ...592..188B}. Such configurations closely resemble equilibrium Bonnor-Ebert spheres \\citep{1955ZA.....37..217E, 1956MNRAS.116..351B}, which are popular models for interpreting observed cloud cores. The best studied example probably is the Bok globule B68 \\citep{2001Natur.409..159A, 2003ApJ...586..286L}. However, supersonic turbulence does not create hydrostatic equilibrium configurations.  Instead, the density structure is transient and dynamically evolving, as the different contributions to virial equilibrium do not balance \\citep{2003ApJ...585L.131V}. Numerical models predict that gravo-turbulent fragmentation of molecular cloud material will lead to cores that are generally triaxial and in many cases highly distorted.  Depending on the projection angle, they often appear extremely elongated, being part of a filamentary structure that may connect several objects \\citep{Klessen04}.  It is the goal of this paper to extend our understanding of pre-stellar cores by entering the complex and difficult world of 3D modeling. The target of our investigation is the pre-stellar core \\rod\\ located in the $\\rho$ Ophiuchi star-forming cluster. This core was selected because it has been observed both in the MIR and mm range and the maps indicate a cloud structure deviating from spherical symmetry. We start by outlining the general strategy of the analysis in Sect.~2. In Sect.~3, we present results of a simultaneous fit of ISOCAM images at 7 and 15 $\\mu$m by a series of 2D Gaussian profiles to obtain the number column density at each image pixel. The necessary background determination scheme is described and discussed. We compare the derived dust density to an IRAM 30m 1.3 mm map in Sect.~4. Presenting and applying a new method to obtain 3D density and temperature information of cores, we derive constraints on the structure of the core \\rod. A discussion of ambiguities and errors of the obtained distributions is given in Sect.~5. In Sect.~6, we compare the derived spatial distribution with distributions typically obtained in gravo-turbulent simulations and conclude our findings in Sect.~7. ",
        "conclusions": "In this paper, we propose a new method to model the three-dimensional dust density and temperature structure of dense molecular cloud cores. The method is based on the fits of multiple continuum images in which the core is seen in absorption and emission and requires only a few computations with a 3D CRT code. The large parameter space of a three-dimensional structure is covered using simulated annealing as optimization algorithm. One of the key points of the proposed method is the use of a $T_{\\overline\\tau}$-relation that has been derived from the 3D CRT runs. It links the mean optical depth at a given wavelength from outside to a spatial point within the core to the dust temperature at this point, making a fast evaluation of the mm emission along the line of sight possible.  We have applied the method to model the dense molecular cloud core \\rod. In the MIR, the core is seen in absorption against a bright background from the photo-dissociation region of the nearby B2V star, illuminating the cloud from behind. Two ISOCAM images at 7 and 15 $\\mu$m have been fitted simultaneously by representing the dust distribution in the core with a series of 3D Gaussian density profiles. The background emission behind the core was interpolated from nearby regions of low extinction. Using simulated annealing, we have obtained a 2D column density map of the core. The column density of the core has a complex elongated pattern with two peaks, with the southern peak being more compact. To retrieve the full three-dimensional structural information, we have calculated the temperature structure of the core with a 3D CRT code, assuming several limiting cases of core extent along the line of sight. It was shown that, for a given position in the core, the relation between the dust temperature and the mean optical depth from outside to this position varies within less than 1 K, when changing the shape of the core along the line of sight. Therefore, the $T_{\\overline\\tau}$-relation was used to estimate quickly the temperature at a given point by a fast direction averaging over the optical depth at just one wavelength. Using this mean, we have varied extent and position of the Gaussian components of the density along the line of sight until a fit of both the MIR and a 1.3mm map, obtained with the IRAM 30m telescope, could be achieved. We have presented the three-dimensional dust density and temperature structure, revealing a condensed southern part and a more extended and complex northern part. We have addressed several sources of errors, namely the background determination method, the assumed dust particle properties, errors in the maps, the ambiguity in the derived distribution, and the approximation error of the $T_{\\overline\\tau}$ method. As we did not reach a perfect fit of the mm image between the two maxima due to the low number of Gaussians, we expect that the detailed 3D structure in-between the density maxima might change when using a higher number of Gaussians. The general structure of a condensed southern maximum and a complex northern multi-clump region, however, is clearly resolved within the limits of the presented analysis. This structure is in general agreement with recent gravo-turbulent collapse calculations for molecular clouds. We speculate that the southern condensation belongs to the category of cloud clumps which are dominated by self-gravity. It may be collapsing to form a star in the near future. The northern clumps, however,  may still have sufficiently large internal or kinetic energy to re-expand and merge into the ambient medium once the turbulent compression subsides.  In this application, we determine about 210 free parameter to fit about 2000 pixels simultaneously, but hidden in the assumptions are many more free parameter like the dust properties or details of the illuminating background. It has to be kept in mind, though, that the hidden parameters are present in any other, more simple modeling as well, and often with more unrealistic approximations. The $\\rho$ Oph cloud is illuminated from behind in contradiction to a 1D boundary condition for the incoming radiation.  The perspective of applying the method is promising. There are a number of well-observed cores where images at more than two wavelengths are available. With every additional image, the ratio of new unknown background parameter versus constrained parameter decreases, as the density parameters are unchanged. This will increase the accuracy of the determined density structure. Moreover, the derived $T_{\\overline\\tau}$-relations for \\rod\\ can be used in other applications incorporating the temperature.  The ultimate goal of applying the method to well-observed cores, however, will be to address the key question of early star formation, namely if the considered cores have in-falling material. The current line observations provide the molecular line emission flux integrated over all moving gas cells along the line of sight. In the general case, gas motion and emissivity of the cells can not be disentangled, and the 1D approximation or shearing layers are assumed to unfold it. Without unfolding it, infall motion can be mixed up with rotational motion leaving it undecided if the core shows any sign for the onset of star formation. This is changed if the core has been investigated with the $T_{\\overline\\tau}$-method. Knowing the full 3D structure in dust density and temperature, the line of sight-integral can be inverted providing the complete kinematical information if the considered line is optically thin and a model for the depletion of the considered molecules is used. This direct verification of infall motion will be subject to a forthcoming publication."
    },
    "0410/astro-ph0410403_arXiv.txt": {
        "abstract": "We compute the rate of diffusive nuclear burning for hydrogen on the surface of a ``magnetar\" (Soft Gamma-Ray Repeater or Anomalous X-Ray Pulsar).  We find that hydrogen at the photosphere will be burned on an extremely rapid timescale of hours to years, depending on composition of the underlying material. Improving on our previous studies, we explore the effect of a maximally thick ``inert\" helium layer, previously thought to slow down the burning rate. Since hydrogen diffuses faster in helium than through heavier elements, we find this helium buffer actually increases the burning rate for magnetars.  We compute simple analytic scalings of the burning rate with temperature and magnetic field for a range of core temperature. We conclude that magnetar photospheres are very unlikely to contain hydrogen. This motivates theoretical work on heavy element atmospheres that are needed to measure effective temperature from the observed thermal emission and constrains models of AXPs that rely on magnetar cooling through thick light element envelopes. ",
        "introduction": "The surface composition of young, isolated neutron stars is not well constrained by theory.  Since the spectrum is formed in the outermost $\\sim 10^{-17} M_\\odot$, it is difficult to predict the surface composition from supernova simulations, especially in light of uncertainties in the amount of fallback, nuclear burning in the nascent star, and subsequent accretion and spallation reactions. On the observational side, the absence of well-understood atomic spectral lines, and the extreme physical conditions at the photosphere impede progress using spectral modeling, though some inferences have been made ( Pavlov, Sanwal \\& Zavlin 2002).  Diffusive nuclear burning (DNB) is a mechanism by which the composition of young neutron star (NS) surface layers can evolve with time (Chang \\& Bildsten 2003; 2004, and Chang, Arras \\& Bildsten 2004, hereafter paper I, II and III respectively). The simple picture of DNB (Chiu \\& Salpeter 1964, Rosen 1968) can be understood as follows (see Fig. 1 of paper I). Rapid sedimentation ($\\sim 1 \\,{\\rm sec}$) allows H to settle above heavier elements. While nuclear burning timescales for H at the surface are far too long to be of interest, a diffusive tail readily extends to depths at which proton captures occur, increasing the burning rate by many orders of magnitude.  Anomalous X-Ray Pulsars (AXP's) and Soft Gamma-Ray Repeaters (see e.g. Woods \\& Thompson 2004 for a review), are a class of strongly magnetized ($B \\sim 10^{14-15}{\\rm G}$), hot ($T_{\\rm bb} \\sim (5-7)\\times 10^6{\\rm K}$) \\footnote{ The relation between $T_e$, the effective temperature used in our theoretical work, and $T_{\\rm bb}$, the blackbody fit, depends on the composition and magnetic field strength and is uncertain.} , young ($P/2\\dot{P}\\sim 10^{3-4}{\\rm yr}$), and slowly rotating ($P\\sim 5-10\\ {\\rm sec}$) neutron stars. We will refer to them as ``magnetars\", following Duncan and Thompson (1992).  In this Letter, we show that the high temperatures and substantial alteration of the equation of state (EOS) by superstrong magnetic fields combine to increase the rate of DNB by many orders of magnitude as compared to other observed neutron stars. This conclusion is robust for a variety of heavy proton capturing elements, even when a He layer exists between H and the proton-capturing element. Consequently, magnetar photospheres are most likely composed of heavy elements, possibly helium. ",
        "conclusions": "We have shown that diffusive nuclear burning removes hydrogen from the photospheres of magnetars on a timescale much shorter than the age. This conclusion holds for a range of proton capturing elements up to Fe. Our new understanding of stratification in Coulomb liquids allows us to include a maximally thick He buffer, which increases the burning rate in the diffusion limited regime.  We conclude that magnetar photospheres will not contain hydrogen unless it can be supplied at a rate \\begin{equation}\\label{eq:H-rate} \\dot{M} \\ga 2 \\times 10^{-16}\\,\\Msolar\\,{\\rm yr}^{-1}\\left(\\frac {y_H}{10^2\\,{\\rm g\\,cm}^{-2}}\\right)\\left(\\frac {\\tau_{\\rm col}} {1 {\\rm day}}\\right)^{-1}. \\end{equation} The burning time is so short that the supply of hydrogen must be continuous over the timescale of the observations.  The relativistic outflow expected from magnetars likely precludes accretion of hydrogen. However, spallation reactions due to magnetospheric currents may generate fresh hydrogen.  The absence of hydrogen will modify our understanding of both the thermal evolution and the thermal spectrum of magnetars.  Atomic physics (e.g. Lai 2001) and spectral modelling of heavy element atmospheres in $B\\sim 10^{14}\\,{\\rm G}$ fields is relatively unexplored.  The light atmosphere models of Heyl \\& Hernquist (1997), Ozel (2001) and Ho \\& Lai (2001) assume fully ionized H and may be significantly different if the photospheric material is He or heavy elements.  Even after $T_e$ has been determined, the question remains as to whether the data require an additional heat source (e.g. field decay; Heyl \\& Kulkarni 1998) or whether the apparent ``hotness\" is due to a thick layer of low opacity material (i.e.  helium; Heyl \\& Hernquist 1997). Models of AXP's that rely on magnetar cooling with a H envelope thicker than even a photosphere are difficult to realize given the rapid timescales we have found for H consumption.  Thermal stability restricts hot ($T>2 \\times 10^8\\,{\\rm K}$) He envelopes to a maximum column of $y_{\\rm He} \\la 10^8\\,{\\rm g\\,cm}^{-2}$.  Since $T_c \\approx 2-3 \\times 10^8\\,{\\rm K}$ (Yakovlev, private communication) for a non-superfluid NS between the ages of $10^3-10^4$ years, the effective temperature of such a maximal envelope ($y_{\\rm He} = 10^8\\,{\\rm g\\,cm}^{-2}$ with C below) is $T_e \\approx 3 \\times 10^6\\,{\\rm K}$ at $B=10^{15}\\,{\\rm G}$. In comparison a thermally unstable pure H or pure He envelope would have $T_e \\approx 5 \\times 10^6\\,{\\rm K}$ and $T_e \\approx 4\\times 10^6\\,{\\rm K}$, respectively. Therefore, passively cooling neutron stars without internal heat sources (such as magnetic field decay; Heyl \\& Kulkarni 1998 and references therein; also see Arras, Cumming \\& Thompson 2004) fall short of the observed luminosity of AXPs by a factor of 3 to 30 (Perna et al. 2001)."
    },
    "0410/astro-ph0410129_arXiv.txt": {
        "abstract": "The first galaxies that formed at a redshift $z\\sim 20$--30 emitted continuum photons with energies between the Ly$\\alpha$ and Lyman limit wavelengths of hydrogen, to which the neutral universe was transparent except at the Lyman-series resonances. As these photons redshifted or scattered into the Ly$\\alpha$ resonance they coupled the spin temperature of the 21cm transition of hydrogen to the gas temperature, allowing it to deviate from the microwave background temperature. We show that the fluctuations in the radiation emitted by the first galaxies produced strong fluctuations in the 21cm flux before the Ly$\\alpha$ coupling became saturated. The fluctuations were caused by biased inhomogeneities in the density of galaxies, along with Poisson fluctuations in the number of galaxies. Observing the power-spectra of these two sources would probe the number density of the earliest galaxies and the typical mass of their host dark matter halos. The enhanced amplitude of the 21cm fluctuations from the era of \\Lya coupling improves considerably the practical prospects for their detection. ",
        "introduction": "The reionization history of cosmic hydrogen, left over from the big bang, provides crucial fossil evidence for when the first stars and black holes formed in the infant universe \\citep{BL01}. The hyperfine spin-flip transition of neutral hydrogen (\\ion{H}{1}) at a wavelength of 21 cm is potentially the most promising probe of the cosmic gas before reionization ended. Observations of this line at a wavelength of $21\\times (1+z)$ cm can be used to slice the universe as a function of redshift $z$ and obtain a three-dimensional map of the diffuse \\ion{H}{1} distribution in it \\citep{Hogan}. The 21cm signal vanished at redshifts $z\\ga 200$, when the small residual fraction of free electrons after cosmological recombination kept the gas kinetic temperature, $T_k$, close to the temperature of the cosmic microwave background (CMB), $T_\\gamma$.  Subsequently, between $200\\ga z\\ga 30$ the gas cooled adiabatically, faster than the CMB, and atomic collisions kept the spin temperature $T_s$ of the hyperfine level population below $T_\\gamma$, so that the gas appeared in absorption \\citep{Scott}. Primordial density inhomogeneities imprinted a three-dimensional power-spectrum of 21cm flux fluctuations on scales down to $\\la 10$ comoving kpc, making it the richest data set on the sky \\citep{Loeb04}. As the Hubble expansion continued to rarefy the gas, radiative coupling of $T_s$ to $T_\\gamma$ dominated over collisional coupling of $T_s$ to $T_k$ and the 21cm signal began to diminish.  However, as soon as the first galaxies appeared, the UV photons they produced between the Ly$\\alpha$ and Lyman limit wavelengths propagated freely through the universe, redshifted or scattered into the Ly$\\alpha$ resonance, and coupled $T_s$ and $T_k$ once again through the Wouthuysen-Field \\citep{Wout, Field} effect by which the two hyperfine states are mixed through the absorption and re-emission of a Ly$\\alpha$ photon \\citep{Madau, Ciardi}. Emission of UV photons above the Lyman limit by the same galaxies initiated the process of reionization by creating ionized bubbles in the neutral gas around these galaxies. While UV photons were blocked at the neutral boundary of the bubbles, X-ray photons propagated farther into the bulk of the intergalactic gas and heated $T_k$ above $T_\\gamma$ throughout the universe. Once $T_s$ grew larger than $T_\\gamma$, the gas appeared in 21cm emission with a brightness temperature that was asymptotically independent of $T_s$ and of $T_k$.  The ionized bubbles imprinted a knee in the power-spectrum of 21cm fluctuations \\citep{Zalda04}, which traced the topology of \\ion{H}{1} as the ionized bubbles percolated and completed the process of reionization \\citep{Fur04}.  Since the first galaxies represented rare peaks in the cosmic density field, their spatial distribution possessed unusually large fluctuations on large scales \\citep{Bar04a}. In this paper we show that the resulting large-scale bias of the first galaxies produced substantial 21cm fluctuations during the intermediate epoch when the first Ly$\\alpha$ photons coupled $T_s$ to $T_k$. Since a relatively small number of galaxies contributed to the flux seen at any given point, Poisson fluctuations were significant as well, and we show that they produced correlated 21cm fluctuations. Owing to the rapid evolution in the density of galaxies with cosmic time, there is a sharp decline in the abundance of sources with distance (corresponding to retarded cosmic time for the source emission due to the light travel time) around each point in the intergalactic medium. Ultimately, there is an effective horizon around each point out to which sources may contribute to 21cm fluctuations. In addition, the atomic physics of resonant scattering by hydrogen fixes different horizon distances for photons emitted in various wavelength intervals between \\Lya and the Lyman limit; these horizons are set by the requirement that a given photon must redshift or scatter into the \\Lya resonance at the point of interest. We incorporate all of these ingredients into our calculation.  The outline of the paper is as follows. In \\S~2 we summarize the basic equations that describe the evolution of the spin temperature. In \\S~3 we review the angular anisotropy of the 21cm power spectrum due to peculiar velocities, as derived in our previous paper \\citep{BL04}. We then calculate the flux of \\Lya photons produced by the earliest galaxies (\\S~4) and its fluctuations (\\S~5). Finally, we demonstrate how the resulting power spectrum of fluctuations in the 21cm flux (\\S~6) can be used to study the properties of the earliest galaxies (\\S~7), and briefly summarize our main results (\\S~8). Throughout the paper we assume the concordance set of cosmological parameters \\citep{WMAP}. ",
        "conclusions": "We have shown that the 21cm flux fluctuations from intergalactic hydrogen are strongly enhanced due to the radiation emitted by the first galaxies at redshifts $\\sim 20$--$30$.  The earliest galaxies emitted photons with energies between the Ly$\\alpha$ and Lyman-limit transitions of hydrogen, to which the neutral universe was transparent except at the Lyman series resonances. As these photons redshifted into a Lyman-series resonance and then degraded into a Ly$\\alpha$ photon, they coupled the spin temperature $T_s$ of the 21cm transition of hydrogen to the gas kinetic temperature $T_k$, allowing it to deviate from the microwave background temperature, $T_\\gamma$. Figures 4--7 show that the fluctuations in the radiation emitted by the first galaxies enhanced the overall power-spectrum of 21cm flux fluctuations by at least a factor of a few during the period when \\Lya coupling was established. The enhancement is sourced by inhomogeneities in the biased distribution of galaxies (as those represent rare peaks in the density field) as well as by Poisson noise in the number of galaxies. The effects of galaxies are more apparent when the power spectrum is split into angular components relative to the line of sight (with the anisotropy sourced by peculiar velocities); in particular, the component $P_{\\rm un-\\delta}$ shown in the lower panels of Figures 4--7 (eq.~\\ref{eq:un-d}) would vanish in the absence of Poisson fluctuations in the distribution of galaxies. For the density-induced fluctuations (upper panels), the enhancement factor is roughly constant on scales $\\ga 10$ comoving Mpc. Because of the rapid evolution in the abundance of galaxies at early cosmic times, and the atomic physics of resonant scattering by hydrogen, the fluctuations in the \\Lya intensity are dominated by the contribution from sources within distances $\\sim 1$--100 comoving Mpc, well within the source horizon out to which a \\Lyb photon could redshift into a \\Lya photon (Fig.~\\ref{radialfig}).  Several experiments are currently being constructed or designed to detect the 21cm fluctuations prior to reionization. For an overview, see http://space.mit.edu/eor-workshop/; for individual experiments, see LOFAR (http://www.lofar.org/), MWA (http://web.haystack.mit.edu/arrays/MWA/index.html), PAST (http://xxx.lanl.gov/abs/astro-ph/0404083), and the future SKA (http://www.skatelescope.org/). Although low-frequency foregrounds are much brighter than the 21cm signal, they are not expected to include sharp spectral features, unlike the 21cm maps that should decorrelate over small shifts in frequency (corresponding to slicing the universe at different redshifts). Consequently, the prospects for extraction of the 21cm signal are promising \\citep{Shaver, Zalda04, Miguel1, Miguel2, Santos}. The signal-to-noise ratio in these experiments is controlled by the sky brightness and grows as the square-root of the integration time.  Hence, the increase in the amplitude of 21cm fluctuations that we predict has important practical implications for observing the 21cm signal from the corresponding redshifts, since the required integration time is reduced by the square of the enhancement factor.  The enhanced level of fluctuations in the 21cm flux lasted only until the \\Lya coupling of $T_s$ and $T_k$ was saturated. Both before and after the \\Lya coupling epoch, the enhancement (compared to the case of a uniform \\Lya flux) had a lower amplitude (see, e.g., Figure~\\ref{zfig1} where \\Lya coupling was achieved at $z\\sim 20$). Detection of the enhancement factor over the available redshift interval (i.e., the corresponding observed frequency band) can be used to infer the abundance and characteristic mass of the earliest galaxies (\\S 7) without any such galaxy being observed individually. The existence of \\Lya coupling can also be used to determine whether $T_k$ is larger than $T_\\gamma$ in this redshift interval. If the 21cm signal appears in emission rather than in absorption [where absorption is present in the interval $30\\la z\\la 200$ \\citep{Loeb04}] then $T_k>T_\\gamma$, most likely due to heating from X-rays emitted by the same galaxies. Altogether, the 21cm fluctuations can be used to probe the level of UV and X-ray radiation emitted by the earliest galaxies, long before reionization was complete. The milestones of \\Lya coupling between $T_s$ and $T_k$ and X-ray heating of $T_k$ above $T_\\gamma$ provide important precursors for the radiative feedback of galaxies on the IGM that culminated with the completion of reionization."
    },
    "0410/astro-ph0410686_arXiv.txt": {
        "abstract": "{By coupling observations of type~Ia supernovae with results obtained from the best available numerical models we constrain the Hubble constant, independently of any external calibrators. We find an absolute lower limit of $H_{\\circ}$ $>$ 50~\\ksm. In addition, we construct a Hubble diagram with UVOIR light curves of 12 type~Ia supernovae located in the Hubble flow, and when adopting the most likely values (obtained from 1-D and 3-D deflagration simulations) of the amount of \\Nif~produced in a {\\it typical} event, we find values of $H_{\\circ}$ $\\geq$ 66$\\pm$8 and $\\geq$ 78$\\pm$9 \\ksm, respectively. Our result may be difficult to reconcile with recent discussions in the literature as it seems that an Einstein-de Sitter universe requires $H_{\\circ}$$\\simeq$ 46~\\ksm in order to fit the temperature power spectrum of the cosmic microwave background and maintain the age constraints of the oldest stars. ",
        "introduction": "Due to their high intrinsic luminosity and apparent uniformity, type~Ia supernovae (hereafter referred to as SNe~Ia) have become an important distance indicator in modern cosmology. Today, they are utilized to determine the value of the Hubble constant ($H_{\\circ}$) and measure the changes of the past cosmic expansion. For this reason a number of substantial observing campaigns have recently been conducted for SNe~Ia at nearby redshifts \\citep[see][ for a list]{leibundgut00}. As a result there is now a considerable number of events available with superb temporal and photometric coverage. However, there has been little effort made to use these high quality data sets to link observations with the physics of SNe~Ia in a systematic way. The purpose of this article is to combine results obtained from theoretical models with modern data in order to constrain the value of $H_{\\circ}$.  Prior attempts to couple observations with explosion models of SNe~Ia in order to determine the value of $H_{\\circ}$ include the pioneering investigations of \\citet{arnett85}, followed by \\citet{branch92}, \\citet{leibundgut92}, \\citet{nugent95} and \\citet{hoflich96}. These works gave promising results, constraining the Hubble constant between 45~$\\leq~H_{\\circ}~\\leq$~105~\\ksm, and revealed that with few assumptions, SNe~Ia used in such a manner provide an attractive way to measure $H_{\\circ}$, while circumventing many problems associated with the extragalactic distance ladder \\citep[see][]{bohm97,livio97}.  Although several progenitor models for SNe~Ia are discussed, the common view favors an accreting carbon oxygen (C-O) white dwarf in a binary system, which undergoes thermonuclear incineration at or near the Chandrasekhar mass \\citep[for reviews see][]{hillebrandt00,arnett96,woosley86}. The energy released from burning to nuclear statistical equilibrium (NSE) at the density and temperatures expected in these explosions completely disrupts the C-O white dwarf, while the subsequent light curve is powered by the Comptonization or ``thermalization'' of $\\gamma$ rays produced from the radioactive decay of $^{56}$Ni $\\rightarrow$ $^{56}$Co $\\rightarrow$ $^{56}$Fe. The single degenerate model became favored over the double degenerate model after strong H$\\alpha$ emission was observed in SN~2002ic \\citep{hamuy03}. The H$\\alpha$ emission in this SN~Ia is thought to come from circumstellar material originating from the white dwarf's evolved companion star \\citep{hamuy03,nomoto04}.  Here we use bolometric light curves of SNe~Ia as a means to link observations with results obtained from models of an exploding C-O white dwarf. Bolometric light curves constructed from observations provide a simple and direct route to probe the complicated explosion physics and radiation transport. As it is typically more straightforward to extract the total flux (hence luminosity) of a SN~Ia from models rather than the flux for individual filters, which require complicated multi-group calculations \\citep{leibundgut92,eastman97,hoflich97,leibundgut00} bolometric light curves provide a sorely needed tool to connect observations with models. In addition, observed bolometric light curves are easily constructed from the integration of broad-band photometry, and near maximum light they reflect the fraction of $\\gamma$ rays thermalized. Consequently $\\sim$80\\% or more \\citep{suntzeff96, suntzeff03} of the thermalized flux from the $\\gamma$ rays is emitted at optical and near-infrared wavelengths (3000-10000~\\AA), therefore what we manufacture from the observed photometry has been coined a {\\it UVOIR} bolometric light curve. It must be noted that by neglecting the small amount of flux outside the UVOIR wavelength regime we introduce a systematic underestimation on the calculated values of H$_{\\circ}$. However, we address this systematic error when placing constraints on H$_{\\circ}$ (see below). The summed UVOIR flux also offers the advantage that we do not need to apply any K-corrections, which are necessary when just using individual filter passbands of SNe~Ia located in the Hubble flow. Essentially, we can reduce the problem to one of energy balances where the energy inputs from the radioactive decays and the losses due to $\\gamma$ ray escape can be compared to the observed wavelength-integrated flux of the SNe~Ia.  In the following we utilize a set of well observed SNe~Ia to demonstrate, via two methods, (see below $\\S$ \\ref{method1} and $\\S$ \\ref{method2}) that --under the assumption that SNe~Ia are a product of the thermonuclear disruption of a Chandrasekhar-mass C-O white dwarf-- it proves to be rather difficult to obtain a value of $H_{\\circ}$ $<$ 50~\\ksm. Our results, along with recent detection of the integrated Sachs-Wolfe (ISW) effect \\citep{boughn04a,boughn04b} observed in the Wilkinson Microwave Anisotropy Probe (WMAP) data, \\citep{bennett03,spergel03} cast doubts on recent discussions in the literature \\citep[see][]{blanchard03,shanks04}, which suggest ``alternatives'' to the concordance cosmological model. Spatially flat, matter-dominated Einstein-de~Sitter models may produce a temperature power spectrum that can fit cosmic microwave background (CMB) observations just as accurately as the best concordance model, which sports a dark energy component. However, Einstein-de~Sitter models require very low values of the Hubble constant (e.g. $H_{\\circ}$ $\\simeq$ 46~\\ksm) and are unable to account for the observed ISW effect.  The structure of this paper is as follows: in $\\S$~\\ref{data} we briefly discuss the basic details and information that we can extract from a UVOIR/bolometric light curve. This is followed by a short discussion of the data we have compiled and the method with which we construct our UVOIR light curves. In $\\S$~3 we discuss different models of SNe~Ia and the \\Nif~yields we adopt for a typical SN~Ia. We then present in $\\S$~\\ref{method1} a method to derive $H_{\\circ}$ through the combination of observations and the theoretical \\Nif~masses calculated in explosion models. Section \\ref{method2} contains a discussion of the classical way to derive $H_{\\circ}$ through the Hubble diagram of SNe~Ia. Contrary to previous methods we employ here the bolometric flux. We conclude in $\\S$~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  Under the main assumptions: (1) that the progenitors of SNe~Ia are C-O white dwarfs that explode at or near the Chandrasekhar mass, and (2) that the amount of \\Nif~synthesized to first order determines the peak luminosity, we are able to use results obtained from state-of-the-art numerical simulations of explosion models to uniquely define the bolometric peak luminosity, and in concert with photometric observations constrain H$_{\\circ}$. The attractiveness of this approach stems from the ability to bypass assumptions that are typically made when one attempts to determine H$_{\\circ}$, i.e. the extragalactic distance ladder and its accumulation of error from rung to rung. We stress that our fitting method does not add any corrections to the data. In other words we {\\it do not} normalize the flux to any decline rate relation (e.g. $\\Delta$$m_{15}$ \\citep{phillips99}, MLCS \\citep{riess96} or stretch \\citep{perlmutter97}).  Uncertainties that marginally affect (in decreasing order of importance) our results include: the abundances of peak Fe group metals -- stable and radioactive -- produced in the explosion models, the redshift peculiar velocities of each SN~Ia, the total absorption, the assumed rise time to bolometric maximum, the exact nature of $\\alpha$, which may slightly vary from SN to SN depending on the exact nature of the opacity and ionization structure of the expanding ejecta, and the amount of flux that we neglect outside of the optical window.  It is still unclear what parameters effect the amount of \\Nif~produced in a SN~Ia explosion. Obvious candidates are the initial conditions prior to explosion. These include the metallicity of the C-O white dwarf, the central density and the ignition mechanism. If there is a considerable fraction of alpha elements such as $^{22}$Ne within the progenitor one would expect more stable isotopes such as $^{58}$Ni and $^{54,56}$Fe to be produced from burning to NSE, thus reducing the \\Nif~yield \\citep[e.g.][]{brachwitz00}. A higher central density on the other hand would lead to a more robust explosion and hence an increased amount of \\Nif. As discussed earlier the explosion mechanism is still uncertain and different deflagration and detonation scenarios produce different amounts of peak Fe group elements. Nevertheless with a larger \\Nif~mass we obtain smaller values of H$_{\\circ}$. Errors attributed to the adopted recession velocity and reddenings produce scatter in our Hubble diagram but have very little effect on our calculations of H$_{\\circ}$ via equations~\\eqref{eq:hubequal} or \\eqref{eq:hub7}. The $\\pm$3 day departure from our adopted rise has no more than a 10\\% effect on our calculations. A change in $\\alpha$ by $\\pm$20\\% can affect H$_{\\circ}$ by 9\\%. Finally, we reiterate from equation~\\eqref{eq:hubequal} that by neglecting $\\sim$10\\% of the flux emitted outside of the optical, we are underestimating H$_{\\circ}$ by at least 5\\% or correspondingly more if more flux is unaccounted in our method.  In $\\S$ \\ref{method2results} we determined a rather high value (85~\\ksm) for the Hubble constant when using results from the MPA model. This indicates that the amount of \\Nif~produced in these 3-D deflagration simulations currently are on the low side. And indeed a large sample of SNe~Ia show that the average distribution of \\Nif~mass is slightly higher ($\\sim$~0.6~M$_{\\sun}$) \\citep{stritzinger05}. This suggests that their models may need more fine tuning in order to produce a larger amount of \\Nif, and hence match observations more accurately.  There have been many attempts since \\citet{kowal68} presented his Hubble diagram to exploit SNe~Ia to determine H$_{\\circ}$. We refer the reader to \\citet{branch98} for a detailed review of previous works that attempt to calculate H$_{\\circ}$ based on SNe~Ia. He concluded, from methods based on physical considerations similar to the methods presented in this work and methods which utilize SNe~Ia that have been independently calibrated by Cepheids, a range in the Hubble constant of 54~$\\leq$~H$_{\\circ}$~$\\leq$~67 \\ksm, with a ``consensus\" on H$_{\\circ}$~$=$~60$\\pm$10~\\ksm. More recent investigations of \\citet{suntzeff99} and \\citet{jha02} give values of H$_{\\circ}$~$=$~64~\\ksm. Finally, \\citep{freedman01,spergel03,altavilla04} have all measured slightly larger values of H$_{\\circ}$~$\\approx~$70~\\ksm with 10\\% accuracy.  Another method independent of the extragalactic distance ladder which combines X-ray imaging of galactic clusters with the Sunyaev-Zel'dovich effect (SZE) has been recently used to place limits on H$_{\\circ}$ \\citep{myers97,mason01, jones03,reese02}. These works have provided the detailed study of 41 clusters giving distances which yield an averaged value of H$_{\\circ}$~$\\approx~ $61$\\pm$3~\\ksm \\citep[for a review see][]{reese03}.  We find from both methods presented here that the Hubble constant must be $>$~50~\\ksm in order to be compatible with current supernova models. In addition, we stress that this lower limit is based on the assumption that 1~M$_{\\sun}$ is an upper limit on the amount of \\Nif~produced in a SN~Ia explosion, and not from our adopted models. This result, along with other methods to measure H$_{\\circ}$ using SNe~Ia calibrated with Cepheids, SZE/X-ray distances and evidence of the ISW effect, strongly suggest that we do not live in a matter dominated universe without some form of cosmological constant or similar agent."
    },
    "0410/astro-ph0410479_arXiv.txt": {
        "abstract": "A new drifting pulsar, PSR J0815+09, was discovered in the Arecibo drift-scan searches \\citep{Mc04}. An intriguing feature of this source is that within the four pulse components in the integrated pulse profile, the sub-pulse drifting direction in the two leading components is opposite from that in the two trailing components. In view that the leading theoretical model \\citep{RS75} for pulsar sub-pulse drifting can only interpret one-direction sub-pulse drifting, the observed \\emph{bi-drifting} phenomenon from PSR J0815+09 poses a great challenge to the pulsar theory. The inner annular gap (IAG), a new type of inner particle accelerator, was recently proposed to explain both $\\gamma$-ray and radio emission from pulsars \\citep{QLWXH04}. Here we show that the coexistence of the IAG and the conventional inner core gap (ICG) offers a natural interpretation to the bi-drifting phenomenon. In particular, the peculiar drifting behavior in PSR J0815+09 can be reproduced within the inverse Compton scattering (ICS) model for pulsar radio emission. ",
        "introduction": "The sub-pulse drifting phenomenon observed in many pulsars has been widely regarded as a powerful tool to probe the pulsar inner magnetospheric structure and radiation mechanism (Ruderman \\& Sutherland 1975, hereafter RS; Gil, Melikidze \\& Geppert 2003). It may be linked to the physical properties of the surface of pulsars (RS75), shedding light on the nature of pulsars, e.g. whether they are normal neutron stars or bare strange stars (Xu, Qiao \\& Zhang 1999).  In the classical drifting model, sub-pulse patterns are manifestations of some sparks passing along a ring (``carousel'') in the polar cap region circulating the magnetic axis (RS75). This interpretation is supported by detailed analysis of drifting data \\citep{DR99,DR01,AD01,GS03}. The RS model is the leading model to interpret the observations quantitatively. More detailed analysis revealed that the circulation speed in a pure vacuum gap is too high when compared with the observations \\citep{GGMG03}. Such a discrepancy may be resolved through introducing partial screening of the gap electric field so that the inner gap is not purely vacuum \\citep{GGMG03,CR80,UM95}. These models are successuful to interpret many observations, including time variations of the drift rate and changes of the apparent drift direction \\citep{GGMG03}.  Recently, a new drifting pulsar, PSR J0815+09, was discovered in the Arecibo drift-scan searches \\citep{Mc04}. An intriguing feature of this source is that within the four pulse components in the integrated pulse profile, the sub-pulse drifting direction in the two leading components is opposite from the one in the two trailing components. We call this ``bi-direction drifting'' or ``bi-drifting'' phenomenon. This phenomenon poses a great challenge to the RS model and its variants, since in all these models, the drifting direction in different spark-rings is expected to be the same. Recently, Qiao et al. (2004) suggested that an inner annular gap (IAG) likely coexists with the conventional inner core gap (ICG) in pulsars (especially when pulsars are bare strange stars), and proposed a phenomenological model to illustrate the origin of the $\\gamma$-ray and radio radiation from pulsars. In this paper, we propose that the bi-drifting phenomenon is a natural consequence of the coexistence of the IAG and the ICG. Furthermore, the complicated drifting patterns in PSR J0815+09 could be reproduced if the pulsar radio emission is dominantly generated through coherent inverse Compton scattering (ICS) processes \\citep{QL98,XLHQ00,QLZH01}.  In \\S 2, we introduce both the ICG and the IAG in pulsars, and the ICS model for radio emission is introduced in \\S 3. Numerical simulations of the observed bi-drifting phenomenon is presented in \\S 4, and the conclusions and discussions are presented in \\S 5. ",
        "conclusions": "By considering both the ICS process and the sparks from the IAG and the ICG, we have demonstrated that the bi-drifting phenomenon observed in PSR J0815+09 is naturally interpreted. This result also lends support to the existence of the IAG. Some parameters and effects related to our simulation are discussed as follows.  (1) The gap heights. Assuming that the spark diameter in the polar gap is the same as the gap height \\citep{GS00}, we have used the resonance inverse Compton scattering induced polar gap model \\citep{ZQLH97} to calculate the gap height to get the spark number in the polar cap. One requirement to the height of the sparks is that some sparks must take place in the IAG. The height of a resonance inverse Compton scattering induced gap in our simulation is consistent with this condition.  (2) The drifting rates. The drifting rates derived from the observation give us insights into the electric structure of the two gaps. The average parallel electric field $E_{\\parallel}$ in the gap can be estimated with $E_{\\parallel}\\sim 4\\pi h(\\rho-\\rho_{\\rm GJ})$. From \\S 2, we know that the perpendicular electric field $E_{\\perp}$ can be estimated as $E_{\\rm \\perp}\\simeq E_{\\parallel}h/\\Delta r$ (see also Gil et al. 2003), where $\\Delta r$ is the transverse spatial dimension of the sparks and has $\\Delta r\\simeq h/2\\simeq1.6 \\times 10^3 $ cm. The drifting velocity of the sparks is therefore $v_{\\rm D}=E_{\\perp} c/B$, where $c$ is the speed of light, and the period for a spark to make a circle in both gaps can be written as $\\hat{P_{3}}=2\\pi r_{\\rm D}/v_{\\rm D}$, or \\begin{equation} \\hat{P_{3}}\\simeq \\frac{B r_{\\rm D} \\Delta r} {2ch^{2}(\\rho-\\rho_{\\rm GJ})} \\label{P3} \\end{equation} The period for a subpulse to reappear at a same phase is defined by $P_3 = \\hat P_3/n$, where $n \\propto r_{\\rm D}$ is the number of sparks in a circle. Since both the IAG and the ICG have the same height $h$ and hence, the same $\\Delta r$, from Eq.(\\ref{P3}) one gets \\begin{equation} \\frac{P_{\\rm 3,ICG}}{P_{\\rm 3,IAG}} \\simeq \\frac { r_{\\rm \\rm D, IAG} \\hat{P_{\\rm 3}}_{\\rm , ICG} } { r_{\\rm D, ICG} \\hat{P_{\\rm 3}}_{\\rm , IAG} } \\simeq \\frac{(\\rho-\\rho_{\\rm GJ})_{\\rm IAG}}{(\\rho-\\rho_{\\rm GJ})_{\\rm ICG}} = \\frac{E_{\\parallel, \\rm IAG}}{E_{\\parallel, \\rm ICG}}, \\label{ele} \\end{equation}  The simulation gives $\\hat{P_{\\rm 3,}} _{\\rm IAG}/\\hat{P_{\\rm 3,}}_{\\rm ICG}\\simeq1.5$, and from \\S 4, we have $r_{\\rm D, IAG}/r_{\\rm D, ICG}\\simeq1.4$ (Fig.~\\ref{cap}). So $E_{\\parallel, \\rm IAG}/E_{\\parallel, \\rm ICG}=P_{3,\\rm IAG} /P_{3,\\rm ICG} \\simeq0.9$. This indicates that the absolute charge density deviation from $\\rho_{\\rm GJ}$ in the ICG is roughly the same as that in the IAG, although a charge deficit presents in the ICG while a charge excess presents in the IAG.  Also the $E_{\\parallel}$ in the ICG is roughly the same as in the IAG but with a different sign. The sub-pulse drifting phenomenon offers a diagnostic tool for the plasma near the polar cap region.  The theoretical drifting period of a vacuum gap for PSR J0815+09 should be order of 1 s (Eq. (\\ref{P3}) and put $\\rho=0$) and is much shorter than the fitted value $\\sim10^2$ s. This has been also noticed earlier by Gil et al. (2003). There are three possible ways to solve this problem. One is to conclude that both the IAG and the ICG are not vacuum gaps, so that only a small charge density deviation from the GJ charge density is allowed. In order to match the observations, only a $\\siml 1\\%$ deviation from the GJ density in both gaps is required. This would naturally reduce the drifting rates, and is consistent with the partial screening picture (Gil et al. 2003). It is also consistent with the XMM-Newton observation results for another drifting pulsar PSR 0943+10 (B. Zhang, D. Sanwal \\& G. G. Pavlov 2004, in preparation). In such a case, an unsteady space charge limited flow may play an important role, and a detailed model is called for. The second reason is overestimating the drifting velocity, because the method above just give the average value of $E_{\\perp} $ which is larger than that at sparking location. The third way to solve the problem is to take into account the drifting dynamics, in which the drifting velocity is not the $E\\times B$ velocity. Again more detailed investigations are needed. It should be noted that the subpulse drifting in the two gaps depends on the global electric property of the pulsar. Two gaps interact with each other via the boundary condition. So a physically reasonable and complete drifting subpulse model should at least includes a global electric solution of pulsars.  It is found that within the observation data span, the four pulse components keep a phase relationship at least within 120 pulses (77s). Our Eq.(2) suggests that $P_3$ in both gaps are roughly the same. A further physical possibility is that the two gaps may interact with each other. Dynamical system theories prove that for two quasi-periodical systems, if there are some small nonlinear interactions, the two systems will be locked into a nearby frequency with a rational ratio \\citep{JBB83}.  Because of the interactions between the two gaps, the two drifting frequencies may be locked into a relative constant ratio in our model, so that the apparent $P_3$ could appear the same for all four pulse components. Detailed interaction dynamics is needed to further address this problem. Some other kinds of interaction between the two gaps are also possible (e.g. Young 2004).  (3) High quality polarization observations are needed to finally verify whether the observed components are the trailing or the leading components.  In summary, the so called \\emph{bi-drifting} phenomenon is a newly observed touch-stone to test the radiation theories and the surface physical properties of the pulsar. Our simulated results support the coexistence of IAG and ICG. Further investigations are needed to address the questions such as how such gaps are generated."
    },
    "0410/astro-ph0410153_arXiv.txt": {
        "abstract": "Due to their intrinsic brightness, supernovae make excellent cosmological probes.   We describe the SEAM method for obtaining distances to Type IIP supernovae (SNe IIP) and present a distance to SN~1999em for which a Cepheid distance exists. Our models give results consistent with the Cepheid distance, even though we have not attempted to tune the underlying hydrodynamical model, we have simply chosen the best fits. This is in contradistinction to the expanding photosphere method (EPM) which yields a distance to SN~1999em that is 50\\% smaller than the Cepheid distance. We emphasize the differences between SEAM and EPM. We show that the dilution factors used in the EPM analysis were systematically too small at later epochs. We also show that the EPM blackbody assumption is suspect.  Since SNe IIP are visible to redshifts as high as $z \\la 6$, with the \\emph{JWST}, SEAM may be a valuable probe of the early universe. ",
        "introduction": "A reliable way to determine accurate distances is a Holy Grail of astronomy and particularly cosmology.  In order to determine the values of the fundamental cosmological parameters, an accurate distance indicator visible to high redshift is required.  Supernovae are extremely bright and hence can be detected at cosmological distances with modern large telescopes.  Due to their homogeneity, SNe Ia had long been thought of as as good distance indicators since they roughly meet the astronomer's definition of a ``standard candle'', that is that the luminosity at peak, $L_{\\mbox{max}}$, is approximately constant. Two Hubble Space Telescope (\\emph{HST}) projects \\citep{WFfinal01,parodi00} were awarded time to use Cepheid variable stars to determine distances to the Virgo cluster and to determine the Hubble constant to 10\\% accuracy. An additional aim of the program of Sandage and collaborators \\citep{parodi00} was to calibrate the luminosity of SNe Ia by obtaining Cepheid distances to galaxies which also were the hosts of SNe~Ia. Distances obtained using Cepheids are considered to be the among the most reliable in astronomy (purely trigonometric methods cannot be used at distances in the Hubble flow), but they are not free of systematic errors and Cepheids are too dim to be observed at large distances. The reliability of SNe Ia as distance indicators improved significantly with the realization that the luminosity at peak was correlated with the width of the light curve \\citep{philm15} and hence that SNe~Ia were correctable candles in much the same way that Cepheids are \\citep{philetal99,goldhetal01,rpk95}. This work and the development of highly efficient search strategies \\citep{perlq097} sparked two groups to use SNe~Ia to measure the deceleration parameter and to discover the dark energy \\citep{riess_scoop98,perletal99}.  All of the work with SNe~Ia is empirical, based on observed SNe~Ia template light curves. Another method of determining distances using supernovae is the ``expanding photosphere method'' \\citep[EPM,][]{kkepm,branepm,eastkir89,esk96} a variation of the Baade-Wesselink method \\citep{baadeepm}. The EPM method assumes that for SNe~IIP, with intact hydrogen envelopes, the spectrum is not far from that of a blackbody and hence the luminosity is approximately given by \\[ L = 4\\pi\\,\\zeta^2\\,R^2\\,\\sigma\\,T^4 \\] where $R$ is the radius of the photosphere, $T$ is the effective temperature,  $\\sigma$ is the radiation constant, and $\\zeta$ is the ``dilution factor'' which takes into account that in a scattering dominated atmosphere the blackbody is diluted \\citep*{hlw86a,hlw86b,hw87}. The temperature is found from observed colors, so in fact is a color temperature and not an effective temperature, the photospheric velocity can be estimated from observed spectra using the velocities of the weakest lines, \\[ R = v\\, t, \\] the dilution factor is estimated from synthetic spectral models, and $t$ comes from the light curve and demanding self-consistency.  Both an advantage and disadvantage of EPM is that it primarily requires photometry. Spectra are only used to determine the photospheric velocity, colors yield the color temperature, which in turn is used to determine the appropriate dilution factor (from model results). This method suffers from uncertainties in determining the dilution factors, the difficulty of knowing which lines to use as velocity indicators, uncertainties between color temperatures and effective temperatures, and questions of how to match the photospheric radius used in the models to determine the dilution factor and the radius of the line forming region \\citep{hamuyepm01,leonard99em02}.  In spite of this the EPM method was successfully applied to SN~1987A in the LMC \\citep{eastkir89,bran87a} which led to hopes that the EPM method would lead to accurate distances, independent of other astronomical calibrators. Recently, the EPM method was applied to the very well observed SN~IIP 1999em \\citep{hamuyepm01,leonard99em02,elmhhamdietal99em03}. All three groups found a distance of 7.5--8.0~Mpc. \\citet{leonard99em03} subsequently used \\emph{HST} to obtain a Cepheid distance to the parent galaxy of SN~1999em, NGC 1637, and found $11.7 \\pm 1.0$~Mpc, a value 50\\% larger than that obtained with EPM.   With modern detailed NLTE radiative transfer codes, accurate synthetic spectra of all types of supernovae can be calculated.  The \\textbf{S}pectral-fitting \\textbf{E}xpanding \\textbf{A}tmosphere \\textbf{M}ethod  \\citep[SEAM,][]{b93j3,b94i1,l94d01,mitchetal87a02} was developed using the generalized stellar atmosphere code \\texttt{PHOENIX} \\citep[for a review of the code see][]{hbjcam99}. While SEAM is similar to EPM in spirit, it avoids the use of dilution factors and color temperatures. Velocities are determined accurately by actually fitting synthetic  and observed spectra. The radius is still determined by the relationship $R = vt$, (which is an excellent approximation because all supernovae quickly reach homologous expansion) and the explosion time is found by demanding self consistency. SEAM uses all the spectral information available in the observed spectra simultaneously which broadens the base of parameter determination. Since the spectral energy distribution is known completely from the calculated synthetic spectra, one may calculate the absolute magnitude, $M_X$, in any photometric band $X$, \\[ M_{X} = -2.5 \\log \\int_{0}^{\\infty} S_{X}(\\lambda)\\, L_{\\lambda}\\, d\\lambda + C_X \\, \\] where $S_{X}$ is the response of filter $X$, $L_{\\lambda}$ is the luminosity per unit wavelength, and $C_X$ is the zero point of filter $X$ determined from standard stars. Then one immediately obtains a distance modulus $\\mu_X$, which is a measure of the distance \\[ \\mu_X \\equiv m_X - M_X - A_X = 5\\log{(d/10\\textrm{pc})}, \\] where $m_X$ is the apparent magnitude in band $X$ and $A_X$ is the extinction due to dust along the line of sight both in the host galaxy and in our own galaxy. \\citet{bsn99em00} found that the early spectra were quite sensitive to the assumed reddening and hence determined a value of $E(B-V) =0.1$ for SN~1999em.  The SEAM method does not need to invoke a blackbody assumption or to calculate dilution factors. ",
        "conclusions": "The SEAM method assumes that supernovae are spherically symmetric, which is not strictly true. However, polarization data indicate that SNe~IIP seem to be more spherically symmetric than other types of core collapse supernovae, most likely because the large intact hydrogen envelope sphericizes the explosion. Thus SNe~IIP appear to be the most promising candidates for using the SEAM method. \\citet{leonard01} found evidence for polarization in SN~1999em at 7--163 days after discovery. Modeled in terms of oblate electron scattering atmospheres, the asphericity was about 7\\%. They found some tendency for increasing polarization with time.  This is consistent with polarization studies of Type Ib/c supernovae where the polarization appears to increase the closer one gets to the central explosion mechanism \\citep{wang02ap03}.  It is difficult to know exactly why the SEAM method gives such a different result from that of EPM. \\citet{leonard99em02} found $t_{\\mbox{exp}} = 5.3$~d, and our date is somewhat earlier. Even with a similar explosion date (see Table~\\ref{tab:dists}) we find a larger distance. Figure~\\ref{fig:epmcomp} compares the color temperature $T_{BV}$, the velocity at the photosphere (defined as $\\tau =2/3$), and the dilution factor, $\\zeta_{BV}$, obtained using $T_{BV}$ with those of \\citet{hamuyepm01}. The results agree very well at early times, but by the 4th epoch the dilution factors disagree by 40\\% and by nearly a factor of 3 at the fifth epoch. Comparing only two epochs, if one mistakenly uses a dilution factor that is too small at the later time, the distance obtained will be too small. With hindsight \\citet{hamuyepm01} recognized this fact when they found that they obtained distances close to the Cepheid value when they restricted their analysis to early times where our dilution factors agree. However, the whole foundation of EPM appears suspect. Figure~\\ref{fig:bbcomp} compares the best fit diluted Planck function with our computed flux at the first epoch where we have fit the observations very well. It is clear that a Planck function does not fit the SED at all. Thus, we find that the diluted blackbody assumption is too simplistic, particularly at later times. That the EPM approach works at early times seems coincidental, but it may be that in the hot early phases the color temperature is reasonably accurate, we will explore this in detail in future work.  \\begin{figure} \\includegraphics[width=0.7\\hsize,angle=90]{f02} \\caption{\\label{fig:epmcomp}The EPM parameters  $v(\\tau =2/3)$, $T_{BV}$, and dilution factor $\\zeta_{BV}$ from our models (filled circles) are compared with those of \\citet{hamuyepm01} (filled triangles). While there is good agreement at early epochs, by the fourth epoch the two results differ by 40\\%.} \\end{figure}  \\begin{figure} \\includegraphics[width=0.7\\hsize,angle=90]{f03} \\caption{\\label{fig:bbcomp}The flux from our model (solid line) compared with the best fit diluted blackbody flux (dashed line).} \\end{figure}   The SEAM method seems clearly superior to EPM since the assumption of black-body emission is never realized in a supernova. SEAM should be testable by the Nearby Supernova Factory \\citep{snfactory02spie} if they follow a dozen or so SNe IIP in the Hubble flow that they will discover.  An independent cosmological probe is highly desirable.  SNe~IIP may be detectable to high redshifts with the James Webb Space Telescope (\\emph{JWST}). With a dataset of spectral models that fit nearby SNe~IIP we will be able to determine the nucleosynthetic history of the first generation of stars."
    },
    "0410/astro-ph0410709_arXiv.txt": {
        "abstract": "{We present detailed predictions of the contributions of the various source populations to the counts at frequencies of tens of GHz. New evolutionary models are worked out for flat-spectrum radio quasars, BL Lac objects, and steep-spectrum sources. Source populations characterized by spectra peaking at high radio frequencies, such as extreme GPS sources, ADAF/ADIOS sources and early phases of $\\gamma$-ray burst afterglows are also dealt with. The counts of different populations of star-forming galaxies (normal spirals, starbursts, high-$z$ galaxies detected by SCUBA and MAMBO surveys, interpreted as proto-spheroidal galaxies) are estimated taking into account both synchrotron and free-free emission, and dust re-radiation. Our analysis is completed by updated counts of Sunyaev-Zeldovich effects in clusters of galaxies and by a preliminary estimate of galactic-scale Sunyaev-Zeldovich signals associated to proto-galactic plasma. ",
        "introduction": "The new broad-band correlators for compact interferometric arrays and multi-beam receivers are  beginning to enable large area surveys at high radio frequencies ($\\ge 15\\,$GHz), hitherto impossible because the beam solid angle scales as $\\nu^{-2}$ (assuming diffraction limited performance) so that surveying large areas of the sky becomes quickly impractical at high frequencies with conventional techniques. The high-frequency surveys will have a major impact on astrophysics:  \\begin{description}  \\item[--] They will open a window on new classes of both Galactic and extragalactic sources, for example on those with strong synchrotron or free-free self-absorption corresponding to both very early phases of nuclear radio-activity (extreme GHz Peaked Spectrum - GPS - sources or high-frequency peakers) and late phases of the evolution of Active Galactic Nuclei (AGNs), characterized by low accretion/radiative efficiency (ADAF/ADIOS sources), as well as to early phases of the evolution of radio afterglows of gamma-ray bursts.  \\item[--] They will provide adequate samples to get a really unbiased view of rare, but very interesting, classes of sources with flat spectrum up to high frequencies, that, at low frequencies, are swamped by more numerous populations which fade away as the frequency increases. One example are blazars, for which much more extended samples are required to properly cover their parameter space. Galactic sources with substantial high-frequency emission include star-forming regions, transient black hole binaries and flaring brown dwarves.  \\item[--] They will allow us to get important information even on the physics of sources known from lower frequency surveys: the high frequency spectral steepenings due to electron aging are informative on the distribution of source ages and on mechanisms for energy losses; the break frequencies marking the transition from optically thick to optically thin regimes depend on the magnetic field intensity, etc..  \\item[--] They will produce surveys of the Sunyaev-Zeldovich effect in distant clusters of galaxies and perhaps on galactic scales, extremely important both to understand the formation of large scale structure and the heating of the intergalactic medium.  \\item[--] They will play a vital role in the interpretation of temperature and polarization maps of the Cosmic Microwave Background (CMB), by allowing us to characterize and remove the contamination by astrophysical foregrounds  \\item[--] They will be essential to calibrate the forthcoming large millimeter array ALMA.  \\item[--] Follow-up high frequency observations will allow us to study the intrinsic polarization of sources, unaffected by Faraday depolarization, as well as their variability properties, essential ingredients to understand their physics.  \\end{description}  \\noindent NASA's WMAP satellite, designed to make accurate measurements of the space distribution of the Cosmic Microwave Background (CMB), is carrying out the first all sky millimeter surveys in five bands (centered at 22.8, 33.0, 40.7, 60.8, and 93.5 GHz; Bennett et al. 2003a). The analysis of the first year observations have yielded a 98\\% reliable catalog of 208 point sources brighter than $\\sim 1\\,$Jy (Bennett et al. 2003b).  Unbiased radio surveys at 15 GHz, covering $520\\,\\hbox{deg}^2$, have been carried out with the Ryle telescope (Taylor et al. 2001; Waldram et al. 2003). Ricci et al. (2004), using a novel wide-band analogue correlator on one baseline of the Australia Telescope Compact Array (ATCA) have surveyed $1216\\,\\hbox{deg}^2$ of the southern sky at 18 GHz. Preliminary source counts at 30 GHz have been produced by the DASI (Kovac et al. 2002), VSA (Grainge et al. 2003), and CBI (Mason et al. 2003) CMB experiments.  The ATCA survey will cover, in the near future, $10^4\\,\\hbox{deg}^2$ at 20 GHz to a flux limit $S_{20{\\rm GHz}} > 40\\,$mJy. The Jodrell Bank Observatory -- University of Torun One Centimetre Radio Array (OCRA) will begin soon a multi-beam survey of the Northern sky at 30 GHz. ESA's {\\sc Planck} satellite will carry out all-sky surveys at 9 frequencies from 30 to 860 GHz with sensitivity several times higher than WMAP.  As a new observational window opens up, it is useful to utilize models exploiting the available information to foretell which source populations the planned surveys will be able to detect as a function of their depth, which redshift range will these survey explore, which information on formation and evolution of each source population will they be able to provide. Answering these questions is a key step towards an optimal planning of the surveys.  Earlier time-honored comprehensive analyses of radio-source evolution (Dunlop \\& Peacock 1990; Toffolatti et al. 1998; Jackson \\& Wall 1999) provided remarkably successful fits to data from surveys at $\\lsim 8\\,$GHz, covering several decades in flux-density. The model by Toffolatti et al. (1998), in particular, was extensively exploited to estimate the radio source contamination of CMB maps at cm/mm wavelengths (e.g.: Mart{\\'\\i}nez-Gonz\\'alez et al. 2003; Vielva et al. 2003; Herranz et al. 2002; Refregier et al. 2000) and turns out to satisfactorily account also for the high frequency counts down to $S \\simeq 20\\,$mJy, with a maximum offset of a factor $\\simeq 0.7$ but only at the brightest fluxes (cf. Fig. 13 of Bennett et al. 2003). On the other hand, the data that have been accumulating in recent years require a more detailed treatment of each of the various relevant sub-populations.  In this paper we present a new analysis encompassing canonical radio source populations as well as a variety of special classes of sources. In Section 2 we describe the model for canonical radio sources, allowing for different evolutionary behaviours for two classes of flat-spectrum sources (flat-spectrum radio quasars and BL Lac objects) and for steep-spectrum sources, while in Sect. 3 we present the data sets used to determine the parameters. In Sect. 4 we work out predictions of counts for star-forming galaxies, extreme GPS sources, ADAF/ADIOS sources, Sunyaev-Zeldovich effects on scales from clusters of galaxies to large galaxies, radio afterglows of $\\gamma$-ray bursts. Our main conclusions are summarized and discussed in Sect.~5.  We have adopted a flat $\\Lambda$CDM cosmology with $\\Omega_{\\Lambda}=0.7$ and $H_0=65\\,\\hbox{km}\\,\\hbox{s}^{-1}\\,\\hbox{Mpc}^{-1}$.  \\begin{figure} \\begin{center} \\includegraphics[height=7cm, width=7cm]{WMAP_counts_int.ps} \\caption{Integral counts of extragalactic sources in the WMAP $K$-band (centered at 22.8 GHz). }\\label{WMAP_counts} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[height=7cm, width=7cm]{WMAP_lat_distr.ps} \\caption{Surface density of WMAP extragalactic sources as a function of Galactic latitude. }\\label{WMAP_lat} \\end{center} \\end{figure}    \\begin{table*} \\begin{center} \\caption{Best fit values of the parameters of the evolutionary models for canonical radio sources.}\\label{tab:parameters} \\begin{tabular}{lcrrccrr} \\hline Source type &\\multicolumn{4}{c}{luminosity function}&\\ & \\multicolumn{2}{c}{evolution}\\\\ \\cline{2-5} \\cline{7-8} & \\multicolumn{1}{c}{$\\log n_0(\\hbox{Mpc}^{-3})$} & \\multicolumn{1}{c}{$a$}& \\multicolumn{1}{c}{$b$}& \\multicolumn{1}{c}{$\\log L_\\ast(\\hbox{erg}\\,\\hbox{s}^{-1}\\,\\hbox{Hz}^{-1}@5\\hbox{GHz},z=0)$} & \\ & \\multicolumn{1}{c}{$k_{\\rm ev}$} & \\multicolumn{1}{c}{$z_{\\rm top}$} \\\\ \\hline FSRQ    & -8.988 & 0.651 &  2.888 & 34.043 & & 0.224 & 2.252 \\\\ BL Lac  & -7.921 & 0.977 &  1.266 & 32.789 & & 0.745 &       \\\\ Steep   & -8.030 & 1.071 &  3.054 & 33.329 & & 4.902 &       \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "A synoptic view of the contributions of the above source populations to the 20 and 30 GHz differential counts is presented in Figs.~\\ref{tot_counts20} and \\ref{tot_counts30}. As expected, and directly demonstrated by the high frequency 9C (Taylor et al. 2001; Waldram et al. 2003) and ATCA (Ricci et al. 2004) surveys, the dominant population at bright flux densities are blazars. These are, however, a composite population, including the rapidly evolving Flat Spectrum Radio Quasars and the slowly evolving BL Lac objects. The former sub-population is prominent above $\\simeq 100\\,$mJy, but its counts converge rapidly at fainter flux densities, eventually falling below those of steep-spectrum sources and approaching those of BL Lacs.  The Figs.~\\ref{tot_counts20} and \\ref{tot_counts30} also show that ongoing high frequency surveys are not far from serendipitous observations of cluster SZ effects, easily recognizable because they show up as negative peaks.  It is well known very deep surveys at 1.4 -- 8.4 GHz that a large fraction of micro-Jy radio sources are moderate redshift active star forming galaxies (Gruppioni et al. 2001; Fomalont et al. 2002; Chapman et al. 2003). Our models are consistent with these results, and indicate that below a few tens of $\\mu$Jy (at 20--30 GHz), there is a further change in the composition of the radio-source population as proto-spheroidal galaxies come into play, suddenly increasing the ratio of star forming to classical radio sources. Since these objects are predicted to be at typical redshifts of 2--3 and to be highly obscured by dust, there appearance will also increase the fraction of optically very faint sources and strongly shift to higher redshifts the redshift distribution.  The emission of star forming galaxies in the spectral region of interest here is dominated by synchrotron and free-free processes at rest-frame wavelengths longer than a few mm, and by thermal dust radiation at shorter wavelengths. The latter emission is characterized by a steeply inverted spectrum, $S_\\nu \\propto \\nu^\\alpha$, with $\\alpha \\sim 4$, causing a strongly negative K-correction and a sharp steepening of the counts. Counts of proto-spheroidal galaxies could be significantly higher than shown in Figs.~\\ref{tot_counts20} and \\ref{tot_counts30} if the mm excess detected in our own Galaxy, combining Archeops with WMAP and DIRBE data (Bernard et al. 2003; Dupac et al. 2003), and in NGC1569 is a general property of the SED of dusty galaxies.  Below $\\simeq 100\\,\\mu$Jy the 20--$30\\,$GHz a significant contribution to the counts may come from Sunyaev-Zeldovich (1972) signals (De Zotti et al. 2004; Rosa-Gonzalez et al. 2004) or free-free emission (Oh 1999) produced by proto-galactic plasma. While no attempts to estimate the counts of proto-galactic free-free sources have been produced yet, we show in Figs.~\\ref{tot_counts20} and \\ref{tot_counts30} the tentative estimates of the counts of SZ effects by De Zotti et al. (2004), that, as pointed out by the authors, should be more conservatively interpreted as (possibly generous) upper limits."
    },
    "0410/astro-ph0410223_arXiv.txt": {
        "abstract": "{We report detailed follow-up observations of the cataclysmic variable HS\\,2331+3905, identified as an emission-line object in the Hamburg Quasar Survey. An orbital period of 81.08\\,min is unambiguously determined from the detection of eclipses in the light curves of HS\\,2331+3905. A second photometric period is consistently detected at $P \\simeq 83.38$\\,min, $\\sim 2.8\\,\\%$ longer than \\Porb, which we tentatively relate to the presence of permanent superhumps. High time resolution photometry exhibits short-timescale variability on time scales of $\\simeq5-6$\\,min which we interpret as non-radial white dwarf pulsations, as well as a coherent signal at 1.12\\,min, which is likely to be the white dwarf spin period. A large-amplitude quasi-sinusoidal radial velocity modulation of the Balmer and Helium lines with a period $\\sim 3.5$\\,h is persistently detected throughout three seasons of time-resolved spectroscopy. However, this spectroscopic period, which is in no way related to the orbital period, is not strictly coherent but drifts in period and/or phase on time scales of a few days. Modeling the far-ultraviolet to infrared spectral energy distribution of HS\\,2331+3905, we determine a white dwarf temperature of $\\Teff\\simeq10\\,500$\\,K (assuming $M_{\\rm wd}=0.6$\\,\\Msun), close to the ZZ\\,Ceti instability strip of single white dwarfs. The spectral model implies a distance of $d=90\\pm15$\\,pc, and a low value for the distance is supported by the large proper motion of the system, $\\mu=0.14\\arcsec\\,\\mathrm{yr^{-1}}$. The non-detection of molecular bands and the low $J$, $H$, and $K$ fluxes of HS\\,2331+3905 make this object a very likely candidate for a brown-dwarf donor. ",
        "introduction": "In the current standard model of cataclysmic variable (CV) evolution \\citep[see ][]{king88-1} the systems evolve to shorter orbital periods, reach a minimum period at $\\sim 70$\\,min and then begin to evolve back to longer orbital periods. As a CV approaches the minimum period, the mass of the donor star decreases until it can no longer sustain hydrogen burning. The low mass secondary star becomes increasingly degenerate, entering the brown dwarf regime. \\citet{kolb93-1} estimated that 70\\% of all CVs should have such ``substellar'' secondaries. Despite these large predicted numbers, only a handful of CVs are currently believed to harbour a brown dwarf donor, with WZ\\,Sge, EG\\,Cnc and EF\\,Eri being among the best candidates \\citep{littlefairetal03-1, patterson98-1, beuermannetal00-1}. So far, for none of them does a compelling proof of the degenerate nature of the star exist. The discrepancy between the observed and predicted numbers of brown dwarf CVs could be due to selection effects: because of the low mass transfer rates predicted for these systems they may be substantially fainter than most CVs and may have very rare outbursts. Alternatively, it may also be that the standard theory is wrong , and that CVs have either not reached the period minimum where their donors become degenerate \\citep{king+schenker02-1}, or they may even die before reaching the period minimum \\citep{patterson98-1}.  In order to be able to  test properly any of these hypotheses we need to reduce the observational biases that afflict the current CV population to a minimum. This is the main purpose of an undergoing large-scale search for CVs using the Hamburg Quasar Survey \\citep[HQS;][]{hagenetal95-1, gaensickeetal00-1, gaensickeetal02-2, gaensickeetal04-1, araujo-betancoretal03-2, rodriguez-giletal04-1}. One of our most recent discoveries, HS\\,2331+3905 (HS\\,2331 thereafter), was selected as a CV candidate on the basis of its spectral characteristics in the HQS. Specifically, its HQS spectrum displays the strong Balmer emission lines that suggest ongoing accretion in CVs.  In this paper we present the results of follow-up spectroscopy and photometry of HS\\,2331 obtained over three years, which confirm the CV nature of this object and unveil a very unusual picture.  In Sec.\\,\\ref{s-obs} we describe the photometric and spectroscopic (ground and space) observations and data reduction and calibration. The analysis of the photometry and ground-based spectroscopy is presented in Sec.\\,\\ref{s-ana_photo} and \\ref{s-ana_spec}. In Sec.\\,\\ref{s-ana_spec} we also present a model for the spectral energy distribution of HS\\,2331, from which we derive a temperature for the white dwarf and distance to the system.  Finally, in Sec.\\,\\ref{s-discussion}, we use our findings in trying to construct a preliminary picture of the nature of HS\\,2331. ",
        "conclusions": "We have discovered a new short orbital period CV, HS\\,2331, which displays an overwhelming  wealth of observable phenomena, and we encourage detailed follow-up studies over all wavelength ranges and time scales to exploit the opportunities that this system offers for an improved understanding of CV evolution and accretion onto white dwarfs."
    },
    "0410/astro-ph0410001_arXiv.txt": {
        "abstract": "In this work we have tested conception of energy dependent power spectra which can be applied for analysis of X-ray dynamic spectra of cataclysmic variables. Using this technique we can study properties of variability in different energy ranges. ",
        "introduction": "Taking  into  account rapidly growing sensitivity of X-ray detectors we can apply to X-ray data more  complicated  techniques of time series analysis. Especially it could  be  useful  to analyse properties of variability in different energy ranges  to  detect different  physical  mechanisms which form behaviour of X-ray  sources.  The idea of energy-dependent time series analysis use binning of X-ray data both  in  time  and  energy  space (so  called ``dynamic spectra'', \\citet{hal2004}). As a result, we can calculate power spectra of light curves for each (narrow enough) energy  bin. Resulting set of power spectra can be showed as a function of  two  variables:  energy  and  time-scale.  Such  product  we called as ``energy-dependent power spectra'' (EDPS). ",
        "conclusions": ""
    },
    "0410/astro-ph0410237_arXiv.txt": {
        "abstract": "One of the principal scientific objectives of the upcoming \\textit{Constellation-X} mission is to attempt to map the inner regions of accretion disks around black holes in Seyfert galaxies by reverberation mapping of the \\fe\\ fluorescence line. This area of the disk is likely radiation pressure dominated and subject to various dynamical instabilities. Here, we show that density inhomogeneities in the disk atmosphere resulting from the photon bubble instability (PBI) can cause rapid changes in the X-ray reflection features, even when the illuminating flux is constant. Using a simulation of the development of the PBI, we find that, for the disk parameters chosen, the \\fe\\ and \\oeight\\ Ly$\\alpha$ lines vary on timescales as short as a few hundredths of an orbital time.  In response to the changes in accretion disk structure, the \\fe\\ equivalent width (EW) shows variations as large as $\\sim$100~eV. The magnitude and direction (positive or negative) of the changes depends on the ionization state of the atmosphere. The largest changes are found when the disk is moderately ionized. The \\oeight\\ EW varies by tens of eV, as well as exhibiting plenty of rapid, low-amplitude changes. This effect provides a natural explanation for some observed instances of short timescale \\fe\\ variability which was uncorrelated with the continuum (e.g., Mrk~841). New predictions for \\fe\\ reverberation mapping should be made which include the effects of this accretion disk driven line variability and a variable ionization state. Reflection spectra averaged over the evolution of the instability are well fit by constant density models in the 2--10~\\kev\\ region. ",
        "introduction": "\\label{sect:intro} Shortly after the prediction of relativistically broadened iron K$\\alpha$ fluorescence lines in the X-ray spectra of accreting black holes \\citep{fab89}, it was realized that if the lines were variable then constraints on the emission geometry and basic black hole parameters could be determined \\citep{stel90,mp92}. These ideas matured into the prospect of iron line 'reverberation mapping', which, given a series of short time-resolved X-ray spectra following a significant flare in the continuum, could accurately derive such fundamental parameters as the black hole mass, spin, and location of the X-ray source(s) \\citep{rey99,yr00,rus00,gl04}. The promise of this technique is such that iron line reverberation mapping of nearby Seyfert~1s is one of the principal science goals of the upcoming \\textit{Constellation-X} observatory.  In recent years it has become obvious that reverberation mapping will be more challenging than perhaps was previously thought. The \\textit{XMM-Newton} telescope has found only a handful of relativistically broadened \\fe\\ lines, which indicates that these components may be in general quite weak and low contrast features. Moreover, the best example of a relativistic line (observed from \\mcg; \\citealt{fab02}) exhibits less variability than the X-ray continuum and shows no correlation with the observed broadband flux (\\citealt{ve01,vf04}, but see \\citealt{pon04}).  These properties seem to apply on long timescales in a wide sample of bright Seyfert~1 galaxies \\citep*{wgy01,mev03}, but there also have been instances of rapid iron line variation with only a slight continuum change (Mrk 841; \\citealt{pet02}), and where the line at least partially follows the continuum \\citep[e.g.,][]{nan97,iwa04}. Clearly, there is a wide range of possible \\fe\\ variability characteristics.  The problem these observations cause with the standard reverberation picture is with the underlying assumption that variations in the \\fe\\ line are directly related to changes in the observed continuum. The observed behavior shows that this assumption cannot be valid at all times, and that there are other effects which are equally important to the production of the \\fe\\ line. One possible perturbation to the previous reverberation calculations is photoionization of the accretion disk surface by the incident X-rays (the work by \\citealt{rey99}, e.g., assumed only neutral reflection from the disk outside of the innermost stable circular orbit). If the disk surface is ionized to a certain level, then the \\fe\\ equivalent width (EW) becomes anti-correlated with the continuum, in contrast to the positive correlation expected if the disk remains neutral \\citep{nay00,bfr02,br02}. It is also possible that the X-ray continuum incident on the disk is different from the observed one. For example, light-bending effects could enhance the illumination \\citep{mf04}, while an outflowing corona could diminish it \\citep{bel99}. While these effects may all be operating at some level, they are still connected with the X-ray continuum, and therefore a relation (positive or negative; strong or weak) between the \\fe\\ line flux and the incident continuum is expected. In this paper, we outline a process by which the \\fe\\ line may vary independently of the illuminating X-ray flux.  For systems accreting at greater than 0.3\\% of the Eddington rate, the innermost annuli of the accretion disk, where the relativistic portion of the iron K$\\alpha$ line is produced, are likely radiation pressure dominated \\citep{ss73}. Numerical results indicate that turbulence driven by the magneto-rotational instability (MRI; \\citealt{bh91}) may involve large density fluctuations in radiation-dominated disks if magnetic pressure exceeds gas pressure and photons diffuse from compressed regions \\citep{tsks03,tur04}. Large density variations may also be caused by the photon bubble instability (PBI; \\citealt{ar92,gam98}), which may develop into a series of shocks propagating through the plasma \\citep{beg01}. These results suggest that radiation-dominated disks have time-varying density inhomogeneities on scales smaller than the disk thickness. In the first paper of this series (\\citealt{btb04}, hereafter Paper I), we showed that density variations in the outermost few Thomson depths can lead to reflection spectra differing from those of uniform slabs. Here, we continue our investigation of X-ray reflection from inhomogeneous accretion disks by computing the evolution of the reflection spectrum as the disk structure is changing due to the PBI. As is shown below, the changing density structure in the disk can affect the EW and line flux of the predicted emission lines, despite the constant illuminating continuum.  The next section details the setup and conditions of the PBI model and reflection calculations. We then present the results of the reflection models, both as a function of time and time-averaged (\\S~\\ref{sect:res}). The results are then discussed in \\S~\\ref{sect:discuss} before we summarize and draw our conclusions in \\S~\\ref{sect:concl}. ",
        "conclusions": "\\label{sect:concl} Our conclusions from the second paper in our study on X-ray reflection from heterogeneous accretion disks can be summarized as follows:  \\begin{enumerate} \\item Changes in the density structure of an accretion disk atmosphere from the photon bubble instability can give rise to rapid variability of emission lines in the X-ray reflection spectrum. These variations would be observed to be independent of the illuminating continuum.  \\item This process gives a natural explanation for any short timescale (e.g., hours) line variability that seemed to be uncorrelated with the X-ray continuum (for example, Mrk~841). However, these observations cannot be used as proof that disk instabilities are causing the line variations.  Nor can it offer an explanation for the uncorrelated behavior of the \\fe\\ line on longer timescales (greater than a day), although we can not rule out such an effect in a different region of parameter space.  \\item The timescale of the variability was found to be as small as a few one-hundredths of an orbital time. Therefore, this effect, combined with any changes in ionization state, could be a serious source of noise for future reverberation studies.  \\item Averaging the individual reflection spectra over the 1.55 orbits most affected by the PBI yielded spectra corresponding to the more common long timescale X-ray AGN observations. These averaged spectra were well fit above 2~\\kev\\ by constant density models.  \\item Photon bubble structures move across the surface of the disk in the radiation MHD calculations, passing in and out of the ray where we extracted density profiles.  Variations in the X-rays reflected from a region containing many structures may be less than those from a single line of sight.  The separations, strengths and orientations of the structures are found to vary with the alignment and strength of the magnetic fields.  The fields are likely to evolve over orbital timescales along with the turbulence inside the disk.  Two and three-dimensional reflection calculations may be useful in studying the consequences for reverberation mapping. \\end{enumerate}"
    },
    "0410/astro-ph0410147_arXiv.txt": {
        "abstract": "The self-gravitating gas in the presence of a positive cosmological constant $ \\Lambda $ is studied in thermal equilibrium by Monte Carlo simulations and by the mean field approach. We find excellent agreement between both approaches already for $ N = 1000 $ particles on a volume $V$ [The mean field is exact in the infinite $N$ limit]. The domain of stability of the gas is found to increase when the cosmological constant increases. The particle density is shown to be an increasing (decreasing) function of the distance when the dark energy dominates over self-gravity (and vice-versa). We confirm the validity of the thermodynamic limit: $ N, V \\to   \\infty $ with $ N/V^{\\frac13} $ and $\\Lambda \\; V^{\\frac23} $ fixed. In such dilute limit extensive thermodynamic quantities like energy, free energy, entropy turn  to be proportional to $N$. We find that the gas is stable till the isothermal compressibility diverges. Beyond this point the gas becomes a extremely dense object whose properties are studied by Monte Carlo. ",
        "introduction": "The self-gravitating gas in thermal equilibrium has been thoroughly studied since many years\\cite{gasn,gas2,Emdsg,Chandrsg,Ebertsg,Bonnorsg, Antonovsg,Lyndsg,Jeans,BT}. As a consequence of the long range attractive Newton force, the selfgravitating gas admits a consistent thermodynamic limit $ N, V \\to   \\infty $ with $ \\frac{N}{V^{\\frac13}} $ fixed. In this limit, extensive thermodynamic quantities like energy, free energy, entropy are proportional to $N$\\cite{gasn,gas2}.  In ref.\\cite{gaslambda} we investigated how a cosmological constant affects the properties of the non-relativistic self-gravitating gas in thermal equilibrium by mean field methods. The mean field approximation becomes exact in the limit when the number of particles becomes infinity.  In the present paper we study the stability properties of the self-gravitating gas in the presence of a cosmological constant by  mean field and Monte-Carlo methods. The use of Monte-Carlo simulations is particularly useful since they are like real experiments and allow to unambiguously determine the stability or instability of the system.  We compute both by Monte Carlo and mean field methods several physical quantities: the equation of state, the particle density and the average particle distance for different values of the dark energy. An excellent agreement between both approaches is found except very close to the collapse transition [see figs. 4-6]. The difference between both approaches turns to be, as expected, of the order $ \\frac1{N} $ where the number of particles $N$ we choose in the  Monte Carlo simulations was $ N = 1000 - 2000 $. The  slightly larger split between  Monte Carlo and mean field near the collapse is due to the fact that there the corrections to the mean field become singular \\cite{gasn,gas2}.  We find that the onset of instability in the canonical ensemble coincides with the point where the isothermal compressibility diverges. At this point we find that the dimensionless parameter $ \\zeta = \\frac{G^3 \\; m^6 N^2 \\; P}{T^4} $ is maximal. ($ \\zeta $ was introduced in ref. \\cite{Lyndsg} for $ \\Lambda = 0 $). We find that  the domain of stability of the gas increases for increasing cosmological constant. The dark energy has an anti-gravity effect that disfavours the collapse pushing apart the particles.  In absence of cosmological constant the particle density $\\rho(r)$ is a {\\bf decreasing} function of the distance in the self-gravitating gas (see for example \\cite{gasn}). In the presence of a positive  $ \\Lambda $ we find that $\\rho(r)$ {\\bf decreases} with $r$ when the self-gravity dominates over the dark energy. This happens for $ X \\equiv \\frac{2 \\Lambda V}{m N} < 1$. For $X>1$ the cosmological constant dominates over the self-gravity and the the particle density turns to {\\bf increase} with the distance. This is a consequence of the repulsive character of the dark energy.  \\medskip  The  Monte-Carlo study of the self-gravitating gas shows the validity of the dilute thermodynamic limit introduced in refs.\\cite{gasn,gaslambda}. Namely,  for $ N, V \\to   \\infty $ with $ \\frac{N}{V^{\\frac13}} $ and $\\Lambda \\; V^{\\frac23} $ fixed, a consistent thermodynamic limit is reached where extensive thermodynamic quantities like energy, free energy, entropy are  shown to be proportional to $N$. Furthermore, the Monte-Carlo simulations allow us to study the condensed phase which turns to be an extremely dense body. The phase transition to collapse is found to be of zeroth order.  \\medskip  The outline of the paper is as follows. Section II presents non-relativistic selfgravitating particles in the presence of the cosmological constant, their statistical mechanical treatment and the mean field approach. In section III we analyze the stability of the selfgravitating gas for zero and nonzero cosmological constant $\\Lambda$ while in section IV we give the physical picture of the gas and its particle density as a function of  $\\Lambda$. Section V contains the results of our Monte-Carlo simulations on the phase diagram, the collapse transition, the particle density and the properties of the collapsed phase. ",
        "conclusions": "The behaviour of the self-gravitating gas is significantly influenced by the cosmological constant (or not) depending on the value of the ratio $X$ defined by eq.(\\ref{x}), $$ X = \\frac{2 \\, \\Lambda \\; V}{m \\; N} = 2 \\; \\left(\\frac{\\mbox{cosmological \\; constant}}{\\mbox{mass}}\\right)_V = 2 \\; \\frac{\\Lambda}{\\rho_V} $$ where $\\rho_V \\equiv N \\; m /V $ is the average mass density in the volume $V$ and where the subscript $\\left( \\right)_V$ indicates mass and cosmological constant inside the volume $V$.  $X$ takes small values for astrophysical objects except for clusters of galaxies where $ X{galaxy clusters} \\sim 0.01 $ and especially for the universe as a whole where  $ X{universe} \\sim 4 $. The smallness of  $X$ stems from the tiny value of the present cosmological constant $\\Lambda = 0.663 \\; 10^{-29}$g/cm$^3$\\cite{lam}. The particle mass density is much larger than $\\Lambda$ in all situations except for the universe as a whole.  However, the non-relativistic and equilibrium treatment does not apply for the universe as a whole. Certainly, the galaxy distribution can be considered as a non-relativistic self-gravitating gas in the presence of the cosmological constant. But the gas of galaxies is certainly not yet at thermal equilibrium today. Perhaps, thermal equilibrium may be reached in some billions of years. Incidentally, at thermal equilibrium the particle density is regular at the center of the distribution in the whole range of $ X $. Namely, there is no cusp  at the center in a situation of thermal equilibrium.   \\begin{figure}[htbp] \\rotatebox{-90}{\\epsfig{file=densx3eta1z5.ps,width=14cm,height=14cm}} \\caption{The local density of the gas $\\rho(x,y,z)$ versus the Cartesian coordinates $x$ and $y$ at $z=5$ in a cube of size $9$ for $X=3$ and $\\eta=1$ from Monte Carlo simulations. The density is larger on the boundary of the cube than at the center. } \\label{densinc} \\end{figure}  \\begin{figure}[htbp] \\rotatebox{-90}{\\epsfig{file=densx0p3eta1z5.ps,width=14cm,height=14cm}} \\caption{The local density of the gas $\\rho(x,y,z)$ versus the Cartesian coordinates $x$ and $y$ at $z=5$ in a cube of size $9$ for $X=0.3$ and $\\eta=1$ from Monte Carlo simulations. The density is larger at the center of the cube than on the boundary.} \\label{densdec} \\end{figure}  \\begin{figure}[htbp] \\rotatebox{-90}{\\epsfig{file=profx0p3eta1p6.ps,width=14cm,height=14cm}} \\caption{Average particle distribution in the gaseous phase from Monte Carlo simulations in a cubic volume for $X=0.3$, $\\eta=1.6$ and $N=2000$. A $+$  denotes one particle.} \\label{profgas} \\end{figure}  \\begin{figure}[htbp] \\rotatebox{-90}{\\epsfig{file=profx0p3eta1p8.ps,width=14cm,height=14cm}} \\caption{Average particle distribution in the collapsed phase from Monte Carlo simulations in a cubic volume for $X=0.3$, $\\eta=1.8$ and $N=2000$. A $+$  denotes one particle.} \\label{profcoll} \\end{figure}"
    },
    "0410/astro-ph0410692_arXiv.txt": {
        "abstract": "We report redshift measurements for 12  BL Lacertae objects from a program aimed at obtaining high signal to noise (up to $\\sim$ 500) optical spectroscopy of a mixed sample of objects. The new observations, gathered with the 8 m ESO Very Large Telescope, allowed us to detect weak spectral features down to  a line equivalent width as small as $\\sim$ 1 \\AA. The new redshifts fall in the 0.2$-$1.3 interval. For nine objects we observe emission lines from the active nucleus. In the remaining three cases absorption lines from the host galaxy are found. For two objects we also  detect absorption lines from intervening  systems. ",
        "introduction": "BL Lac objects are active galactic nuclei (AGN) exhibiting strong nonthermal emission, superluminal motion, rapid and large flux and polarization variability. In contrast with other classes of AGN, their spectra are characterized by the absence or extreme weakness of emission lines.  In the standard model of BL Lacs, the weak emission lines (if present) are generated by fluorescence, as in other types of AGN. The line  equivalent width (EW) is reduced by the strong beamed continuum, caused by the alignment with our line of sight of the relativistic jet produced by the nucleus. On the other hand absorption lines  can be produced either from spectral features of the stellar population of the host galaxy or from intervening halos as in the case of quasars (mainly for high redshift sources). Detectability of the absorption features depends inversely on the brightness state of the central source. Because of the weakness of the emission lines (usually EW$\\lesssim$5 \\AA) and the relatively bright central continuum source with respect to the emission from the host galaxy, in many cases the redshift of these objects is unknown and/or very uncertain. The obvious implication of this is that the distance of the sources remains undetermined, hampering a proper evaluation the physical parameters of the objects.  In the past decade a number of projects were carried out to derive the redshift of BL Lac objects either for selected targets or to obtain as much redshift information as possible for complete samples of BL Lacs. Apart from \\citet{heidt04} all previous works  \\citep[e.g.][]{stickel93, veron93, bade94, veron94, marcha96,drinkwater97, landt01, rector01, londish02, carangelo03, hook03} were based on optical spectra collected with $\\leq$4 m class telescopes and are therefore limited by either relatively low S/N ($\\lesssim$50) or to bright objects (m$_V<$15). As a result of this in spite of the use of relatively large aperture telescopes in a number of cases the redshift remained unknown. Since the detection of spectral features critically depends on the S/N of the spectra it is clear that a further step towards the knowledge of the redshift of BL Lacs requires the use of 8m class telescopes.  With these aims in mind we have carried out a project to secure optical spectra of BL Lacs of still unknown or uncertain redshift with the highest possible S/N using VLT in service mode. Such a program can be executed even during non photometric sky and poor seeing conditions and does not interfere with higher priority programs performed at large aperture telescopes. We selected sources from various lists of BL Lacs \\citep[e.g.][]{padovani95a, veron03} with 15$<$m$_{V}<$22, and $\\delta<$20$^{\\circ}$ thus ensuring to obtain high S/N spectra with reasonable exposure times at VLT. Under these assumptions we constructed a sample of $\\sim$60 targets. 31 have been observed in the first two campaigns while the others are scheduled. Of these observed sources 12 have featureless spectra, 2 turn out to be galactic stars, 2 are early type galaxies, 3 are high redshift quasars and 12 are BL Lacs for which the redshift has been measured by us.  In this paper we report the results concerning these 12 objects. A full account of the observations for the featureless BL Lacs and the other sources will be presented in a later publication. The outline of this work is as follows: in section 2 we describe the observations and the data analysis; the results for individual objects are presented in section 3 and  finally we give in section 4 a brief discussion of our findings.  To compute absolute quantities we adopted a cosmology with H$_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.3, $\\Omega_{\\lambda}$=0.7. ",
        "conclusions": "In this work we have described high S/N spectroscopic observations of 12 BL Lac objects from which the redshift, in the range 0.2$<$z$<$1.3, has been derived. While the targets discussed in this paper do not represent a statistical sample, some considerations on the results summarized in table \\ref{tab:lines} can be proposed.   In three cases we have detected faint (EW are $\\sim$1 \\AA) absorption lines  from the host galaxy. These refer to nearby objects with 0.3$<$z$<$0.5. Only for one of them (1RXS J022716.6+020154) the host galaxy was also detected by imaging. The others are therefore good candidates for follow up imaging studies.   For nine objects the redshift derives from broad and/or narrow emission lines in the range 0.2$<$z$<$1.2. For six objects two or more emission lines are observed, providing secure determination of the redshift. Only in three cases a single broad emission line is detected. The redshift for these objects derive from the plausible identification of this line as MgII 2798 \\AA \\ and because of this should be regarded as tentative.  For seven BL Lac objects we have detected the MgII 2798 \\AA \\ emission line with a  luminosity in the range from  0.5 10$^{42}$erg s$^{-1}$ to 7.0 10$^{42}$ erg s$^{-1}$, (mean luminosity 1.9 10$^{42}$erg s$^{-1}$). These luminosities are about one order of magnitude less than the characteristic luminosity of Mg II line in normal quasars \\citep{puchnarewicz97}. In Figure \\ref{fig5} we compare our new measurements with previous data  \\citep{scarpa97,stickel93} for BL Lacs and blazars over the the continuum vs line luminosity plane. Our sources cover the lower luminosity region of this plane and conform to the behavior of the objects in the class. A  correlation between the line and the continuum luminosities is apparent over four dex although part of it is induced by the correlation with the redshift (partial correlation coefficient L$_{line}$-L$_{cont}$ = 0.70 after removing the effect due to the correlation with redshift). At any given continuum luminosity there is a spread of about 2 dex in line luminosity which translates into a difference of about two orders of magnitude in the observed EW. Although the objects appear uniformly distributed in EW, this difference  has contributed  to a different classification of the sources (BL Lacs, HPQs, see also \\citet{scarpa97})."
    },
    "0410/astro-ph0410371_arXiv.txt": {
        "abstract": "Identifying the processes that determine the initial mass function of stars (IMF) is a fundamental problem in star formation theory. One of the major uncertainties is the exact chemical state of the star forming gas and its influence on the dynamical evolution. Most simulations of star forming clusters use an isothermal equation of state (EOS).  We address these issues and study the effect of a piecewise polytropic EOS on the formation of stellar clusters in turbulent, self-gravitating molecular clouds using three-dimensional, smoothed particle hydrodynamics simulations. In these simulations stars form via a process we call gravoturbulent fragmentation, i.e., gravitational fragmentation of turbulent gas. To approximate the results of published predictions of the thermal behavior of collapsing clouds, we increase the polytropic exponent $\\gamma$ from 0.7 to 1.1 at some chosen density $n_{\\rm c}$, which we vary from from $4.3\\times10^4\\,\\mathrm{cm^{-3}}$ to $4.3\\times10^7\\,\\mathrm{cm^{-3}}$.  The change of thermodynamic state at $n_{\\rm c}$ selects a characteristic mass scale for fragmentation $M_{\\rm ch}$, which we relate to the peak of the observed IMF.  We find a relation $M_{\\mathrm{ch}} \\propto n_{\\mathrm{c}}^{-0.5\\pm0.1}$.  Our investigation supports the idea that the distribution of stellar masses largely depends on the thermodynamic state of the star-forming gas.  The thermodynamic state of interstellar gas is a result of the balance between heating and cooling processes, which in turn are determined by fundamental atomic and molecular physics and by chemical abundances. Given the abundances, the derivation of a characteristic stellar mass may thus be based on universal quantities and constants. ",
        "introduction": "Although the IMF has been derived from vastly different regions, from the solar vicinity to dense clusters of newly formed stars, the basic features seem to be strikingly universal to all determinations \\citep{KRO01b}. Initial conditions in star forming regions can vary considerably. If the IMF depends on the initial conditions, there would thus be no reason for it to be universal.  Therefore a derivation of the characteristic stellar mass that is based on fundamental atomic and molecular physics would be more consistent.    There are many ways to approach the formation of stars and star clusters from a theoretical point of view. In particular models that connect stellar birth to the turbulent motions ubiquiteously observed in Galactic molecular clouds have become increasingly popular in recent years. See, e.g., the reviews by \\citet{LAR03} and \\citet{MAC04}.  The interplay between turbulent motion and self-gravity of the cloud leads to a process we call gravoturbulent fragmentation: Supersonic turbulence generates strong density fluctuations with gravity taking over in the densest and most massive regions \\citep[e.g.][]{LAR81, FLE82, PAD95, PAD97, KLE98, KLE00b, KLE01c, PAD02}. Once gas clumps become gravitationally unstable, collapse sets in. The central density increases and soon individual or whole clusters of protostellar objects form and grow in mass via accretion from their infalling envelopes.  However, most current results are based on models that do not treat thermal physics in detail. Typically, a simple isothermal equation of state (EOS) is used. The true nature of the EOS, thus, remains a major theoretical problem in understanding the fragmentation properties of molecular clouds.  Recently \\citet{LI03} conducted a systematic study of the effects of a varying polytropic exponent~$\\gamma$ on gravoturbulent fragmentation. Their results showed that $\\gamma$ determines how strongly self-gravitating gas fragments. They found that the degree of fragmentation decreases with increasing polytropic exponent $\\gamma$ in the range $0.2 < \\gamma < 1.4$ although the total amount of mass in collapsed cores appears to remain roughly consistent through this range. These findings suggest that the IMF might be quite sensitive to the thermal physics. However in their computations, $\\gamma$ was left strictly constant in each case. Here we study the effects of using a piecewise polytropic equation of state and investigate if a change in $\\gamma$ determines the characteristic mass of the gas clump spectrum and thus, possibly, the turn-over mass of the IMF. ",
        "conclusions": "Using SPH simulations we investigate the influence of a piecewise polytropic EOS on the gravoturbulent fragmentation of molecular clouds. We study the case where the polytropic index $\\gamma$ changes from 0.7 to 1.1 at a critical density $n_{\\mathrm{c}}$, and consider the range $4.3\\times10^4\\,\\mathrm{cm^{-3}} \\le n_{\\mathrm{c}} \\le 4.3\\times10^7\\,\\mathrm{cm^{-3}}$.   A simple scaling argument based on the Jeans mass $M_\\mathrm{J}$ at the critical density $n_\\mathrm{c}$ leads to $M_{\\mathrm{J}}\\propto n_{\\mathrm{c}}^{-0.95}$ \\citep[see][]{JKLLM04}. If there is a close relation between the average Jeans mass and the gravoturbulent fragmentation spectrum, a similar relation should hold for the characteristic mass $M_{\\mathrm{ch}}$ of protostellar cores.  Our simulations qualitatively support this hypothesis, however, with the weaker density dependency $M_{\\mathrm{ch}} \\propto n_{\\mathrm{c}}^{-0.5\\pm0.1}$.  So indeed, the density at which $\\gamma$ changes from below unity to above unity defines a preferred mass scale. Consequently, the peak of the resulting mass spectrum decreases with increasing critical density. The distribution not only shows a pronounced maximum but also a power-law tail towards higher masses, similar to the observed IMF.  Altogether, supersonic turbulence in self-gravitating molecular gas generates a complex network of interacting filaments. The overall density distribution is highly inhomogeneous. Turbulent compression sweeps up gas in some parts of the cloud, but other regions become rarefied.  The fragmentation behavior of the cloud and its ability to form stars depend on the EOS.  However, once collapse sets in, the final mass of a fragment depends not only on the local Jeans criterion, but also on additional processes. For example, protostars grow in mass by accretion from their surrounding material.  In turbulent clouds the properties of the gas reservoir are continuously changing. And in addition protostars may interact with each other, leading to ejection or mass exchange. These dynamical factors modify the resulting mass spectrum, and may explain why the characteristic stellar mass depends on the EOS more weakly than expected from simple Jeans-mass scaling arguments.   Our investigation supports the idea that the distribution of stellar masses depends, at least in part, on the thermodynamic state of the star-forming gas.  If there is a low-density regime in molecular clouds where the temperature $T$ sinks with increasing density $\\rho$, followed by a higher-density phase where $T$ increases with $\\rho$, fragmentation seems likely to be favored at the transition density where the temperature reaches a minimum. This defines a characteristic mass scale. The thermodynamic state of interstellar gas is a result of the balance between heating and cooling processes, which in turn are determined by fundamental atomic and molecular physics and by chemical abundances. The theoretical derivation of a characteristic stellar mass may thus be based on quantities and constants that depend mostly on the chemical abundances in the star-lpwh klforming  cloud."
    },
    "0410/astro-ph0410621_arXiv.txt": {
        "abstract": "In this exploratory study, we investigate the bounds on the equation of state of dark matter. Modeling dark matter as a fluid component, we take into account both positive and negative fixed equations of state. Using CMB, supernovae Ia and large scale structure data we find constraints on the equation of state in a modified $\\Lambda$CDM cosmology. We obtain $-1.50  \\times 10^{-6} < w_{dm} < 1.13 \\times 10^{-6}$ if the dark matter produces no entropy and $ -8.78 \\times 10^{-3}< w_{dm} < 1.86 \\times 10^{-3}$ if the adiabatic sound speed vanishes, both at 3 $\\sigma$ confidence level. ",
        "introduction": "In this note, we would like to present bounds on the dark matter equation of state derived from cosmological observations. In other words: given the information from present cosmological observations, how cold is cold dark matter? In this context, we do not want to limit ourselves to a non-negative equation of state since the nature of dark matter is not yet clear. There are a number of particle dark matter candidates and numerous experiments attempting to detect dark matter particles ( see \\cite{Bertone:2004pz} and references therein). However, dark matter may not be a particle at all. We will try to keep an open mind and check if a negative equation of state can be ruled out by cosmological observations. Previous studies focused on the power spectrum properties of warm dark matter \\cite{Colombi:1995ze} and mixed dark matter models \\cite{Davis:1992ui}, but we would like to emphasize that the present work does not consider the bounds on warm dark matter. We will consider both positive and negative equations of state $w_{dm}$ and show that the equation of state of dark matter is already strongly constrained by current observations of the CMB, supernovae Ia and large scale structure. We will not concern ourselves with bounds from other than cosmological observations.  This is meant as an exploratory study, not as rigid modeling of dark matter, and we will therefore work only from a fluid perspective and leave the question open how one could obtain a negative equation of state of dark matter. For clarity, we will only allow for a constant equation of state $w_{dm}$, even though from a particle perspective it is clear that $w_{dm}$ varies with time for warm and mixed dark matter models. We have chosen this approach as there is no simple possibility to give a corresponding particle motivation for $w_{dm}<0$, as specifying the time evolution of $w_{dm}$ would be model dependent.  In models where dark matter interacts with other components of the universe, such as in coupled quintessence models \\cite{Amendola:1999er}, one may obtain a negative equation of state for dark matter. It may be possible to obtain $w_{dm} <0$ by other methods as well, but we are not aware of any such model.  In this work we will use two simple models for dark matter: one with no entropy production and one with vanishing adiabatic sound speed, both with fixed equation of state.  We obtain bounds for a constant equation of state of dark matter of $-1.50  \\times 10^{-6} < w_{dm} < 1.13 \\times 10^{-6}$ if there is no entropy production and $ -8.78 \\times 10^{-3}< w_{dm} < 1.86 \\times 10^{-3}$ if the adiabatic sound speed vanishes, both at 3 $\\sigma$ confidence level.  For this investigation we will assume that the universe is flat and contains a cosmological constant type dark energy component with equation of state $w_{de}=-1$,  dark matter with a variable equation of state $w_{dm}$, baryons, photons and massless neutrinos. We do not include the tensor part in our analysis. In the conclusions, we will address the issue what will change if we relax the flatness assumption and if we have an equation of state different from $-1$.  The plan of this note is as follows: in Section \\ref{sec::pert}, we will introduce the perturbation equations for dark matter with an arbitrary equation of state (which we will from now on call modified dark matter for brevity). In Section \\ref{sec::gamma} we introduce a model with vanishing entropy production, in Section \\ref{sec::sound} we give a different model with vanishing adiabatic sound speed. In Section \\ref{sec::MCMC} we present the bounds on $w_{dm}$ obtained from a Markov Chain Monte Carlo simulation, followed by our conclusions. ",
        "conclusions": "\\label{sec::conclusions} Since this is an exploratory study we have not attempted to formulate a model with realistic variation in the equation of state but chose to show the main effects and bounds derived from current observations using a constant equation of state. For the $\\Gamma=0$ model, it is clear that the main constraint comes from the matter power spectrum. If we would include measurements at smaller scales, the constraints on the equation of state would be even more restrictive. We may conclude that the $\\Gamma=0$ model is unlikely to be realistic. The situation is not so clear for the $c_{ad}^2=0$ model. More accurate measurements of the CMB, especially in the large multipole region, should give tighter constraints. Based on the data we cannot conclude that this model is ruled out. It would be necessary for formulate a specific model before more progress can be made concerning the question of a possible negative equation of state of dark matter.  How strongly dependent are these results on our assumption of flatness and dark energy equation of state $w_{de}=-1$? From Fig. \\ref{fig::power_gamma} we can readily see that the main constraint on the $\\Gamma=0$ model comes from the large scale structure data; since the CMB spectrum does not change much in the allowed parameter range, we may conclude that relaxing the flatness assumption would not make much difference on the constraints. The same is true for an equation of state different from $-1$. From previous studies it is known that for $w_{de}>-1$, structure growth is suppressed at small scales \\cite{Caldwell:2003vp,Doran:2001rw,Ferreira:1997hj}. But for the $\\Gamma=0$ model, this can only make a small difference, since the growth suppression cannot  ameliorate the strong deviation from the LSS measurements at small scales as is readily apparent in Fig. \\ref{fig::power_gamma}.  The situation for the  $c_{ad}^2=0$ case is different. Here, relaxing the flatness assumption and allowing for open or closed universes  will lead to a weaker constraint on $w_{dm}$. The first peak position is sensitive to the geometry of the universe, but increasing or decreasing $w_{dm}$ can in principle shift this peak to be in agreement with the WMAP data, as may be seen in Fig. \\ref{fig::power_sound}. We therefore expect also that the constraint on the total energy $\\Omega_{tot}$ will be weaker than in  the standard $\\Lambda$CDM case. Concerning the possibility that the equation of state of dark energy $w_{de}>-1$ we may say that here, too, the constraints on the dark energy equation of state and $w_{dm}$ will be less stringent. As mentioned above, the main impact of $w_{de}>-1$ is through a suppression of structure growth at small scales, but this may be compensated by decreasing $w_{dm}$ (see Fig. \\ref{fig::power_sound}). It is therefore apparent that allowing for $w_{de} \\neq -1$ will result in weaker constraints on $w_{dm}$.  Relaxing the flatness and $w_{de}=-1$ will therefore have a negligible impact for the $\\Gamma=0$ model but may lead to a significant relaxation of constraints for the $c_{ad}^2=0$ model."
    },
    "0410/astro-ph0410284_arXiv.txt": {
        "abstract": "s{ The cosmological tests are tight enough now to show that the Friedmann-Lema\\^\\i tre $\\Lambda$CDM cosmological model almost certainly is a useful approximation. This means general relativity theory passes significant tests of the extrapolation of some fifteen orders of magnitude from the length scales of the precision tests of gravity physics.} ",
        "introduction": "Cosmology is tested by checking its consistency with  observational constraints that are substantially more numerous than the freedom to adjust the theory and the interpretations of the observations. We are reaching this goal for some aspects of cosmology: particularly impressive is the abundance of evidence that the matter density parameter is in the range \\beq 0.15\\lap\\Omega _m\\lap 0.3. \\label{eq:omegam} \\eeq This certainly is not precision cosmology, but the consistency of estimates that depend on many different aspects of the theory and a broad variety of astronomical methods show that it almost certainly is accurate cosmology, that is, that we have not been misled by systematic errors. That means general relativity theory passes significant empirical tests on the scale of the observable universe. The measurements supporting this remarkable advance are reviewed in Sec.~3. After brief comments on the more modest network of evidence for the detection of $\\Lambda$, in Sec.~4, I conclude with a qualitative assessment of our understanding of the physics of spacetime and gravity on the scales of cosmology, and the case for developing a more systematic framework for judging the empirical evidence. ",
        "conclusions": ""
    },
    "0410/hep-ph0410064_arXiv.txt": {
        "abstract": "Recently new data from the Cosmo-LEP project appeared, this time from DELPHI detector. They essentially confirm the findings reported some time ago by ALEPH, namely the appearence of bundles of muons with unexpectedly high multiplicities, which so far cannot be accounted by present day models. We argue, using arguments presented by us some time ago, that this phenomenon could be regarded as one more candidate for the presence in the flux of cosmic rays entering the Earth's atmosphere from outer space nuggets od Strange Quark Matter (SQM) in form of so called {\\it strangelets}. \\vspace{1pc} ",
        "introduction": "Recently new data from Cosmo-LEP program, this time from DELPHI detector, has been reported \\cite{CosmoDelphi}. Among other things they have confirmed the findings reported before by ALEPH \\cite{CosmoAleph}, namely that one observes bunches of cosmic muons (i.e., produced at the top of the Earth's atmosphere) of unexpected large multiplicities (up to $N_{\\mu} = 150$). Their origin is so far unexplained and no model used in Monte Carlo (MC) programs simulating cascades of cosmic rays (CR) in the atmosphere is able to account for this phenomenon. In \\cite{CosmoDelphi} the expectation was made that source of this discrepancy can eventually come directly from the elementary interaction model used in MC. However, in our opinion, which we would like elaborate here in more detail, it could rather come (at least to a large extent) from the projectile initiating the cascade. Namely, as we have already done in many places on other occassions \\cite{S,RWWAPPB}, we shall argue that the abovementioned results of both experiments can be regarded as yet another signal of the presence in the flux of CR entering the Earth's atmosphere of nuggets of Strange Quark Matter (SQM) called {\\it strangelets}. In this way results of \\cite{CosmoDelphi} and \\cite{CosmoAleph} would just continue a long list of other phenomena explanable in this way like anomalous cosmic ray burst from {\\it Cygnus X-3}, extraordinary high  luminosity gamma-ray bursts from the {\\it supernova remnant N49} in the Large Magellanic Cloud or {\\it Centauro} (to mention only the most interesing and intriguing examples, for more details see \\cite{RWWAPPB,SS} and references therein). In \\cite{RWW} we have already provided successful explanation of ALEPH observations by using notion of strangelets and assuming their flux being the same as obtained from analysis of all previous signals of strangelets present in the literature. (Actually, at that time ALEPH results were circulated only as conference papers, however, the final results presented in \\cite{CosmoAleph} turned out to be identical to those addressed in \\cite{RWW}).  It is worth to remind here that CosmoLEP data are very important because: $(a)$ the high multiplicity cosmic muon events ({\\it muon bundles}) are potentially very important source of information about the composition of primary CR because muons transport in essentially undisturbed way information on the first interaction of the cosmic ray particle with atmosphere; $(b)$ such events have never been studied with such precise detectors as provided by LEP program at CERN, nor have they been studies at such depth as at CERN \\cite{LEP} (ranging between $30$ and $140$ meters what corresponds to muon momentum cut-off between $15$ and $70$ GeV). ",
        "conclusions": "To conclude: we propose to regard the Cosmo-LEP data on CR muons obtained so far as an additional possible signal of the possible SQM admixture present in the primary CR flux. We would like to add here that such admixture would also contribute to CR flux at energies greater than GZK cut-off \\cite{RWWAPPB,M} explaining therefore this phenomenon in a quite natural way\\footnote{If it will be finally confirmed by experiment \\cite{GZK}.}. This makes strangelets interesting subject to investigate in the future.  We would like to close with the following remark. With the flux of strangelets as estimated by us and used here (equal to $F_S/F_{total} = 2.4\\cdot 10^{-5}$ in the enrgy range of tens of GeV) the energetic spectrum of strangelets should fall like $\\sim E^{-2.4}$, i.e., with spectral index being much smaller than for protons. Actually, this result agrees nicely with $A$-dependence of the spectral index of CR's obtaine when fitting the world CR data \\cite{SI}."
    },
    "0410/astro-ph0410216_arXiv.txt": {
        "abstract": "We report on the investigation of the spatial distribution of the flickering sources in the dwarf nova V2051~Oph with eclipse mapping techniques.  Low-frequency flickering originates in the gas stream and is related to the mass transfer process, whereas high-frequency flickering arises in the accretion disk and is probably connected to magneto-hydrodynamic turbulence. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410020_arXiv.txt": {
        "abstract": "We elaborate on the analytical model of Ritter, Zhang, \\& Kolb (2000) which describes the basic physics of irradiation-driven mass transfer cycles in semi-detached compact binary systems. In particular, we take into account a contribution to the thermal relaxation of the donor star which is unrelated to irradiation and which was neglected in previous studies. We present results of simulations of the evolution of compact binaries undergoing mass transfer cycles, in particular also of systems with a nuclear evolved donor star. These computations have been carried out with a stellar evolution code which computes mass transfer implicitly and models irradiation of the donor star in a point source approximation, thereby allowing for much more realistic simulations than were hitherto possible. We find that low-mass X-ray binaries (LMXBs) and cataclysmic variables (CVs) with orbital periods $\\la 6$hr can undergo mass transfer cycles only for low angular momentum loss rates. CVs containing a giant donor or one near the terminal age main sequence are more stable than previously thought, but can possibly also undergo mass transfer cycles. ",
        "introduction": "The possible importance of irradiating the donor star of a semi-detached compact binary by accretion luminosity for its long-term evolution has first been pointed out by Podsiadlowski (1991). Subsequently, the stability of mass transfer with irradiation feedback has been studied in some detail by King et al. (1996, 1997, hereafter KFKR96 and KFKR97), and by Ritter, Zhang, \\& Kolb (2000, hereafter RZK00). In KFKR96 and KFKR97 it was also shown that mass transfer in cataclysmic variables (CVs) and low-mass X-ray binaries (LMXBs) can become unstable against irradiation feedback, and that in case of instability mass transfer proceeds in cycles in which episodes of irradiation-driven mass transfer alternate with low-states during which mass transfer is essentially shut off. Evolutionary calculations of mass transfer cycles which were based on a homology model for the donor and a simplified irradiation model, have been presented in KFKR97 and RZK00. Here we elaborate on the work done in previous studies in two respects: First, we have developed a binary evolution code which allows us to simulate the evolution of compact binaries with mass transfer cycles more realistically than was possible in previous studies, and second, we elaborate on the stability analysis of KFKR96, KFKR97 and RZK00, thereby taking into account a contribution to the thermal relaxation of the donor star which is unrelated to irradiation and which has been neglected in previous studies. ",
        "conclusions": ""
    },
    "0410/astro-ph0410350_arXiv.txt": {
        "abstract": "Typical observational gamma-ray burst (GRB) spectra are discussed and, in this connection, what is the origin of the compactness problem and how it was solved at first. If the threshold for $e^{-}e^{+}$ pair production depends on an angle between photon momenta, then another solution of the compactness problem is possible. We discuss a possibility of the $\\gamma$-rays collimation and the dependence of photon beaming on photon energies. The list of basic assumptions of the scenario describing the GRB source with energy $< 10^{49}$\\,ergs is adduced: the matter is about an alternative to the ultrarelativistic fireball if {\\it all} long-duration GRBs are related or physically connected with {\\it normal/unpeculiar} core-collapse supernovae (SNe). Namely, we consider the questions about radiation pressure and how the jet arises on account of even small asymmetry of the radiation field in a compact GRB source. The possibility of a new approach to explanation of the GRB phenomenon is shown. Possible mechanisms of their generation in regions of size $< 10^8$cm are discussed (a compact model of GRBs). Observational consequences of the compact GRB energy release are considered.  \\keywords {gamma rays: bursts -- gamma rays: theory} ",
        "introduction": "There are direct and indirect observational arguments in favor of physical connection between massive or core-collapse supernovae (SNe) and long duration gamma-ray bursts (GRBs), and the list of publications on the topic is ever-increasing. At first this connection was justified by the fact that all GRB host galaxies turned out to be star-forming or star-bursting galaxies with high massive star-forming rates (e.g., Djorgovski et al. 2001; Frail et al. 2002; Sokolov et al. 2001). There are more and more occurences of SN signs in the GRB afterglow light curves and spectra for GRB 970228 (Galama et al. 1999), GRB 970508 (Sokolov et al. 1998), GRB 980326 (Bloom et al. 1999), GRB 990712 (Bjornsson et al. 2001), GRB 991208 (Castro-Tirado et al. 2001), GRB 000911 (Lazzati et al. 2001), GRB 011121 (Bloom et al. 2002), GRB 021211 (Della Valle et al. 2003), GRB 030329 (Hjorth et al. 2003; Stanek et al. 2003), GRB 031203 (Thompsen et al. 2004). A comprehensive analysis of SN light in GRB afterglows has been recently done by Zeh, Klose, and Hartmann (2004), and see references therein. If there were more cases of clear and indisputable spectral and photometric signs of association between normal core-collapse SNe (Ib/c type and others) and GRBs, it would be a direct proof of the connection between GRBs and massive stars. The increasing statistics of the GRB-SN associations can impose direct and strong observational constraints on GRB beaming and, hence, we could have an {\\it observational} estimation of a real total energy reservoir of GRB sources.   The purpose of this paper is to describe the basic assumptions of a scenario of a GRB source with energy $< 10^{49}$\\,ergs: the matter is about an alternative to the ultrarelativistic fireball if {\\it all} long-duration GRBs are physically connected with {\\it normal/unpeculiar} core-collapse supernovae (SNe). In Section 2 it should be discussed the typical GRB spectra, which is the origin of the compactness problem, and how it was solved at early stages of GRB studies. Sect. 3 concerns with another attempt of  solving the compactness problem, namely, the dependence of the threshold for $e^{-}e^{+}$ pair production on the angle between photon momenta, a soft collimation and the dependence of this collimation on GRB photon energy. In Sect. 4 we discuss and justify observably a rather strong collimation of GRB radiation (for a small number of hard photons in GRB spectra) reaching near-earth detectors. In Sect. 5 we consider the radiation pressure and how a jet or bullet arises if the reason of the relativistic jet is a powerful light pressure of the collimated/non-isotropic prompt radiation of the GRB source. In section 6 we are examining in outline some  possible compact mechanisms of the GRB phenomenon in a compact GRB model. Sect. 7 (Concluding remarks) concerns observational consequences of the ``compact\" GRB energy release, and what new we will see in the sky if the compact GRB source model is true indeed: particulary, we would like to adduce in outline some re-analysis of the observational results of soft X-ray flashes (XRFs), GRBs with strong X-ray excess in spectra (X-Ray-Rich GRB = XRR GRB), normal or classical GRBs, obtained at first with BeppoSAX and then with HETE-2 (Amati et al. 2002; Lamb et al. 2003a).   So, we will try to understand the soft (in the meaning of photon energies) observational GRB spectrum without involving huge kinematical motions of the radiating plasma {\\it a priori}, or without an enormous Lorentz factor, $\\Gamma \\gg 1$. ",
        "conclusions": "So, XRFs can be not collimated at all or slightly collimated (XRR GRB), but with the low total bolometric energy of $\\sim 10^{47}$ ergs. Since  most probably these are actually the explosions of massive SNe at distances of 100 Mpc (Norris 2003; Norris \\& Bonnel 2003), they can be observed much more frequently than it is predicted by the standard fireball GRB model. One should try to find early spectral and photometrical SN features. Then, in general, the observational problem of XRF/XRR/GRB identification becomes a special section in the study of cosmological SNe. (It will be recalled that the GRB~030329/SN 2003dh was a XRR GRB but not a classical GRB.)   As to normal/classical GRBs and especially those ones with many heavy quanta in spectra, it is possible to obtain directly from formula (1) a (kinematical) estimate of the limit collimation of this gamma radiation, which, in turn, independently agrees with the observational ratio (2) of the yearly rates $N_{GRB}/N_{SN} \\sim 10^{-5} - 10^{-6}$. If the matter concerns quanta with $E\\sim 100MeV$ of distant and the most distant GRBs, then from $1 - cos\\theta_{12}\\approx 0.5MeV^2/(100MeV\\cdot100MeV) = 0.5\\cdot 10^{-4}$ it follows that the radiation of such GRBs turns out to be the most collimated. Such photons must be radiated in the cone of an opening of $\\approx 0.5^o$ and be detected in the spectra of the rather distant GRBs with $z\\sim 1$ and farther because of geometrical factor only.   Thus, a natural consequence of our compact model of the GRB source is the fact that distant bursts ($z\\gtrsim 1$) turn out to be harder ones, while close ``GRBs\" ($z \\sim 0.1$) look like XRF and XRR GRBs with predominance of soft X-ray quanta in their spectra (though the factor $1+z$ also works). Naturally, the effects of observational selection due to finite sensitivity of GRB detectors should be also taken into account. For example, the soft spectral component of the distant (classical) GRBs is ``cut\" out by the detector sensitivity threshold. And the isotropic X-ray burst, simultaneous with the GRB, can be simply not seen in distant (classical) GRBs because of the  low total/bolometric luminosity of the source in the compact GRB model ($< 10^{49}$ ergs). Actually, XRF and XRR GRBs have lower values of $E_{iso}$ (so called isotropic equivalent), than GRBs (Amati et al. 2002; Lamb et al. 2003a, 2003b, 2003c).   As a result, only radiation within a narrow solid angle near a selected direction (in the GRB source) is observable for GRB detectors: the soft spectral range remainder for the prompt XRR GRBs with $z\\lesssim 1$ (having climbed over the equipment threshold) and hard (and even heavy) quanta below the threshold of the pair production (1) for classical GRBs with $z\\gtrsim 1$. (Close XRFs with $z\\lesssim 0.1$ can not have the $\\gamma$-ray quanta in their spectra at all for the definit equipment threshold). Though from the review by Postnov (1999) it follows that ``typical\" GRBs are seen in the range of 30 keV --- 100 MeV, but it turns out (and it was known before) that  most GRBs are much softer (see Sect. 2). Not without reason Lamb et al. (2003) were just amazed by this important {\\it observational} result of BeppoSAX and HETE-2 missions. We mean the detection of obvious XRFs and XRR GRBs first by BeppoSAX (Amati et al. 2002) and then by HETE-2. In our compact model of GRB source it (the Amati law) can be a ``simple\" consequence of formula (1) + collimation (most probably) by magnetic field on the surface of the compact object.    In the scenario of jet formation, which was discussed in this paper and which was also used to interpret the GRB 970508 optical transient (OT) light curves (Sokolov 2001b) an isotropic X-ray, optical and radio emission of the {\\it afterglow} of the GRB OT is possible. In X-ray lines (Piro et al. 1999; Yoshida et al. 2001; Piro et al. 2000; Antonelli et al. 2000) it was so for sure. At that an initial assumption was just a possibility of the small GRB collimation (2), which follows from the comparison of the rates of GRBs and SN explosions in distant galaxies. It means that the close relation between GRBs and SNe was taken as a basic assumption. {\\it All} long GRBs are always accompanied by SN explosions, which are sometimes observed, and sometimes not (Sokolov 2001a, 2002). In other words, the long GRB is the beginning of a massive star collapse or the beginning of SN explosion, and GRBs must always be accompanied by SN explosions (of Ib/c type or of {\\it other} types of massive SNe). Then in any case the total energy release at the burst in \\g-rays can be {\\it not more} than the total energy released by any SN ($ <$ or $\\sim 10^{49}$\\,ergs) in all {\\it electromagnetic} waves. (It is interesting that the total energy release in X-ray emission lines observed with BeppoSAX, ASCA, Chandra for GRB 970508, GRB 970828, GRB 991216, GRB 000214 is of the same order --- see the collected data in the paper by Ghisellini et al. (2002) ``Emission lines in GRBs constrain the total energy reservoir.\")   But with so ``low\" total energy of the GRB explosion ($\\lesssim 10^{49}$ ergs) the only possibility to see GRB at cosmological distances ($z\\gtrsim 1$) is the detection of at least the most collimated part of this energy ($1-10$\\%) leaving the source within the solid angle of $\\Omega_{beam} \\sim (10^{-5} - 10^{-6})\\cdot 4\\pi$. The rest can be inaccessible for GRB {\\it detectors} with a limit sensitivity of $\\sim 10^{-7} erg \\cdot s^{-1}\\cdot cm^{-2}$. Certainly, it does not concern the 10~000 times more sensitive X-ray {\\it telescopes} which were used to make sky surveys with the Ariel V, HEAO-1, Einstein satellites (Heise et al. 2001). For limit sensitivity of $\\sim 10^{-11} erg\\cdot s^{-1}\\cdot cm^{-2}$ in the band of $0.2-3.5$ keV the X-ray observatory (Einstein) recorded Fast X-ray Trasients (unidentified with anything) at a rate of $\\sim 10^6 yr^{-1}$ all over the sky. It agrees well with an average rate of the massive SNe explosions in distant galaxies, but for the present, GRB-detectors see only $\\sim 10^{-4}$ part of this huge number of the distant SN explosions as GRBs.   It is natural that at the total/bolometric energy of ``GRB\" $\\sim 10^{47} - 10^{49}$\\,ergs and at the GRB energy (3) released in the narrow cone (2), ``the fireball\" also looks in quite a different way. As to the compactness problem solved by the fireball model for GRB energies of $10^{52} - 10^{53}$\\,ergs, there is no such a problem for ``\\g-burst\" energies $\\sim 10^{47} - 10^{49}$\\,ergs. In any case, allowing for the low \\g-ray collimation from the surface of the compact object -- GRB/XRR/XRF source, which is necessary for the angular dependence of $e^-e^+$ pair production (1), this problem is solved under quite different physical conditions in the GRB-source than that supposed by Piran (1999). In the scenario: {\\it massive star} ---$>$ {\\it WR} ---$>$ {\\it pre-SN $=$ pre-GRB}, in which only a small part of the most collimated radiation with the collimation (2) goes to infinity and, correspondingly, with the total energy of $10^4 - 10^6$ times less than in the standard theory, the source can actually be of a size $\\lesssim 10^8$\\,cm. It means that at the energies of up to $\\sim 10^{49}$\\,ergs the old (``naive\") estimate of the source size resulting directly from the time variability of GRB can be quite true. Thus, the point can be that the burst energy is much less than in the standard fireball model.   To the above-said we add that the suggested compact GRB scenario allows also predicting the behavior of superluminal radio components (which, e.g., have been observed recently for GRB 030329 (Taylor et al. 2004)). As we discussed in Sect. 5, most likely there is no considerable deceleration of the jet/bullet (with the Lorentz factor of order 10). Hence we expect that the superluminal radio components related to the jet have the following properties: \\\\ 1) the radio component will move with the constant observed superluminal velocity; \\\\ 2) the characteristic observed velocity of the superluminal component is of the order of the Lorentz factor, i.e. of order $10 c$.   {\\bf Acknowledgements.} Authors are sincerely grateful to Yu. A. Shibanov for fruitful discussions, T. N. Sokolova for the manuscript editing. The work was supported by RFBR grants: 04-02-16300, 01-02-17106, 03-02-17223, Program of the Presidium of RAN ``Non-stationary Phenomena in Astronomy\", Program of Russian Science and Education Ministry."
    },
    "0410/astro-ph0410399_arXiv.txt": {
        "abstract": "High spatial and spectral resolution spectroscopy of the OH/IR supergiant VY CMa and its circumstellar ejecta reveals evidence for high mass loss events from localized regions on the star occurring over the past 1000 years. The reflected absorption lines and the extremely strong K I emission lines show a complex pattern of velocities in the ejecta. We show that the large, dusty NW arc,  expanding at $\\approx$ 50 km s$^{-1}$ with respect to the embedded star, is kinematically distinct from the surrounding nebulosity and was ejected about 400 years ago. Other large, more filamentary loops were probably expelled as much as 800 to 1000 years ago while knots and small arcs close to the star resulted from more recent events 100 to 200 years ago.  {\\it The more diffuse, uniformly distributed gas and dust is surprisingly stationary with little or no velocity relative to the star.} This is not what we would expect for the circumstellar material from an evolved red supergiant with a long history of mass loss. We therefore suggest that the high mass loss rate for VY CMa is a measure of the mass  carried out by these specific ejections accompanied by streams or flows of gas through low density regions in the dust envelope.  VY CMa may thus  be our most extreme example of stellar activity, but our results also bring into question the evolutionary state of this famous star. In a separate Appendix, we discuss the origin of the very strong K I and other rare emission lines in its spectrum. ",
        "introduction": "The powerful infrared source and OH maser VY Canis Majoris is one of a few very luminous cool hypergiant stars that define the empirical upper luminosity boundary in the HR diagram \\citep{HD94,deJ98}. It  is  thought to represent a very rapid stage of evolution characterized by a  high mass loss rate. The red supergiant stage is a few hundred thousand years long, or less for the most massive stars, but the transition time  from  main sequence to RSG and later from RSG back to a yellow hypergiant may be only a few thousand years \\citep{Sch92,Sch93}. These evolved late-type stars are  major sources of gas, dust and molecules that contribute to the replenishment of the interstellar medium. VY CMa, in particular, is ejecting large amounts of gas and dust at a prodigious rate. With its extended and complex circumstellar nebula, it offers a unique opportunity to study its high mass loss episodes and the detailed kinematics of this process.  VY CMa is a special case even among the rare  hypergiants. It is associated with a region of recent star formation, the young cluster NGC 2362 and H II region Sharpless 310, and the dark cloud L1667, and has been considered a pre--main sequence star (Herbig 1970a,b; Lada \\& Reid 1978) as well as an evolved massive star \\citep{Hy69,RMH75,GW77}. The lack of Li I $\\lambda$6707 \\citep{Wall01} plus its powerful double--peaked maser emission, like other evolved stars, is strong evidence that it is a massive post--main sequence object. The measured $^{12}$C/ $^{13}$C ratio \\citep{Ner89} however is intermediate between what is measured for normal M supergiants and what is expected for the ISM. At a distance of 1.5 kpc \\citep{GH72,LR78, Mar97} it is one of the most luminous red supergiants known ($L \\approx 5 \\times 10^{5} L_{\\odot}$). With its very visible asymmetric nebula 10$\\arcsec$ across and high mass loss rate  ($2 - 3 \\times  10^{-4} M_{\\odot} yr^{-1}$) \\citep{Dan94}, VY CMa is also one of the most important for understanding the high mass loss episodes in massive star evolution.  Interferometric maps of the OH and H$_{2}$O emission reveal a complex structure interpreted as either an expanding disk or bipolar outflow (Bowers, Johnston, \\& Spencer 1983; Richards, Yates \\& Cohen 1998) with expansion velocities of 30 -- 40 km s$^{-1}$ . Speckle imaging \\citep{Witt98} revealed an asymmetric resolved component 0$\\farcs$1 across, elongated approximately north-south. A geometry consistent with the maser maps and the  infrared images \\citep{Monn99, Smi01} is an axially symmetric model with a slightly  flattened optically thick envelope or an equatorial disk and a bipolar flow/rotational axis aligned roughly northeast-southwest and inclined 15 to 30$\\arcdeg$ to our line of sight \\citep{MB80}.  However, our multi-wavelength HST/WFPC2 images of VY CMa  revealed a complex asymmetric distribution of material with no clear axisymmetric structures (See Figures 1,2, \\& 3 in Smith et al. (2001), and with no evident  disk-like structure in the reflected optical light extending out to 2$\\arcsec$ or more from the star as proposed by Herbig (1970a,b), Efstathious \\& Rowan-Robinson (1990), and Monnier et al (1999). Instead, the circumstellar nebula is dominated by a prominent, nebulous  arc to the northwest and two bright filamentary loops or arcs to the southwest of the star, plus some relatively bright condensations or clumps of dusty knots near the star, all of which are evidence for multiple and asymmetric mass loss events. The apparent random orientations of the arcs suggest that they were produced by localized events not strongly aligned to the star's axis or equator. The models for disks are based on modelling the spectral energy distribution and therefore may be due reflected light from the prominent tail and the bright condensations near the star rather than a true disk. The arcs and loops  also do not have the appearance of continuous outflow streams as suggested by Monnier et al. (1999),  and are too massive to  have been ejected by radiation pressure on dust. The initial ejection must have been induced by other processes. We therefore  speculated that the arcs may be expanding loops caused by localized activity on the star's ill-defined surface, involving magnetic fields and convection. Although this suggestion was  conjectural, large starspots do occur on the surface of $\\alpha$  Ori (Gilliland \\& Dupree 1996; Kluckers et al. 1997; Uitenbroek, Dupree \\& Gilliland 1998; Dupree, Lobel \\& Gilliland 1999)  with accompanying outflows and chromospheric activity \\citep{LD00,LD01}. Non-radial pulsational instability is potentially an alternative ejection mechanism, but the distinction between pulsation and convection  may be blurred in VY CMa.  To learn more about the morphology, kinematics, and possible origin  of VY CMa's complex ejecta, we have obtained long-slit spectra to map the emission and  absorption lines across several structures, including the prominent NW  arc and the two bright filamentary loops  to the southwest. Our results reveal the presence of a complex pattern of velocities in the the strong K I emission line   and a significant velocity gradient across the dusty northwest arc. We find that; \\\\  % $\\bullet$ Most of the gas and dust in the nebula, however,  is nearly stationary or moving very slowly with respect to the star. This is not only surprising but may be difficult to explain for an evolved, massive  star. \\\\  % $\\bullet$ Embedded in this nebula are several discrete structures that are kinematically distinct from the surrounding stationary nebulosity. The NW arc and the two prominence--like arcs are expanding at 40 -- 50 km s$^{-1}$ relative to the star. \\\\  % $\\bullet$ They were ejected in separate events over several hundred years from different locations on the star. \\\\  % $\\bullet$ A strong absorption line in the reflected spectrum in the ejecta is identified with Ca I]. Although absorption from a semi-forbidden line seems unlikely, we demonstrate that the identification is probably correct.\\\\ The strong K I emission in the star and the ejecta has never been adequately explained. In the Appendix, we discuss a possible mechanism for forming the K I emission and other rare emission lines, including Ca I].   In the next section we briefly describe the observations and the data reduction procedures.  The spectrum of the central star and the origin of its  peculiar emission lines are discussed in \\S 3 and in the accompanying Appendix. The kinematics of the gas and dust, the expansion of the arcs and loops, and the evidence for stationary ejecta  are described in \\S 4.  In the final two sections we summarize the evidence for localized high mass loss events and comment on the perplexing nature of VY CMa. ",
        "conclusions": "Our observations confirm that  VY CMa has experienced several mass loss events, presumably from localized active regions. These clouds or loops of gas and dust are now expanding outward into what appears to be a surrounding nebula of surprisingly stationary gas and dust. This stationary or slow moving material material is difficult to explain around a star that has had a long history of high mass loss. We can only  speculate that the more diffuse, uniformly distributed gas has been created by a continuous, but slow general outflow of material from the overall surface. It therefore  seems  likely that the high mass loss for rate for VY CMa is really a measure of the mass carried out by these specific ejections perhaps accompanied by streams or flows of emitting gas through holes or gaps in the dust envelope.  VY CMa is the only M--type  supergiant known to have a large, visible reflection nebula. Schuster \\& Humphreys (2005)  have obtained HST/WFPC2 images of several cool hypergiant stars including the OH/IR M supergiants NML Cyg, S Per and VX Sgr to search for associated nebulosity. The results are basically negative or yield only marginal detections, except for NML Cyg which has a small bean-shaped nebula that has probably been shaped by the strong ionizing winds from the nearby hot stars in the Cygnus X complex \\citep{MJ83}. In many ways VY CMa is most like IRC +10420, another high mass--loss, strong OH/IR source believed to be in a post--red supergiant stage and with a complex circumstellar environment \\citep{TJJ,RMH97}. Its extended reflection nebula has a variety of structures including knots and condensations, ray-like features and numerous arcs and scalloped loops which in the outer parts closely resemble the arcs observed in VY CMa's ejecta. Spectroscopy of its ejecta though showed a uniformly expanding outflow of gas \\citep{RMH02}; there was no evidence for stationary gas or dust in its ejecta.  The referee has suggested that the slow moving gas is due to gravitationally bound material in circular orbits (Jura and Kahane 1999). This may be the case but the examples of orbiting material about evolved stars come from a very different class of objects; carbon-rich red giants and post AGB stars which are binaries or presumed binaries (Jura, Bahm, and Kahane 1995, Bujarrabel et al. 2003) with ``reservoirs'' of long -- lasting molecular material. At the distances where the motionless K I emission is observed (4000 - 10000 AU from the star), the orbital period will be several hundred thousand to a million years, exceeding the expected duration of the red supergiant stage for a star this luminous. Furthermore, the OH and H$_{2}$O maser observations show that the clouds are expanding at 30 -- 40 km s$^{-1}$, greater than the expected escape velocity from the dust shell. Therefore we think this explanation is less likely for VY CMa if it is an evolved red supergiant.  VY CMa is  a very perplexing star experiencing discrete mass loss episodes that may be the primary mechanism by which it is losing so much mass. We have also uncovered some puzzling  questions about the origin of its extended diffuse nebula with the slow moving gas and dust. If it is an evolved massive star, it is apparently in a unique stage, with no obvious counterparts. Alternatively, the apparently ``stationary''  ejecta with embedded streams of gas and massive outflows is also consistent with a pre-main sequence state for VY CMa."
    },
    "0410/astro-ph0410485_arXiv.txt": {
        "abstract": "The limited angular resolutions and sensitivities historically available below 300 MHz have made it difficult to define the low end of the electron energy distribution, N($\\gamma$).  We extrapolate down from the well observed segments of radio spectra with almost complete ignorance of what N($\\gamma$) is actually doing.  We do not know if there is a low energy cutoff or if there are other deviations from extrapolated power laws.  The result is that we really do not have a good estimate of the total energy density and pressure of the relativistic plasmas we study.  The situation is even worse for Inverse Compton (IC) X-ray emission, several flavors of which rely on electrons of Lorentz factors, $\\gamma$, of 1000, 300, or in some cases of order 50.  If our assumed extrapolations are wrong, some IC emission models may have to be abandoned.  We present several examples and demonstrate that the Long Wave Array (LWA) should have sufficient sensitivity and resolution to obtain meaningful constraints on N($\\gamma$) at low energies. ",
        "introduction": "At this meeting, we are honoring a man who has not lost sight of his life-long passion for low frequency radio astronomy.  Erickson's exploits are legendary and described in detail in other contributions in this volume.  We note that twenty years ago, November 1984, Erickson and Cane organized a Green Bank Workshop on Low Frequency Radio Astronomy and that 10 of the 43 participants of the '84 workshop managed to survive the intervening 20 years and are registered at this meeting.  In section~\\ref{sec:cl} we review what sort of science we were doing at that time, from the point of view of a guest observer with the TPT at the Clark Lake Radio Observatory.  In addition to celebrating Erickson's birthday, we have come to evaluate and nourish the LWA proposal.  In my opinion, if the LWA does nothing more than tell us something new about the low end of relativistic electron spectra, it will have been worth the effort! Thus, in section~\\ref{sec:ne} we review some of the reasons why these low energy electrons are important.  Briefly, most of the energy density resides below $\\gamma\\approx$5000, i.e. the low end of the electron energy distribution which is generally not sampled by radio observations with arcsec synthesized beams.  In section~\\ref{sec:ic}, we focus on several of the IC processes proposed for X-ray emission which rely on the presence of electrons with $\\gamma\\leq$~1000.  What is normally not discussed in the literature is the extrapolation required from the 'radio-sampled' section of N($\\gamma$) to the energies required for the particular IC model being proposed.  We consider various examples of IC X-ray sources and show what LWA resolutions and frequency coverage could do to mitigate the uncertainty of the current extrapolations.  We use the standard definition of the spectral index, $\\alpha$: flux density,\\\\ S$_{\\nu}\\propto\\nu^{-\\alpha}$. ",
        "conclusions": "\\begin{itemize}  \\item{Will LWA data provide constraints on acceleration mechanisms?}  Yes - to the extent that we can determine $\\gamma_{min}$  \\item{Will these data test any current IC emission models?}  Yes - they will allow us to confirm or to disallow IC/CMB with beaming.  If we can demonstrate a low E cutoff, this model will become untenable .  The standard interpretation of IC/CMB X-rays from radio lobes will also be tested.  \\end{itemize}  The focus of this paper has been on the uncertainty involved in extrapolating electron spectra to low energies, and the ramifications of discovering that our assumed extrapolations may be wrong.  A different outcome however, would be that LWA observations convince us that instead, the extrapolations are valid and that there is little or no change in the observed power law down to low values of $\\gamma$.  In that case, the IC X-ray data will provide measurements of the amplitude of N($\\gamma$) for lobes and clusters of galaxies, and constraints on the amplitude and $\\Gamma$ combination for jets.   For fairly clean examples of IC/CMB from lobes like 2038-272 and 3C351, we see that data from the LWA will permit us to sample the electron spectra at the same energies that we observe with Chandra via IC X-rays.  The IC data provide the amplitude, k$_e$, of N($\\gamma$) for various values of $\\gamma$ below 1500.  The synchrotron degeneracy between B and $k_e$ will thus be resolved and we will then have a good estimate of $<B>$ {\\itshape without the uncertainty of any extrapolation}.  Furthermore, our estimates of the total energy (in fields and particles), the non-thermal pressure, and the energy density will be much more robust and thus we should begin to answer the question of what is the value of $\\gamma_{min}$.  The equipartition constraint will have one less unknown, and pressure balance arguments will then give us constraints on the filling factor and the (1+k) factor (contribution from relativistic protons)."
    },
    "0410/astro-ph0410166_arXiv.txt": {
        "abstract": "Here we present the New York University Value-Added Galaxy Catalog (NYU-VAGC), a catalog of local galaxies (mostly below $z \\approx 0.3$) based on a set of publicly-released surveys matched to the Sloan Digital Sky Survey (SDSS) Data Release 2. The photometric catalog consists of 693,319 galaxies, QSOs and stars; 343,568 of these have redshift determinations, mostly from the SDSS. Excluding areas masked by bright stars, the photometric sample covers 3514 square degrees and the spectroscopic sample covers 2627 square degrees (with about 85\\% completeness). Earlier, proprietary versions of this catalog have formed the basis of many SDSS investigations of the power spectrum, correlation function, and luminosity function of galaxies. Future releases will follow future public releases of the SDSS. The catalog includes matches to the Two-Micron All Sky Survey Point Source and Extended Source catalogs, the IRAS Point Source Catalog Redshift Survey, the Two-degree Field Galaxy Redshift Survey, the Third Reference Catalog of Galaxies, and the Faint Images of the Radio Sky at Twenty-centimeters survey.  We calculate and compile derived quantities from the images and spectra of the galaxies in the catalogs (for example, $K$-corrections and structural parameters for galaxies). The SDSS catalog presented here is photometrically calibrated in a more consistent way than that distributed by the SDSS Data Release 2 Archive Servers and is thus more appropriate for large-scale structure statistics, reducing systematic calibration errors across the sky from $\\sim 2\\%$ to about $\\sim 1\\%$. We include an explicit description of the geometry of the catalog, including all imaging and targeting information as a function of sky position. Finally, we have performed eyeball quality checks on a large number of objects in the catalog in order to flag errors (such as errors in deblending). This catalog is complementary to the SDSS Archive Servers, in that NYU-VAGC's calibration, geometrical description, and conveniently small size are specifically designed for studying galaxy properties and large-scale structure statistics using the SDSS spectroscopic catalog. ",
        "introduction": "\\label{motivation}  New, large galaxy datasets such as the Sloan Digital Sky Survey (SDSS; \\citealt{york00a}) and the Two-Micron All Sky Survey (2MASS; \\citealt{skrutskie97a}) give astronomy a view of the local universe with an unprecedented combination of completeness and detail. These surveys promise to refine our understanding of properties of galaxies, their relationships with environment, and their evolution.  However, these catalogs consist of terabytes of data and are therefore currently too unwieldy simply to download onto a small workstation and investigate directly. Furthermore, simple matched catalogs between many surveys exist only in the form of web interfaces such as NED\\footnote{{\\tt http://nedwww.ipac.caltech.edu/}}.  As indispensable as such interfaces are for studying a small number of individual objects on which one desires all the knowledge in the literature, they are not ideal for huge batch jobs designed to run automatically on relatively homogeneous data sets. Finally, none of the databases interfacing to these huge catalogs yield any expression of the window function on the sky of the included surveys, leaving this critical determination to the user.  For these reasons, we have created a galaxy redshift catalog designed to aid in the study of the local universe. Early versions of this catalog have proven invaluable to the SDSS team, as they form the basis for the work studying the luminosity function, power spectrum, correlation function, and a number of other galaxy property and galaxy clustering statistics; {\\it e.g.} \\citet{blanton03c, tegmark04a, zehavi02a, hogg03b, hoyle03a, shen03a, baldry04a, pope04a}.  Non-SDSS investigators have also used public data extracted from this catalog to investigate galaxy properties and their evolution; {\\it e.g.} \\citet{trujillo04a, rudnick03a}.  The catalog is small enough (tens of Gbytes) to easily store on any modern machine, it contains matches between many major public catalogs, and it contains an explicit description of its window function. The current version of the catalog uses the SDSS imaging survey as its basis, matching sources in other catalogs to that master list, and only tracking the SDSS geometrical information. In the future we plan on incorporating other surveys more fully by explicitly including their geometrical information.  For most of the catalogs in the sample, we simply use the official public releases. However, we give the SDSS (of which the authors are participants) a special status in this catalog. First, it is the only survey for which we track detailed information about the window function.  Second, for the SDSS we use an independent set of reductions of the public data (Schlegel et al. in preparation).  These reductions are significantly improved over the reductions made available through the Data Archive Server (DAS) or Catalog Archive Server (CAS) on the SDSS Data Release 2 (DR2; \\citealt{abazajian04a}) web site\\footnote{{\\tt http://sdss.org/dr2/}}.  The essential improvement is that errors in the large-scale relative photometric calibration of the survey have a lower variance and less structure, making the catalog more appropriate for studies of large-scale structure (as we demonstrate below). In addition, these reductions include an explicit expression for the survey geometry which the Archive Servers do not. Secondarily, the naming and unit conventions in our catalog are different, as we describe in this paper. For these reasons, expect the quantities for objects in this catalog and the corresponding objects in the Archive Server catalogs to be very similar but not identical.  This paper describes only the outline and general principles of the catalog. We leave the detailed documentation to a set of regularly updated web pages listed in Table \\ref{webpages}.  The paper contains: in Section \\ref{window}, a short description of the geometrical expressions used in the catalog and how to use them; in Section \\ref{constituents}, a description of the constituent catalogs and how they are matched, focusing on the SDSS; in Section \\ref{other}, a description of other derived quantities; in Section \\ref{lowz}, the presentation of a low redshift galaxy catalog based on this data; in Section \\ref{tools}, a short desription of some of the tools we have used to create this dataset which might be useful to the user as well; and in Section \\ref{summary}, a summary.  Note that throughout we use the environmental variable \\vagcdir{} to denote the root location of the NYU-VAGC; see the online documentation listed in Table \\ref{webpages} for the actual root location. Similarly, for the large-scale structure samples we describe below, we use the environmental variable \\lssdir{}. ",
        "conclusions": "\\label{summary}  Here we have presented a catalog of galaxies combining information from SDSS, 2MASS, FIRST, PSC$z$, RC3, and 2dFGRS. The main improvements of this catalog over the standard SDSS release (in addition to the matches to other catalogs) are a better calibration and an explicit description of the geometry, including all of the information necessary to perform large-scale structure analyses with the catalog. We have also included structural measurements, $K$-corrections, peculiar velocity corrections, and quality checks of many objects. The catalog is fully documented on the web site listed in Table \\ref{webpages}.  As we continue to develop the NYU-VAGC, we expect to add considerable functionality. For example: \\begin{enumerate} \\item We plan to extract images and, where possible,  measure structural parameters for galaxies detected in other surveys which overlap the SDSS. \\item We plan to include parameters for NYU-VAGC galaxies from other surveys ({\\it e.g.} the Spitzer Space Telescope and the Galaxy Evolution Explorer). \\item We plan to continue adding SDSS data as it is released, and improving the treatment of deblended galaxies. \\item The SDSS ubercalibration procedure will continue to improve, as the SDSS collaboration takes more and more overlapping data. \\end{enumerate}  We note that this project will be more successful if users provide feedback about how the catalog could be improved, since we do not expect that we can predict from pure thought what astronomers will find useful."
    },
    "0410/astro-ph0410600_arXiv.txt": {
        "abstract": "We present a new analysis of the expected magnetospheric radio emission from extrasolar giant planets for a distance limited sample of the nearest known extrasolar planets. Using recent results on the correlation between stellar X-ray flux and mass-loss rates from nearby stars, we estimate the expected mass-loss rates of the host stars of extrasolar planets that lie within 20~pc of the Earth. We find that some of the host stars have mass-loss rates that are more than $100$ times that of the Sun, and given the expected dependence of the planetary magnetospheric radio flux on stellar wind properties this has a very substantial effect.  Using these results and extrapolations of the likely magnetic properties of the extrasolar planets we infer their likely radio properties.  We compile a list of the most promising radio targets, and conclude that the planets orbiting Tau Bootes, Gliese~86, Upsilon Andromeda and HD\\,1237 (as well as HD\\,179949) are the most promising candidates, with expected flux levels that should be detectable in the near future with upcoming telescope arrays. The expected emission peak from these candidate radio emitting planets is typically $\\sim 40-50$~MHz. We also discuss a range of observational considerations for detecting extrasolar giant planets. ",
        "introduction": "\\label{sec1}  The reasons for searching for and studying extrasolar giant planets (EGPs) at radio wavelengths are many. First, to provide another means of direct detection of EGPs, second, if a detection with sufficient strength is made, it will provide evidence of the presence and indeed strength of the magnetic field of the EGP. Many current planet-finding techniques are indirect in nature, relying on seeing the effect of the planet on the host star (such as the Doppler, transit, astrometric or microlensing methods). In the future, direct detection of the planet may be made using space-based interferometric techniques, or through coronographic techniques. Radio detection provides another possible means of directly observing the planet (or the planet's near environment).  In principle, the detection of radio emission could also provide information on the existence of any satellites around the EGP, the rotation period, and indeed provide constraints on the inclination of the EGP (though these last reasons are very speculative and will require high quality observations). For the time being, mere detection of EGPs at radio wavelengths is a sufficiently worthwhile goal.  The solar system object that leads to the expectation that EGPs may be detectable at radio wavelengths is Jupiter (and, to a lesser extent, the other planets that have been detected at radio wavelengths, namely the Earth and the other Jovian planets -- see Bastian, Dulk \\& Leblanc 2000). At decameter wavelengths, Jupiter is an extremely bright source, detectable with very modest equipment. At these wavelengths the Sun/Jupiter contrast can be $\\sim 1$ and indeed the polarised nature of the emission from Jupiter (and we expect from EGPs) will allow discrimination as to whether the emission is planetary or stellar in origin. We shall discuss the Jovian radio emission in Section~\\ref{sec1p1}.  The radio emission from the solar system planets is related to their magnetospheres, and a radio detection of an EGP will give direct evidence of a magnetosphere. However, indirect evidence for a planetary magnetosphere has already come to light in the case of HD\\,179949\\,b (Shkolnik, Walker \\& Bohlender 2003), where variations in the Ca{\\small II} H and K lines seem to be phase locked with the orbit of the planet (which has a period of 3.093~days and a mass of $M_p\\sin i = 0.98M_{Jup}$). A number of other stars with short period EGPs were also studied in the same way with no other comparable results.  HD~179949 is more distant than the sample covered in this paper (at a distance of $D=27$~pc), but because of the inferred existence of a magnetosphere we will include it in the sample discussed below.  \\subsection{Origin of the Jovian Radio Emission} \\label{sec1p1}  In the magnetospheres of planets, intense radio emission can result from electron cyclotron maser radiation. In this process keV electrons in the auroral regions of the planet radiate at the gyro-frequency of the magnetic field lines. The radiation is emitted in a hollow cone region, and is highly structured in time and frequency.  The requirements for this process are simply a magnetic field and a source of energetic (keV) electrons. The presence of a magnetic field will be extremely likely in EGPs, and we would expect the magnetic moment of the more massive EGPs to exceed that of Jupiter (Section~\\ref{sec2p4p1}). The energetic electrons can be provided by a range of means (in the case of solar system objects); the solar wind, auroral processes or planet-moon magnetic coupling (in the case of Io/Jupiter) and there is no reason to expect this to be different for EGPs.  The electron-cyclotron maser emission is seen to be very sporadic in the case of Jupiter and the other planets, and the emission is 100\\% circularly or elliptically polarised. For a discussion of the extremely variable S-burst decametric emission and the rather more slowly varying L-type emission, see Queinnec \\& Zarka (2001).  In the case of Jupiter, the peak emission occurs at a frequency of $\\sim 10-20$~MHz (though it is seen to extend up to $\\sim 40$~MHz), while for the other magnetised planets it is at correspondingly lower frequencies, in the range of $0.1-1$~MHz. For planets other than Jupiter, this emission is not observable from the ground (with the ionospheric cut-off typically in the range of $\\sim 2-15$~MHz, depending on the location on Earth, day or night, location in the solar cycle etc), and were detected by space-based instruments. Jupiter has also been detected as a radio source at higher (GHz) frequencies, but at a much lower level.  One further constraint, is that the local plasma frequency in the source region must be lower than the gyro-frequency, which is the case for most solar system objects. The exception to this is probably Mercury, where the local plasma frequency in the solar wind exceeds the gyro-frequency (the surface magnetic field at the equator is $B=0.0033$~G). As discussed in Section~\\ref{sec3}, there will also be issues with this associated with planets with stronger magnetic fields immersed in more dense stellar winds, or orbiting at very close distances to the host star.  For an overview of the emission mechanisms and the radio properties of the solar system planets see Bastian \\etal\\ (2000), and references therein.  \\subsection{Previous Searches for EGP Radio emission} \\label{sec1p2}  There have been a small number of reported searches (or related proposals for searches) for radio emission from extrasolar planets, with no detection as yet. The searches reported include:  \\begin{enumerate} \\item Winglee, Dulk \\& Bastian (1986) used the Very Large Array ({\\it VLA}) to observe six nearby stars at 333 and 1400~MHz. These objects are not part of the current list of EGP hosts (with the possible exception of Lalande~21185). None were detected.  \\item Bastian \\etal\\ (2000) reported on {\\it VLA} observations of a range of EGP and brown dwarf candidates at 333~MHz and 1465~MHz, with a smaller number of observations at 74~MHz. The relevant objects covered included 51~Peg, Ups~And, 55~Cnc, 47~UMa, Tau~Boo, 70~Vir and Lalande~21185. These object were all observed at 333~MHz and 1465~MHz, while 47~UMa was also observed at 74~MHz. Of these, the 74~MHz observations are the most relevant and the {\\sl rms} noise level for the non-detection the 47~UMa observation was quoted as 76~mJy for the entire observing run and 3.24~Jy for a 10~sec period (important if the emission is extremely variable, which if Jupiter is anything to go by, is likely to be the case).  \\item Farrell \\etal\\ (2003) observed Tau~Boo with the {\\it VLA} at 74~MHz, with a quoted upper flux limit of 0.12~Jy (see also Farrell, Desch \\& Zarka 1999 for an earlier analysis of the likely radio emission from Tau~Boo).  \\item Ryabov, Zarka \\& Ryabov (2003) reported on progress on observations of EGPs, using the UTR-2 array in the Ukraine. Observations of 20 nearby EGPs have been undertaken with no detections reported so far.  \\item While not a search as such, Butler (2003) reports on the prospect of radio detection of EGPs with upcoming instruments such as the Square Kilometer Array ({\\it SKA}) or the Low-Frequency Array ({\\it LOFAR} \\footnote{For a description of the proposed {\\it LOFAR} instrument see http://www.lofar.org}).  \\end{enumerate}  We shall discuss these observational limits on the radio emission from EGPs, in the light of the new predictions presented here, in Section~\\ref{sec4}.  The rest of the paper is set out as follows: in Section~2, the basic theory of radio emission from EGPs is introduced, including a discussion of the influence of the planetary parameters and characteristics of the stellar winds of the host star on magnetospheric radio emission from EGPs; in Section~3, the results from the modelling are presented, and in Section~4, these results are discussed, along with comments on the likely observability of the radio emission from EGPs. ",
        "conclusions": "\\label{sec4}  In this paper we have presented a critical discussion of the expected radio properties of EGPs in the light of both the planetary and stellar characteristics of the host star. As a result of this we can make the following statements about the likely levels of EGP radio emission.  \\begin{enumerate}  \\item Only very nearby EGPs are likely to be detected by their radio emission. The $1/D^2$ fall-off makes the problem very hard. There are only 28 EGPs currently known within 20~pc (orbiting 22 different host stars) and only 3 within 10~pc (one orbiting Eps Eri and two orbiting Gliese~876).  \\item The prime factor in determining whether an EGP is radio bright is its orbital period. Shorter period planets will also show more radio emission, with the radio flux scaling as $A^{-4/3}$. However, this effect will not continue indefinitely as if the planet is very close to the star the incident wind will impinge on the magnetosphere at lower velocities. Also, eventually plasma-frequency screening will begin to play a role for very short period planets.  \\item Higher levels of radio emission are likely to be associated with more massive planets, with the radio flux scaling as $M_p^{2/3}$. The frequency of the radio emission from the more massive planets will likely be shifted to more amenable wavelengths. We note that the most massive planet within 20~pc is HD\\,39091\\,b (with $M_p\\sin i=10.35M_{Jup}$), though we do not expect this planet to be a bright radio source, as it is in a very long period orbit.  \\item Higher levels of EGP radio emission are expected from stars with higher coronal X-ray surface fluxes (and hence higher mass-loss rates). Of the nearby stars (within 20~pc) a fair number have been detected at X-ray wavelengths and have mass-loss rates in the range of $\\dot M_\\ast = 0.1-100 {\\dot M_\\odot}$. The nearest star harbouring an exoplanet (Eps Eri) does have a strong wind, but also a long period planet, which offsets this. Tau~Boo, which remains the most promising candidate, also has a very strong wind, with $\\dot M_\\ast=83\\dot M_\\odot$, although HD\\,1237 is another promising candidate, has $\\dot M_\\ast=86\\dot M_\\odot$.  \\item The peak frequency of emission in the five brightest cases range from $\\sim 10$~MHz up to 50~MHz, though these values are uncertain. Certainly, very low frequency observations will be required for a detection, with even the 74~MHz {\\it VLA} observations possibly being at too high a frequency. As with Jupiter, we do expect emission from the EGPs at $\\nu>\\nu_{peak}$, but it will likely be falling off very rapidly.  \\item Two of the five predicted brightest systems (Tau~Boo and Gliese~86) are in binary stellar systems (there seems to be a dearth of massive planets in short period orbits in single stellar systems and the apparent opposite trend in binary systems, see Zucker \\& Mazeh 2001).  \\item In some cases, orbital eccentricity may play a major role in detectability of the system. The prime example from this study is HD\\,1237, where the expected flux will vary by a factor 4.5 over the 134 day orbit, reaching over 20~mJy.  \\end{enumerate}  On the basis of the results presented here, there remains the prospect of detecting nearby EGPs (particularly with {\\it LOFAR} or {\\it SKA}). The best candidates (in order of expected brightness), with their predicted peak frequencies) are:  \\begin{enumerate} \\item Tau Bootes (48~MHz) \\item Gliese 86 (46~MHz) \\item HD\\,179949 (11~MHz) \\item Upsilon Andromeda (8~MHz) \\item HD\\,1237 (40~MHz) \\end{enumerate}  Previous searches for EGPs (see Section~\\ref{sec1p2}) have tended to focus on rather higher frequencies than these (with the 74~MHz observations at the {\\it VLA} being the most sensitive). The emission from Jupiter falls off very sharply with increasing frequency, which means that careful (multi-frequency) observations are needed.  On a positive note, all five of these objects have fluxes that are potentially detectable, particularly with upcoming {\\it LOFAR} telescope, which will operate at lower frequencies, more appropriate for detecting EGPs.. The quoted sensitivities for {\\it LOFAR} (for a 1 hour integration with the full array, single polarisation, 4~MHz bandwidth), are 3~mJy at 10~MHz, 1.6~mJy at 30~MHz and 1.0~mJy at 75~MHz (see http://www.lofar.org). However, not all of them are visible from the anticipated {\\it LOFAR} site (HD\\,1237 and Gliese~86 are not -- Gliese~86 is just visible from the {\\it VLA}). The other 3 targets are visible from this site.  As a further consideration, it is notable that Shkolnik \\etal\\ (2003) also detected night-to-night activity on Tau~Bootes and Upsilon Andromeda, as well as HD\\,179949 (though in these cases the synchronicity with the planet was not obvious). This activity was not present in the standard stars observed, nor in the 51~Peg system. This activity could also be related to magnetospheric activity.  There are however, many uncertainties in this analysis, such as the extrapolations from our own solar system to extrasolar systems, in particular how the radio flux from short period systems will depend on the stellar/planetary parameters. In addition, we have assumed a scaling for the magnetic moment with mass. The close-in planets may well be rotationally locked to their host star. What this lack of rotation will do to any planetary dynamo and magnetic field is uncertain, but is likely to reduce it. If this turns out to be the case then HD\\,1237 is perhaps the most promising target, with a reasonably long period planet (and likely interesting orbital variability).  As a further consideration, the expected radio bright EGPs will preferentially be orbiting young, rapidly rotating and X-ray luminous stars (with high mass-loss rates). This means these planets themselves will be young too.  Given that it may take planets a reasonable length of time to generate a self-sustaining magnetic field, it is possible that this may limit the radio brightness of EGPs, with increasing magnetic field strength of the EGPs being coupled with decreasing mass-loss rate of the host star. Further study of the likely temporal evolution of magnetic fields in EGPs is needed.  We have also assumed that the host stars are not radio emitters at similar wavelengths, but we do note that radio flare stars may well confuse the issue in some cases. This may be an issue for some M stars, but is much less likely to be an issue for F, G and K-stars (Bastian 1990). In addition, while the winds of massive stars can also generate both thermal and non-thermal radio emission, the winds considered here are unlikely to be strong enough to generate any confusing emission.  In addition to the expected flux, the optimum frequency to observe these sources is also uncertain.  The predicted peak frequencies are around 50~MHz for Tau~Boo and Gliese~86 (which is still at a lower frequency than the 74~MHz observations of Tau~Boo), and rather lower for HD\\,179949 and Ups~And.  The peak frequency is, however, very uncertain, and depends on the planetary magnetic moment (and planetary mass), which for all of these objects is uncertain due to the $\\sin i$ factor (and the assumptions made in this model). We have assumed a mean value of the inclination, and it could be for these objects that $i$ is closer to $90^\\circ$, which will reduce the expected flux as well as making the peak frequency difficult to observe (though the converse is also possible, making the situation easier). The peak frequency also depends on the planetary radius, which we may be underestimating for the very close planets (those with orbital separations of 0.05~AU).  Gliese~86\\,b has an orbital separation of 0.11~AU, putting in the category of a class~IV EGP, and so, while it is possible we are underestimating the planetary radius, the effect is likely to be less than for Tau~Boo.  On a positive note, given the anticipated parameters of the {\\it LOFAR} instrument, from the estimates presented here, there should be ample sensitivity to detect a substantial number of EGPs. These EGPs will have a range of period/mass characteristics, and thus the observations will initiate the observational study of extrasolar magnetospheres.  So, in summary, we have performed a detailed analysis of the expected radio emission from nearby extrasolar giant planets. In particular, new results allow us for the first time to be quantitative about the expected stellar wind properties of the host stars of EGPs, which in turn has important consequences on their radio detectability. We predict that a number of EGPs will emit at radio wavelengths at levels that are likely to be observable with {\\it LOFAR}. We do note that there remain a number of issues that could reduce the observability of these objects, and care should be taken when choosing to observe the candidates, but nonetheless there does remain the prospect of a detection and we have proposed the five most promising targets on which to undertake a search."
    },
    "0410/astro-ph0410436_arXiv.txt": {
        "abstract": "The radio emission during 201 X-ray selected solar flares was surveyed from 100 MHz to 4 GHz with the Phoenix-2 spectrometer of ETH Zurich. The selection includes all RHESSI flares larger than C5.0 jointly observed from launch until June 30, 2003. Detailed association rates of radio emission during X-ray flares are reported. In the decimeter wavelength range, type III bursts and the genuinely decimetric emissions (pulsations, continua, and narrowband spikes) were found equally frequently. Both occur predominantly in the peak phase of hard X-ray (HXR) emission, but are less in tune with HXRs than the high-frequency continuum exceeding 4 GHz, attributed to gyrosynchrotron radiation. In 10\\% of the HXR flares, an intense radiation of the above genuine decimetric types followed in the decay phase or later. Classic meter-wave type III bursts are associated in 33\\% of all HXR flares, but only in 4\\% they are the exclusive radio emission. Noise storms were the only radio emission in 5\\% of the HXR flares, some of them with extended duration. Despite the spatial association (same active region), the noise storm variations are found to be only loosely correlated in time with the X-ray flux. In a surprising 17\\% of the HXR flares, no coherent radio emission was found in the extremely broad band surveyed. The association but loose correlation between HXR and coherent radio emission is interpreted by multiple reconnection sites connected by common field lines. ",
        "introduction": "`Flares are highly individualistic', already noticed Smith \\& Smith (1963) reporting on optical observations. They recapitulated the first century of flare observations and were impressed by the different morphology of each H$\\alpha$ flare. Today's observers of the non-thermal aspects of the flare phenomenon, particularly of the hard X-rays (HXR) and decimeter radio emissions, have a similar experience. The pattern of the various emissions in the radio spectrogram combined with the X-rays seems to be equally disparate in each flare.  In meter waves, the classification into 5 types has been sufficient to characterize, interpret and model the emissions. A comprehensive review of the meter-wave emissions may be found in a volume edited by McLean \\& Labrum (1985). The situation is not clear, however, in decimeter waves where in principle a similar number of types exists but seems to represent still largely morphological classes in spectrograms rather than physical phenomena. The decimetric radio emission has been classified by G\\\"udel \\& Benz (1988) and Isliker \\& Benz (1994) into type III-like bursts, pulsations, diffuse continua, narrowband spikes, decimetric type IV bursts, and high-frequency continuum generally attributed to gyrosynchrotron emission. Except for the last one, these classes are rather broad and may well include emissions originating by different processes. Each includes several subclasses or various fine structures. The above 6 broad classes have basically been confirmed by other observers including Allart et al. (1990) and Bruggmann et al. (1990) at frequencies up to 8.5 GHz, and enriched by further subclasses (e.g. Jiricka, Karlicky, \\& Meszarosova, 2002). The details of the various burst classes and subclasses are not reviewed here. The reader is referred to the original literature quoted and to the review by Bastian, Benz, \\& Gary (1998).  Individual radio burst types have been studied extensively for the association with HXRs. Kane (1981) notes that only 3\\% of meter-wave type III bursts are correlated with HXR emission, but the correlation increases with higher start frequency of the bursts. Dennis et al. (1984) and Aschwanden et al. (1985) find occasional associations between reverse-drift type III bursts with HXR peaks. Aschwanden et al. (1995) took a large sample of HXR peaks and find associated type III bursts around 300 MHz in 31\\%  of all cases. Some cases of correlation between the start or enhancement of a type I bursts (also called `noise storm') and soft X-rays were reported by Lantos et al. (1981), Raulin \\& Klein (1994) and Willson, Kile, \\& Rothberg (1997). HXR emission was found coincident with only 20\\% of all decimetric pulsations by Aschwanden, Benz, \\& Kane (1990). However, the more irregular subtype of drifting pulsations was reported as well associated in several recent case studies (Kliem, Karlick\\'y, \\& Benz, 2000; Kahn et al., 2002; Saint-Hilaire \\& Benz, 2003). Finally, 95\\%  of all decimetric spike events were found associated with HXR emission, but only 2\\% of the HXR events were associated with spikes (G\\\"udel, Aschwanden \\& Benz, 1991).  General surveys on coherent radio emissions during X-ray flares are rare. Aschwanden et al. (1985) used observations by ISEE-3 in the two channels 26 - 43 keV and 5.8 - 6.9 keV. The radio data consisted of analog recordings on film and were limited to the 100 - 1000 MHz range. The prime list was selected from radio events. Simnett \\& Benz (1986) did the opposite and selected strong HXR events ($>$ 1000 cts/s at $>$ 25 keV) from the HXRBS/SMM list. They found no radio emission in 15\\% of the events, using data from the same analog spectrometer.  Still missing is a broadband survey of radio emissions during `X-ray selected' flares. X-ray emission is the widely used measure for flare energy and thus selects the more energetic flares. The key questions are: Which flares emit what type of radio emission and in which phase of the flare? What are the general patterns for impulsive HXR flares?  In this study, two large and homogeneous data bases of X-ray and radio observations are used to address the question of HXR-radio association. The RHESSI detector resolves well the 10 - 20 keV range where most of the flare photons are detected. It is the first time that the Phoenix-2 radio spectrometer was used to survey its entire possible band from 100 MHz to 4 GHz for an extended period. The band covers the entire frequency range of the plasma frequency and electron cyclotron frequency in the regions where the flare energy is thought to be released. The goal is not only to identify all bursts in the meter and decimeter spectrogram or in HXR lightcurves, but patterns in the combined data set. Such patterns are needed for a fresh look at the decimeter burst emission processes. Although many decimeter burst emission mechanisms have been proposed (review by Benz, 2002), few are supported by further observations. Thus the investigation of decimeter flare emissions must be considered to be still in an exploratory phase.  The ultimate goal of the comparison with HXR emission is to make coherent radio emissions useful diagnostic tools for the understanding of the flare processes. Thus the strategy of this survey is to start from HXRs as a proxy for energy release and particle acceleration. ",
        "conclusions": "Hard and soft X-rays indicate the dominant energy release in flares. Radio emissions, on the other hand, have inspired flare physics for the past 50 years. In meter waves, electron beams and shock waves have been first discovered by their radio signatures. Centimeter waves have revealed the presence of relativistic electrons and are essential in present estimates of flare magnetic fields. At decimeter wavelengths, coherent radio emissions have been the starting point for various acceleration theories. Scenarios have emerged in the past decades where the coherent radio emissions fitted nicely into a standard flare model.  This comparison of X-rays with decimeter bursts has confirmed that the two emissions are relatively well associated, but surprisingly revealed them to be much more individualistic than simple scenarios would suggest. RHESSI resolves the low-energy spectrum, where small flares and flare elements can be more easily discovered and compared with radio emissions. Combined with the very broadband radio survey of Phoenix-2 in meter and decimeter wavelengths, the timing of two emissions was compared in a comprehensive way for the first time. Decimetric type III bursts and decimetric pulsations, spikes and continua are found preferentially in the peak hard X-ray phase, but often not simultaneously enough to suggest a common electron population or even a common acceleration process. The fact that HXRs and radio emissions do not correlate in detail suggests that there are other acceleration processes taking place in addition to the energetically dominant acceleration apparent in HXRs. Thus, the survey presented here indicates that the flare phenomenon is not a locally well-defined process but may include several parts of the active region yielding different HXR and radio signatures.  If type III bursts occur, most are found in the rise or peak HXR phase. Their electron beams can be followed in the range of decimetric frequencies, suspected to include the plasma frequency of main flare energy release. Imaging type III emissions at various frequencies will make it possible to trace back these electron beams to their origin. Type III bursts thus have the potential to reveal energy release sites, secondary in amount but primary in time.  The other decimetric emission processes are still unknown. Nevertheless, some of the coherent flare radio emissions are suspected to be indicators for flare electron acceleration as they occur in the rise or peak phase of HXRs. Again, the temporal correlation is not sufficient to identify the radio and X-ray sources in a straightforward way. Therefore, the decimetric source regions do not seem to coincide always with the main flare energy release site nor to correlate with the acceleration process in a quantitative way.  The decimeter flare emissions thus appear to be a largely unexplored field full of unexplained phenomena and unsolved issues of radiation processes. This comparative study has become possible only thanks to the recent developments in data retrieval and analysis tools. Further progress can be expected from new imaging instruments dedicated to solar observations, allowing to compare the radio sources with coronal structures and the flare geometry. Future instrumentation at higher sensitivity may possibly reveal the missing link between radio and X-ray signatures of non-thermal electrons."
    },
    "0410/astro-ph0410746_arXiv.txt": {
        "abstract": "The Atacama Large Millimeter Array (ALMA) consists of 64 antennas of 12\\,m diameter that will initially observe in 4 frequency bands between 84 and 720 GHz with spatial resolutions down to 0\\farcs01 and velocity resolutions as fine as 0.05 km/s. These technical requirements are based on three primary science goals. We illustrate two of these requirements: (i) the ability to detect spectral line emission from a Milky-Way type galaxy at $z$=3, and (ii) the ability to provide precise images at an angular resolution of 0\\farcs1. Finally, we present a possible large extragalactic project with ALMA: molecular line studies of submm galaxies. ",
        "introduction": "ALMA will be a flexible observatory supporting a breath of research in the fields of planetary, galactic and extra-galactic astronomy. The three primary science requirements have been defined as: \\begin{enumerate} \\item The ability to detect spectral line emission from CO or C{\\small I} in a normal galaxy like the Milky Way at a redshift of $z=3$, in less than 24 hours of observation. \\item The ability to image the gas kinematics in protostars and protoplanetary disks around young Sun-like stars at a distance of 150~pc, enabling one to study their physical, chemical and magnetic field structures and to detect the gaps created by planets undergoing formation in the disks. \\item The ability to provide precise images at an angular resolution of 0\\farcs1. Here, the term precise images means representing within the noise level the sky brightness at all points where the brightness is greater than 0.1\\% of the peak image brightness. This requirement applies to all sources visible to ALMA that transit at an elevation greater than 20\\degr. \\end{enumerate}  We now concentrate on the two requirements important for extragalactic astronomy (see J. Richer's contribution for the galactic requirements).  \\subsection{Spectral line imaging of normal galaxies at $z$=3} An estimate of the technical requirements to achieve this science goal can be made by using experience learned from existing millimeter arrays, which have collecting areas between 500 and 1000~m$^2$. These arrays now routinely detect CO emission from high redshift galaxies and quasars \\citep[see ][for a review]{car04a}. Fainter molecular lines such as HCN \\citep[\\eg][]{car04b} and atomic Carbon \\citep{wei04a} are now also within the reach of existing telescopes. These observations take one to two days of total observing time, and are only possible for the most luminous sources and/or with the aid of gravitational lensing. In normal, unlensed galaxies, these lines would be a factor of 20 to 30 fainter.  The sensitivity of ALMA for a given integration time is essentially controlled by three major terms: (1) the atmospheric transparency, (2) the noise performance of the detectors, and (3) the total collecting area.  The location of ALMA on the Chajnantor plateau at an altitude of 5000m will minimize the contribution from the atmosphere compared to existing millimeter observatories. The noise level of the detectors can be reduced be a factor of two, and will then approach the fundamental quantum limit. An important factor of $\\sqrt 2$ will be gained by the requirement that ALMA support front end instrumentation capable of measuring both states of polarization. The remaining factor of 7 to 10 can only be gained by increasing the collecting area. Hence, an ALMA requirement is a collecting area $>$7000~m$^2$. \\vspace{1.5cm}  A more specific calculation of the requirement to have the ability to detect CO emission from a Milky Way type galaxy at $z$=3 is as follows. At cosmological distances, the 10\\,kpc disk of the Milky Way is much smaller than the primary beam of existing millimeter antennas, so a single observation would be sufficient. The flux density sensitivity in an image from an interferometric array can be written as $$\\Delta S = \\frac{4 \\sqrt 2 k T_{\\rm sys}}{\\gamma \\epsilon_q \\epsilon_a \\pi D^2 \\sqrt{n_p \\frac{N(N-1)}{2}\\Delta\\nu \\Delta t}}~{\\rm W m^{-2} Hz^{-1},}$$ where $T_{\\rm sys}$ is the system temperature, $\\epsilon_a$ is the aperture efficiency, $\\epsilon_q$ is the correlator quantization efficiency, $D$ is the antenna diameter, $n_{\\rm p}$ is the number of simultaneously sampled polarizations, $N$ is the number of antennas, $\\Delta \\nu$ is the bandwidth, $\\Delta t$ is the integration time, and $\\gamma$ is a gridding parameter that we set to unity. For ALMA, we shall assume $\\epsilon_q$=0.95, $n_{\\rm p}$=2, and $\\sqrt{N(N-1)}\\simeq N$ so this equation then simplifies to $$\\Delta S= \\frac{2.6\\times 10^6 T_{\\rm sys}}{\\epsilon_a N D^2\\sqrt{\\Delta\\nu \\Delta t}}~{\\rm mJy}.$$  $T_{\\rm sys}$ and $\\epsilon_a$ vary between the different atmospheric bands of ALMA. A surface accuracy of 20\\,$\\mu$m should be achievable, providing aperture efficiencies between 0.75 and 0.45 (see Table~\\ref{table:eff+Tsys}). $T_{\\rm sys}$ depends on several atmospheric and instrumental parameters. Table~\\ref{table:eff+Tsys} lists estimates of the achievable $T_{\\rm sys}$ at an elevation of 50\\degr\\ using state-of-the-art receivers. With these assumptions, one can thus calculate the required collecting area parametrized by the product $ND^2$ to achieve a certain flux density level.  \\begin{figure*}[ht] \\centering \\includegraphics[width=0.5\\linewidth,angle=-90]{atmospheric_transmission.ps} \\caption{Atmospheric transmission at Chajnantor with the ALMA frequency bands indicated. Initially, only bands 3, 6, 7 and 9 will be used .\\label{fig:atmosphere}} \\end{figure*}  \\begin{table} \\begin{center} \\caption{Aperture efficiencies $\\epsilon_a$ and estimated system temperatures $T_{\\rm sys}$ for ALMA at 50\\degr elevation.}\\vspace{1em} \\renewcommand{\\arraystretch}{1.2} \\begin{tabular}[h]{ccc} \\hline Frequency & $\\epsilon_a$ & $T_{\\rm sys}$ \\\\ GHz & & K\\\\ \\hline 35 & 0.75 &   35 \\\\ 110 & 0.74 &   67 \\\\ 230 & 0.72 &  107 \\\\ 345 & 0.69 &  251 \\\\ 409 & 0.67 &  425 \\\\ 675 & 0.54 & 1050 \\\\ 850 & 0.45 & 1150 \\\\ \\hline \\\\ \\end{tabular} \\label{table:eff+Tsys} \\end{center} \\end{table}  The total CO luminosity of the Milky Way in the CO(1--0) transition $L^{\\prime}_{\\rm CO(1-0)}$=3.7$\\times$10$^8$\\,K\\,km\\,s$^{-1}$pc$^2$ has been estimated by \\citet{sol89}. The CO luminosities seen in higher CO transitions by COBE \\citep{ben94,wri91} are slightly higher. In the following, we shall adopt $L^{\\prime}_{\\rm CO(1-0)}$=5$\\times$10$^8$\\,K\\,km\\,s$^{-1}$pc$^2$. From this, we can calculate the expected received flux density following \\citet{sol92} as $$S_{\\rm CO} = 3.08\\times 10^{-8}\\frac{L^{\\prime}_{\\rm CO}\\nu_{\\rm rest}^2(1+z)}{\\Delta v_{\\rm rest}D_L^2},$$ where $S_{\\rm CO}$ is the flux density in Jy, $\\nu_{\\rm rest}$ is the rest frequency of the transition in GHz, $D_L$ is the luminosity distance in Mpc, and $\\Delta v_{\\rm rest}$ is the rest line width in km\\,s$^{-1}$.  In the Milky Way, most of the CO emission arises in clouds of several tens of Kelvin kinetic temperature. In Milky Way-like galaxies at $z$=3, the gas will be somewhat warmer due to the higher background temperature from the cosmic microwave background radiation. At $z$=3, we can observe the J=3--2 or J=4--3 transitions, which fall inside the atmospheric transmission regions. We thus need to consider the expected luminosity of these higher order transitions, using a proper radiative transfer model that takes the higher background temperatures into account. Several such models have been proposed \\citep{silk97,com99,pap00}. However, they reach significantly different conclusions. There have also been a few observations of CO ladders in high redshift quasars \\citep{car02,ber03} and Ultra Luminous InfraRed Galaxies \\citep{wei04b}, all showing CO intensities up to the 6--5 or 7--6 transition. However, the CO emission from quasars and ULIRGs is likely to be dominated by the central component, while our requirement is to detect the outer regions in Milky Way type galaxies. In a sample of 28 nearby galaxies, \\citet{mau99} finds $I_{\\rm CO(3-2)}$/$I_{\\rm CO(1-0)}$ ratios between 0.2 and 0.7, very different from the ones found in the high-$z$ quasars and ULIRGs. Given these uncertainties, we shall assume $L^{\\prime}_{\\rm CO(3-2)}$/$L^{\\prime}_{\\rm CO(1-0)}$=1.  For a standard $\\Lambda$CDM cosmology with H$_0$=71\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$=0.27 and $\\Omega_{\\Lambda}$=0.73, the luminosity distance at $z$=3 is $D_L$=26\\,Gpc. Assuming an intrinsic width $\\Delta v$=300\\,km$^{-1}$, the expected flux density of the CO(3--2) line is thus 36\\,$\\mu$Jy. Requesting a 5$\\sigma$ detection in a 75\\,km\\,s$^{-1}$ channel in 12 hours of on-source integration time (corresponding to 16 hours of total telescope time), we thus require an $ND^2$ of $\\sim$7300\\,m$^2$. This can be achieved with the ALMA array of 64 12\\,m antennas. Of course, larger values of $N D^2$ are always desirable, as they would allow one to resolve the line flux density into more pixels (higher angular or spectral resolution) or image to higher S/N more quickly.  Next to CO, ALMA should also be able to observe other lines such as C{\\small I}, N{\\small II} and C{\\small II} lines at cosmological redshifts. These lines will provide important probes of the IMF and the Lyman continuum luminosity from the most luminous stars in early galaxies. However, because the evolution of their luminosity as a function of redshift and Hubble type is less known than for CO, we did not use them to determine the total aperture requirement of the array.  \\subsection{Precise high-resolution imaging} The requirement for ALMA to obtain precise images at an angular resolution of 0\\farcs1 follows from the need to complement contemporary facilities such as the James Webb Space Telescope, the extended VLA, and adoptive optics imaging on large ground based telescopes.  To obtain high fidelity images with an interferometer requires a sufficient number of baselines to adequately cover the $uv$ plane. To reach such excellent images limited by dynamic range requires that 50\\% of the ($u,v$) cells be filled \\citep{mor96}. This fraction of occupied cells (FOCC) is calculated out to the longest array baselines. The FOCC is a function of hour angle coverage. Obviously, one would like to observe sources within a limited hour angle range to avoid large system temperature variations that would corrupt the images. Especially in the submillimetre windows, such variation limit the hour angle coverage to approximately 2 hours. \\citet{hol98} presents a detailed analysis of the variation of FOCC as a function of array configuration and hour angle coverage. To achieve an FOCC$>$0.5 in a configuration with a maximum baseline of 3000\\,m in 2 hours of hour angle coverage requires a collecting length $ND>$560. This can be achieved with the ALMA array of 64 12\\,m antennas.  As for any array, ALMA will also be prone to the short spacing problem. Because one cannot measure the smallest spatial frequencies, below approximately the antenna diameter, the interferometer will not be sensitive to sources more extended than $\\sim$2/3 of the primary beam. A key requirement of ALMA is therefore the ability to observe in total power mode. To reach a similar S/N level in total power would require 4 antennas optimized for total power measurements (using a nutating secondary), each observing 4 times longer than the array (hence only 25\\% of the projects will have total power information).  Even after the combination with the total power data, there will still be a gap in the $uv$ plane, located in a ring of approximately half the antenna diameter. This gap will be filled in by the Atacama Compact Array, a set of twelve 7\\,m antennas. ",
        "conclusions": ""
    },
    "0410/astro-ph0410093_arXiv.txt": {
        "abstract": "We present a non-LTE model atmosphere analysis of \\chandra\\  HRC-S/LETG and XMM-\\newton\\  RGS spectroscopy of the prototypical supersoft source CAL~83 in the Large Magellanic Cloud. Taken with a 16-month interval, the \\chandra\\  and XMM-\\newton\\  spectra are very similar. They reveal a very rich absorption line spectrum from the hot white dwarf photosphere, but no spectral signatures of a wind. We also report a third X-ray off-state during a later \\chandra\\  observation, demonstrating the recurrent nature of CAL~83. Moreover, we found evidence of short-timescale variability in the soft X-ray spectrum. We completed the analysis of the LETG and RGS spectra of CAL~83 with new NLTE line-blanketed model atmospheres that explicitly include 74 ions of the 11 most abundant species. We successfully matched the \\chandra\\ and XMM-\\newton\\  spectra assuming a model composition with LMC metallicity. We derived the basic stellar parameters of the hot white dwarf, but the current state of atomic data in the soft X-ray domain precludes a detailed chemical analysis. We have obtained the first direct spectroscopic evidence that the white dwarf is massive ($M_{\\rm WD} \\ga 1 M_\\odot$). The short timescale of the X-ray off-states is consistent with a high white dwarf mass. Our analysis thus provides direct support for supersoft sources as likely progenitors of SN~Ia. ",
        "introduction": "\\label{Intro}  Because of the overall importance of Supernovae type Ia (SNe Ia) in astrophysics and in cosmology, the identification of their progenitors is a pressing issue. Theoretical models suggest that SNe~Ia arise from the thermonuclear explosion of a carbon-oxygen (CO) white dwarf (WD) that has grown to the Chandrasekhar mass, \\Mch\\  \\citep{hoyle60, arnett69}. To date, the most promising formation channel resulting in SNe~Ia involves accreting WD's that sustain steady nuclear burning close to their surface.  When they form, CO WD's have masses between 0.7 and 1.2\\,$M_\\sun$ depending on the star initial mass \\citep{weidemann87}. SN~Ia progenitors, therefore, must be in close binaries where the WD can gain several tenths of a solar mass donated by a companion to reach \\Mch. Because we lack a direct determination of the progenitor properties, we have to rely on indirect arguments to determine their nature.  Two main scenarios have been proposed, involving either the merger of two WD's \\citep[the double-degenerate (DD) scenario;][]{iben84, webbink84}, or a single WD accreting from a normal companion \\citep[the single-degenerate (SD) scenario;][]{whelan73}. The DD scenario provides a natural way to explain the absence of hydrogen lines in SN~Ia spectra, but hardly explains why SNe~Ia could be ``standard candles'' since the explosion should depend on the respective properties of the two WD's. Moreover, while the DD scenario provides an obvious way to form a more massive object, possibly reaching \\Mch, the few double WD systems identified to date that could merge in a Hubble time have a total mass smaller than \\Mch\\ \\citep{saffer98}. New large-scale searches for DD systems, however, have recently yielded promising results \\citep{napiwotzki02, napiwotzki03}, and good candidates might be discovered soon. Finally, the merger scenario has been suggested to lead to an accretion-induced collapse rather than to a SN~Ia event \\citep{nomoto85, saio85}.  The SD scenario is thus generally favored today. The principal issue of the SD scenario is, however, to ascertain if the WD can at all accrete sufficient mass to lead to a SN~Ia event. This topic has been hotly debated in the last 10 years to sort out the competing effects of mass ejection and accretion. \\cite{nomoto79} first pointed out that an accreting WD would quickly ignite a small amount of accreted material in a hydrogen shell-burning and undergo a weak shell flash. Most of the accreted mass then accumulates in an extended, supergiant-like envelope that might be lost during a subsequent common-envelope phase. \\cite{hachisu96, hachisu99} argued that starting at a certain critical accretion rate a strong wind develops, regulating the mass transfer between the mass-losing star and the accreting WD. The WD accretes at most at the critical rate, $\\dot{M}_{\\rm cr} = 9.0\\,\\times\\,10^{-7}\\,(M_{\\rm WD}/M_\\odot - 0.50)\\,M_\\odot$\\,yr$^{-1}$, and the rest is blown off in the wind. At a later stage the accretion rate drops, and the strong wind stops; the final outcome depends on the mass accreted before this point, because stronger flashes will develop at lower accretion rates. \\cite{cassisi98} made the opposite argument. The basic physical processes are the same, but they argued that dynamical helium-burning flashes probably would hamper accretion and, therefore, the WD most likely would not reach \\Mch. The most recent theoretical calculations support the idea that the accreting WD may eventually reach \\Mch. \\cite{yoon03} have built the first detailed binary star evolution model where the WD gains enough mass from a $1.6\\,M_\\odot$ helium star. This model involves higher accretion rates (few $10^{-6}\\,M_\\odot$\\,yr$^{-1}$). Subsequently, \\citet{yoon04} have shown that rotation tends to stabilize the helium-burning shell, thus increasing the likelihood that accreting WD would reach the stage of central carbon ignition.  Owing to the difficulties met by theoretical works, an observational approach is essential to characterize SN~Ia progenitors. The most promising candidates are Close Binary Supersoft X-ray Sources (CBSS) that were revealed by ROSAT as a new class of close binaries with ultrasoft X-ray spectra showing no emission above $0.5 - 1$\\,keV \\citep{trumper91}. \\citet[][vdH92]{vdH92} argued that such a soft emission is not consistent with accreting neutron stars or black holes, which show emission peaking at $1 - 10$\\,keV. They proposed the now classical model of an accreting WD, where a relatively high-mass WD $(0.7 - 1.2\\,M_\\odot)$ sustains steady burning of the hydrogen-rich accreted material. They argued that accretion has to occur at a finely-tuned rate, $1.0\\,\\times\\,10^{-7}\\la\\dot{M}\\la 4.0\\,\\times\\,10^{-7}\\,M_\\odot$\\,yr$^{-1}$. At lower rates, hydrogen burning is unstable and occurs in flashes, while an extended envelope forms at higher rates. The stellar luminosity is then dominated by hydrogen burning which liberates an order of magnitude or more of energy than accretion itself. Typical temperatures, $kT\\approx 30 - 80$\\,eV, are derived from blackbody simulations and WD model atmospheres after correction for interstellar extinction \\citep{kahabka97}. Based on the vdH92 model, \\cite{rappaport94} discussed the formation and evolution of CBSS, reproducing their typical luminosities, effective temperatures and orbital periods. They estimated that the rate of Galactic SNe~Ia associated with the evolution of CBSS might reach 0.006~yr$^{-1}$. Most recently, \\cite{ivanova04} extended the vdH92 work, developing a semianalytical model to investigate the evolution of binaries consisting of main-sequence stars with WD companions in the thermal mass-transfer phase. They accounted for the stabilizing effect of the WD wind, and characterized the different conditions leading to different outcomes, double WD's, (sub-)\\Mch\\  SNe~Ia, or accretion-induced collapse. To evolve towards a SN~Ia explosion, \\cite{ivanova04} argued that the WD and the donor star should be initially relatively massive, $M_{\\rm WD}\\ga 0.8\\,M_\\odot$ and $M_{\\rm d}\\ga 2\\,M_\\odot$, with a mass ratio $M_{\\rm d}/M_{\\rm WD}$ smaller than about 3.  The SD scenario recently gained further observational support. \\cite{hamuy03} reported the detection of a narrow H$\\alpha$ emission line in SN2002ic, providing the first evidence of hydrogen-rich circumstellar material associated with a SN~Ia. This detection might be the missing smoking gun for the SD scenario, though \\cite{livio03} argued that this conclusion might be premature based on the absence of a detection in all other SN Ia spectra.  The exact nature of SN~Ia progenitors thus remains an open problem.  Determining the properties of CBSS may, therefore, provide the strongest empirical case supporting the SD scenario. \\chandra\\  and XMM-\\newton 's capabilities in obtaining high-resolution spectra in the soft X-ray domain now open the possibility to determine spectroscopically the stellar parameters of the WD's in CBSS, in particular the surface gravity and hence the WD mass. We present in this paper a non-LTE (NLTE) analysis of the \\chandra\\  and XMM-\\newton\\  spectra of the prototypical CBSS CAL~83. We start by summarizing previous results on CAL~83 in \\S2, and the new observations are discussed in \\S3 and \\S4. The NLTE model atmospheres are detailed in \\S5. Sect.~6 presents the spectrum analysis, leading to a conclusion that the WD in CAL~83 is massive (\\S7). ",
        "conclusions": "\\label{Conclu}  Table~\\ref{ResuTbl} summarizes our results. Our spectroscopic analysis indicates that CAL~83 contains a massive WD, $M_{\\rm WD} = 1.3\\pm 0.3 M_\\odot$. Obviously, the upper limit is well over \\Mch, but a more interesting issue to tie CBSS and SN~Ia progenitors is the lower limit for the mass.  Because low-mass models do not provide a match to the observed spectrum that is as good as the fits achieved with high-mass models, we have concluded that $M_{\\rm WD} > 1.0~M_\\odot$ is a robust lower limit for the WD mass. Nevertheless, we need to emphasize here that this conclusion depends on our current description of the soft X-ray opacities, particularly between 20 and 40\\,\\AA. We believe that the OP data already provide a good description and, although progress in this regard would be very helpful, we do not expect changes drastic enough to modify our conclusion. The issue of IS extinction is potentially more serious, but we have argued that we have an accurate estimate of the hydrogen column density toward CAL~83. In addition, our model is consistent with the WD surface having a hydrogen-rich composition with LMC metallicity. Finally, our analysis does not reveal any evidence of an outflow from the WD.  We have reported a third X-ray off-state, showing that CAL~83 is a source undergoing unstable nuclear burning. This is consistent with its low luminosity. The short timescale of the off-states (about 50 days) provides a supporting evidence that CAL~83 WD is massive ($M_{\\rm WD}\\ga 1.35 M_\\odot$). A better characterization of the off-states is required to definitively establish the mechanism(s) responsible of the off-states. Moreover, we confirm that the X-ray luminosity is fairly constant on a long-term basis during on-states. However, on the other hand, we found variations up to 50\\% of the soft X-ray luminosity in very short times (less than 0.1 orbital period). Simultaneous X-ray and optical photometry covering continuously a duration longer than a full orbital period would be very valuable to study stochastic accretion processes.  Within the model of SN~Ia progenitors proposed by \\cite{hachisu96}, our results would place CAL~83 in a late stage, after the strong accretion and wind phase when the accretion rate drops below the critical rate for sustaining steady nuclear burning. Our results therefore make CAL~83 a very likely candidate for a future SN~Ia event. We plan to conduct a similar NLTE analysis of the few other CBSS already observed with \\chandra\\  LETG spectrometer to further support our results and quantitatively study differences between these objects. In particular, the second typical CBSS, CAL~87, exhibits a very different spectrum characterized by emission lines \\citep{greiner04}. A detailed spectroscopic analysis should reveal either if the different spectrum arises from a different geometry (edge-on {\\it vs.\\/} face-on system) or if CAL~87 is in an earlier evolutionary stage."
    },
    "0410/astro-ph0410570_arXiv.txt": {
        "abstract": "The identification of pre-acceleration mechanisms for cosmic ray ions in supernova remnant shocks is an important problem in astrophysics. Recent particle-in-cell (PIC) shock simulations have shown that inclusion of the full electron kinetics yields non-time-stationary solutions, in contrast to previous hybrid (kinetic ions, fluid electrons) simulations.  Here, by running a PIC code at high phase space resolution, ion acceleration mechanisms associated with the time dependence of a supercritical collisionless perpendicular shock are examined. In particular the components of $\\int \\mathbf{F} \\cdot \\mathbf{v} dt$ are analysed along trajectories for ions that reach both high and low energies. Selection mechanisms for the ions that reach high energies are also examined. In contrast to quasi-stationary shock solutions, the suprathermal protons are selected from the background population on the basis of the time at which they arrive at the shock, and thus are generated in bursts. ",
        "introduction": "Understanding the initial acceleration mechanisms for Galactic cosmic rays remains an  outstanding problem in astrophysics. From energy balance considerations, supernova remnants (SNRs) provide the most likely kinetic energy source to sustain the cosmic ray population. The local acceleration of electrons has been indirectly observed at the expanding shock front of SNRs (see, for example, Ref.~\\cite{koyama1995}). However protons form the majority constituent of Galactic cosmic rays, and  until recently observational evidence to link SNRs to local ion acceleration has been lacking. X-ray and $\\gamma$-ray data from supernova remnant RX J1713.7-3946 \\cite{enomoto2002} show energy spectra that can only be explained by accelerated ions.  Several mechanisms are postulated to accelerate particles at SNR shocks.  Fermi acceleration \\citep{fermi1949}, which arises as a particle repeatedly scatters off turbulent structures on either side of the shock, is in principle capable of accelerating ions to high energies \\cite{bell1978}. However, to work effectively an initial suprathermal population is required so that particles can re-cross the shock front \\cite{jokipii1987}. The identification and analysis of pre-acceleration mechanisms that can select and initiate the energisation of completely non-relativistic ions at SNR shocks from the background plasma is the subject of this paper.  The Rankine-Hugoniot relations \\citep{tidmanandkrall} can be used to determine the discontinuity in bulk plasma parameters across a collisionless shock; that is, a shock where the particle mean free path is much greater than length scales of interest. These relations are derived by applying the magnetohydrodynamic (MHD) conservation equations in the far upstream and downstream limits, away from the shock. Further conditions are imposed by the fact that a shock must also increase entropy, so that no subsonic flow can spontaneously become supersonic. For an Alfv\\'{e}nic Mach number $M_A \\gtrsim 3$, the shock is supercritical in that the increase in entropy, and in ion heating, required by the Rankine-Hugoniot relations is achieved via the ion kinetics; at least in part, by reflection of a fraction of upstream ions at the shock. The generic supercritical, quasi-perpendicular, collisionless shock in which ions reflect and gyrate in a foot region upstream has been suggested by hybrid (particle ions, fluid electrons) simulations (see, for example, Refs.~\\cite{quest1985,quest1986,burgess1989}), and confirmed by in-situ observations of the Earth's bow shock \\cite{sckopke1983}.  To study the acceleration of ions and electrons, a fully kinetic treatment is in principle necessary for both species, and this can be closely approximated by particle-in-cell (PIC) techniques. Physical mechanisms operating on electron kinetic lengthscales and timescales are significant both for aspects of macroscopic structure (for example, the shock ramp width scales as $c/\\omega_{pe}$), and for microscopic processes affecting the ions (such as caviton formation and dissolution). Whether such effects are important in any given scenario can be estimated, to some extent, by analytical means, as we discuss in detail below. Importantly, however, it is known that inclusion of the full electron kinetics can significantly alter the dynamics of the shock. For example, hybrid simulations for certain parameters \\cite{quest1986,burgess1989} produce time-stationary shock solutions, whereas for the same parameters PIC simulations reveal a dynamic, reforming, shock structure. Furthermore the extent to which an individual ion responds to phenomena on electron kinetic scales must depend on that ion's cyclotron radius, and hence its energy. It follows that for studies of ion acceleration at shocks, as in the present paper, retention of full electron kinetics is desirable in order to resolve fully the shock dynamics (see also Refs.~\\cite{scholer2003,lee2004}) and the ion dynamics. We have previously presented  results of PIC code simulations \\cite{schmitz2002a,schmitz2002b,lee2004} that have high resolution in real space and phase space, over relatively long run times, for parameters relevant to shocks at supernova remnants. Our most recent results \\cite{lee2004} show that the time-dependent electromagnetic fields at the reforming shock can accelerate inflow ions from background to suprathermal energies. This provides a source population which may subsequently be accelerated to produce high energy cosmic rays.   In the present paper, we focus on the specific nature of this ion acceleration mechanism. This requires careful examination of the physics of the interaction between particles and fields as they evolve over time. We first introduce a methodology for simplifying the raw data, obtained on spatio-temporal scales spanning those of the electrons and ions, into data  suitable for examining events on the spatio-temporal scales of interest for ion acceleration. We then examine the detailed dynamics of the ion interactions with the shock front, including a comparison between ions that eventually reach the highest and lowest energies downstream. The time at which particles are incident on the temporally evolving shock structure is found to be a key discriminant in the subsequent energisation. ",
        "conclusions": "We have examined in detail the dynamics of suprathermal ions generated in PIC code simulations of quasi-perpendicular reforming shocks. Importantly, this energisation is not found in stationary shock solutions. We find that: \\begin{enumerate} \\item The shock structure reforms on a timescales of the order of the local ion cyclotron period. This is shown clearly if the electromagnetic fields are cast in the form of a potential, after removing small scale effects, to leave only terms relevant on ion spatio-temporal scales. \\item The time-dependence of the shock dominates the selection of which ions are accelerated to suprathermal energies. Ions that reach the shock when its ramp, and hence potential, are maximal, are reflected and subsequently gain energy by drifting in the time-dependent fields tangential to the shock front. \\item  This selection is in contrast to a time-stationary shock, where the selection mechanism depends upon the initial ion velocity perpendicular to the magnetic field, those  ions coming from the tails of the distribution being preferentially reflected, and so energised. \\item These factors lead to high energy ion creation occurring in bursts. \\end{enumerate}  The present simulations are conducted in a $1x3v$ geometry at a low $m_i/m_e$. Whilst the work in Ref. \\cite{lembege1992} shows no major differences in  higher spatial dimensions for PIC simulations, hybrid simulations in $2x3v$, for longer run times, show that the current that can now exist along the shock front can lead to current-driven instabilities \\cite{winske1988}. These instabilities may act to change the shock structure in the $y$-direction, and so alter ion and electron dynamics across the shock front, affecting both ion and electron acceleration. Simulations with more realistic $m_i / m_e$ ratios \\cite{lembege2002,scholer2003} show alterations in the electron scale physics in the foot region. However, we have shown that the electron scale physics has little effect on ion spatio-temporal scales.  The plasmas simulated here are pure hydrogen, in that there are only two species, protons and electrons. Both species have a Maxwellian distribution with the same temperature. The addition of pickup ions to the simulation, for example in relation to the heliospheric termination shock, where hybrid simulation have already been carried out \\cite{lipatov1999,lever2001}, would allow the acceleration processes relevant to anomalous cosmic ray production to be examined. This will be the subject of future work.  The fundamental plasma physics processes underlying the ion acceleration from background to suprathermal energies (10 to 20 MeV), reported in the SNR shock simulations of Ref.~\\cite{lee2004}, have been elucidated in the present paper. Specifically, we have explained the role of, and interplay between, the key elements anticipated at the end of the Appendix to Ref.~\\cite{lee2004}. We have shown that, while an electron fluid approximation captures some of the key physics, the shock reformation dynamics arising from our fully kinetic PIC treatment are central to the ion acceleration mechanism. This work provides a clear first principles explanation for the ion acceleration that is observed in our simulations, which appears to be a strong candidate injection mechanism for Galactic cosmic ray protons."
    },
    "0410/hep-ph0410103_arXiv.txt": {
        "abstract": "I summarize the present status of global analyses of neutrino oscillations, including the most recent KamLAND and K2K data, as well as the latest solar and atmospheric neutrino fluxes.  I give the allowed ranges of the three--flavour oscillation parameters from the current worlds' global neutrino data sample, their best fit values and discuss the small parameters $\\alpha \\equiv \\Dms/\\Dma$ and $\\sin^2\\theta_{13}$, which characterize the strength of CP violation in neutrino oscillations. I briefly discuss neutrinoless double beta decay and the LSND neutrino oscillation hint, as well as the robustness of the neutrino oscillation results in the presence of non-standard physics.  \\vspace{1pc} ",
        "introduction": "The discovery of neutrino oscillations by combining data from solar~\\cite{solarNF}, atmospheric~\\cite{atmNF}, reactor~\\cite{reacNu04} and accelerator~\\cite{accelNu04} neutrino studies has marked a turning point in our understanding of nature and has brought neutrino physics to the center of attention of the particle, nuclear and astrophysics communities. This culminates a heroic effort which started over four decades ago and now firmly establishes the existence of small neutrino masses, confirming theoretical expectations which date back to the early eighties. Neutrino mass arise from the dimension-five operator $\\ell \\ell \\phi \\phi$ where $\\phi$ the \\21 Higgs doublet and $\\ell$ is a lepton doublet~\\cite{Weinberg:1980bf}.  Nothing is known from first principles about the mechanism that induces this operator, its associated mass scale or flavour structure.  Its most popular realization is the seesaw mechanism which induces small neutrino masses from the exchange of heavy states. Although inspired by unification, the effective model-independent description of the seesaw at low-energies is in terms of the \\21 gauge structure and contains a small effective Higgs triplet contribution to the neutrino mass~\\cite{schechter:1980gr}. Such general seesaw has 24 parameters, 3 more than the tripletless case~\\cite{schechter:1980gr}~\\footnote{In the mass basis these correspond to the 12 mixing angles and 12 CP phases (both Dirac and Majorana-type) that characterize the full 3$\\times$6 charged current seesaw lepton mixing matrix~\\cite{schechter:1980gr}.  These are exactly the same parameters involved in the description of leptogenesis~\\cite{Fukugita:1986hr}, though in this case the use of the weak basis seems more convenient. Note also that current nomenclature (type-I versus II seesaw) is opposite from the original one in \\cite{schechter:1980gr}.}.  The structure of the three-flavour lepton mixing matrix in various gauge theories of neutrino mass was given in~\\cite{schechter:1980gr} and it was also argued that, on \\emph{general} grounds, massive neutrinos should be Majorana particles, leading to neutrinoless double beta decay (\\nbb)~\\cite{Wolfenstein:1981rk}.  Current neutrino oscillation data are well described by the simplest unitary lepton mixing matrix neglecting CP violation. The effect of Dirac CP phases in oscillations and Majorana phases in \\nbb constitute the main challenge for the years to come.     Here I focus mainly on the determination of neutrino mass and mixing parameters in neutrino oscillation studies, a currently thriving industry~\\cite{industry}, with many new experiments currently underway or planned. The interpretation of the data requires good calculations of solar and atmospheric neutrino fluxes~\\cite{Bahcall:2004fg,Honda:2004yz}, neutrino cross sections and experimental response functions, as well as a careful description of neutrino propagation both in the Sun and the Earth, including matter effects~\\cite{mikheev:1985gs,wolfenstein:1978ue}.  After summarizing the latest (post-Neutrino-2004) global analysis of 3-neutrino oscillation parameters~\\cite{Maltoni:2004ei}~\\cite{Goswami2004} I briefly discuss their impact on future \\nbb searches~\\cite{Bilenky:2004wn}. This writeup combines my plenary talk and the WG1 parallel session talk on non-standard scenarios of neutrino conversion. ",
        "conclusions": "Experiment is ahead of theory in neutrino physics: despite the great progress achieved recently we are very far from a ``road map'' to the ultimate theory of neutrino properties.  We have no idea of the underlying neutrino mass generation mechanism, its characteristic scale or its flavor structure. We have still a long way to go!"
    },
    "0410/hep-ph0410117_arXiv.txt": {
        "abstract": "#1{% \\gdef\\@bstract{% \\parbox{\\textwidth-0pc}{% \\centerline{\\bf Abstract}\\penalty1000% \\kern.2cm% \\noindent% \\renewcommand\\baselinestretch{1.0}% {#1}}} }   \\def\\ps@paper{\\let\\@mkboth\\@gobbletwo% \\ifnum\\draftcontrol=1 \\def\\@oddfoot{\\hbox to \\textwidth{\\tiny \\versionno \\hfil\\tiny\\draftdate}% \\hskip -\\textwidth \\hbox to \\textwidth{\\hfil\\rm\\thepage\\hfil}}% \\else\\def\\@oddfoot{\\hbox to \\textwidth{\\hfil\\rm\\thepage\\hfil}} \\fi \\let\\@evenfoot\\@oddfoot }  \\def\\body{\\clearpage \\pagestyle{paper} } \\newenvironment{acknowledgments}{% \\vskip 3.25ex \\noindent {\\bf Acknowledgments} }   \\def\\@version#1{\\ifnum\\draftcontrol=1 \\typeout{}\\typeout{#1}\\typeout{} \\vskip3mm\\centerline{\\hbox{\\fbox{\\normalsize{\\tt DRAFT -- #1 -- } {\\draftdate}}}}\\vskip3mm \\fi} \\let\\version\\@version \\long\\def\\eqlabel#1{\\ifnum\\draftcontrol=1 \\tag@false  % \\tag*{(\\theequation) \\hbox to -0.2cm{\\hspace{0cm}\\small{#1}\\hss}} \\refstepcounter{equation} \\edef\\@currentlabel{\\theequation} \\ltx@label{#1}          % \\else \\label{#1} \\fi } \\let\\st@bibitem\\@bibitem \\let\\st@lbibitem\\@lbibitem \\ifnum\\draftcontrol=1 \\def\\@bibitem#1{% \\st@bibitem{#1}\\a@@label{#1}\\ignorespaces} \\def\\@lbibitem[#1]#2{% \\st@lbibitem[#1]{#2}\\a@@label{#2}\\ignorespaces} \\def\\a@@label#1{% \\gdef\\a@lab{\\smash{\\normalfont\\small#1}} \\ifvmode \\if@inlabel \\global\\setbox\\@labels\\hbox{% \\llap{\\a@lab\\let\\a@lab\\relax \\kern\\@totalleftmargin\\kern\\marginparsep}% \\box\\@labels}% \\fi \\fi} \\fi   \\documentclass[12pt,letterpaper]{article}  \\usepackage{amsmath,amssymb,array,calc,rotating,epsfig,psfrag} \\usepackage[nosort]{cite}  \\ifnum\\draftcontrol=1 \\tolerance=1000 \\fi  \\renewcommand\\baselinestretch{1.25} \\setlength{\\paperheight}{11in} \\setlength{\\paperwidth}{8.5in} \\setlength{\\textwidth}{\\paperwidth-2.4in}     \\hoffset= -.3in   % \\setlength{\\textheight}{\\paperheight-2.4in}   \\topmargin= -.6in %  \\renewcommand{ The Wilkinson Microwave Anisotropy Probe microwave background data suggest that the primordial spectrum of scalar curvature fluctuations is suppressed at small wavenumbers. We propose a UV/IR mixing effect in small-field inflationary models that can explain the observable deviation in WMAP data from the concordance model.  Specifically, in inflationary models where the inflaton couples to an asymptotically free gauge theory, the radiative corrections to the effective inflaton potential can be anomalously large. This occurs for small values of the inflaton field which are of the order of the gauge theory strong coupling scale. Radiative corrections cause the inflaton potential to blow up at small values of the inflaton field.  As a result, these corrections can violate the slow-roll condition at the initial stage of the inflation and suppress the production of scalar density perturbations. }       \\makepapertitle  \\body  \\version\\versionno ",
        "introduction": "High precision observational cosmology places tight constraints on the cosmological parameters of our universe. The Wilkinson Microwave Anisotropy Probe (WMAP) results measuring anisotropies in the Cosmic Background Radiation (CBR) \\cite{cmb1,cmb2}, when combined with data on high redshift supernovae and large scale structure, strongly support the ``concordance model'' for our universe\\footnote{The concordance model is a spatially flat Universe with an adiabatic, nearly scale invariant spectrum of initial fluctuations. In what follows we assume a $\\Lambda$CDM model as a realization of the concordance model.}.  One of the intriguing observations of the CBR anisotropy spectrum is the suppression of its low-$\\ell$ multipoles compared to the predictions of $\\Lambda$CDM model. The low-$\\ell$ multipoles of the temperature-temperature (TT) angular power spectrum correspond to large angular scales --- they encode the information about the small wavenumbers in the spectrum of primordial density perturbations. Consequently they provide a window on the detailed features of the inflaton potential during the part of inflation \\cite{guth} that gives rise to observables in structure formation and the microwave background, produced roughly 60-50 e-foldings before the end of inflation. Unfortunately, it is not possible to disentangle the suppression of low-$\\ell$ modes from cosmic variance limitations. In this paper we assume the suppression is a physical effect.  Not surprisingly, this suppression of low-$l$ multipoles (along with the possibility of a running spectral index) was the focus of much recent discussion\\cite{e1,e2,e3,e4,e5,e6,e7,e8,e9,k,e10,e11,e12,e13,e14,e15,e16}. Most previous work attempted to identify a new ultraviolet (UV) physics responsible for the suppression of the primordial power spectrum at small wavenumbers. In this paper we point out a simple low-energy (albeit strongly coupled) field-theoretic phenomenon that produces a 'feature' in the inflaton potential that serves to explain the suppression of low-$\\ell$ CBR anisotropies\\footnote{Effects of strongly coupled gauge theory dynamics on inflation were also studied in \\cite{w}.}. The observed effect is specific to small-field inflationary models \\cite{r0,r}.  Here, small-field models are defined to be those models of inflation in which the initial value of the scalar field is small (e.g. near zero) as it starts rolling down the potential.   We find a phenomenon that modifies the usual power spectrum of density fluctuations from inflation, \\begin{equation} \\left(\\delta^{(0)}_H\\right)^2\\propto \\left(\\frac{k}{k_s}\\right)^{n_s-1}\\,, \\end{equation} to \\begin{equation} \\delta_H^2=\\left(\\delta_H^{(0)}\\right)^2\\ f\\left(\\frac{k}{k_s}\\right)\\,, \\end{equation} where the superscript $0$ refers to the unmodified spectrum, and the form-factor $f(\\phi)$ has the following behavior: \\begin{equation} f(k) \\rightarrow 0 \\,\\,\\,\\, {\\rm (small} \\,\\, k) \\end{equation} \\begin{equation} f(k) \\rightarrow 1 \\,\\,\\,\\, {\\rm (large} \\,\\, k) . \\end{equation} Here $k_s$ corresponds to the scale of perturbations that are produced $60$ e-foldings before the end of inflation, which corresponds to the present horizon size, $k_s\\sim (4000 Mpc)^{-1}$.  The suppression of large-scale power arises because the potential blows up (becomes infinite) at small values of the field, so that the field is not slowly rolling at all and production of density fluctuations is suppressed; see Fig. (1).  This blow-up of the potential is sharply localized at 60 e-foldings before the end of inflation, so that ordinary density fluctuations and structure formation ensue just afterwards (at 60-50 e-foldings before the end). While the sharpness of the feature is generic to our mechanism, its location at exactly 60 e-folds before the end requires fine tuning\\footnote{ The same fine-tuning problem is acknowledged in all previous work on the small-$\\ell$ suppression of CBR power spectrum we are aware of \\cite{e1,e2,e3,e4,e5,e6,e7,e8,e9,k,e10,e11,e12,e13,e14,e15,e16}. Unfortunately, neither of the proposed models provided the solution to it. It would be very interesting to identify physical phenomena that alleviates the latter fine tuning problem.}. The origin of the blow-up of the potential is due to the radiative corrections for an inflaton coupled to an asymptotically free gauge theory (analogous to QCD).    \\begin{figure}[ht] \\begin{center} \\epsfig{file=figure0.eps,width=0.9\\textwidth} \\caption{ Effective inflaton potential (solid line) in small-field inflationary models with inflaton coupled to asymptotically free gauge theory. The slow-roll tree-level potential (dashed line) receives large radiative corrections whenever the inflaton expectation value is close to the gauge theory strong coupling scale $\\phi\\sim \\Lambda\\ll v$. In the sharply localized regime $|\\phi-\\Lambda|\\ll v$ the field is not slowly rolling, and production of density fluctuations is strongly suppressed.  } \\label{cases0} \\end{center} \\end{figure}   Thirty years ago Coleman and Weinberg  computed the one-loop effective potential for (classically) conformally invariant $\\SU(2) \\times \\U(1)$ gauge theory with no quadratic mass term.  However, due to the massive top quark, the contributions due to the $t$-quark's Yukawa coupling constant have subsequently been found to dominate over those of the $\\SU(2)\\times \\U(1)$ gauge coupling constants. Moreover, when there is a large Yukawa coupling, subsequent leading logarithm terms to the one-loop effective potential are too large to neglect. In the next section we review a technique developed in Ref.  \\cite{a42,a4} for all-order summation of leading logarithm terms for the effective potential, including coupling to asymptotically free theory (QCD). Radiative symmetry breaking as an explanation for the Higgs mass is revived in this context, since a prediction for the Higgs mass above 200 GeV can now be accommodated \\cite{a42}.  The physics characterizing this model extends to more general  scalar potentials, including that of any inflaton coupled to an asymptotically free gauge theory. We explicitly compute the leading contribution to the effective inflaton (Higgs) potential and show that radiative corrections become anomalously large for small vacuum expectation values of the inflaton. This is a reflection of a new UV/IR mixing effect in the model: even though the inflationary potential has a typical GUT scale height, the effective QCD coupling in radiative corrections to the inflaton potential is evaluated at the field value of the inflaton. The effect is thus most profound if the inflaton value at the beginning of inflation is close to the gauge theory's strong coupling scale.  Thus, for a sufficiently small initial inflaton value, the QCD coupling develops a perturbative Landau pole that strongly enhances radiative corrections.  In section 3 we model these radiative corrections in small-field inflationary scenarios and demonstrate that they suppress the primordial power spectrum of scalar density fluctuations for small wavenumbers.  We use the CMBEASY package \\cite{cmbeasy} to relate the latter suppression to the small multipole suppression in the CBR spectrum of anisotropies as observed by WMAP.  We summarize  our results in section 4. ",
        "conclusions": "WMAP data indicates the suppression of small-$\\ell$ multipoles in the anisotropy power spectrum of the CBR. It is difficult to isolate this suppression from the cosmic variance limitations. However, when taken seriously, such suppression can be due to a suppression of power in the spectrum of primordial scalar density fluctuations during inflation.  As the generation of quantum fluctuations during inflation is most efficient in the slow-roll regime, the CBR anisotropy measurements provide a signature for the breakdown of the slow-roll conditions during the first few e-foldings of the single field inflationary models. In this paper we have identified a low-energy field-theoretic phenomenon that arises generically in small-field inflationary models in which the inflaton is coupled to an asymptotically free GUT-like gauge theory. Here, the small vacuum expectation values of the inflaton field probe the strongly coupled dynamics of this gauge theory. This produces large radiative corrections to the tree-level inflaton potential which break its slow-roll conditions.  The effect is most profound if the inflaton scale at the beginning of inflation is close to the gauge theory's strong coupling scale.  In this paper we have used summation of leading logarithms in the perturbation theory as a model of a strongly coupled dynamics of an inflaton and asymptotically free gauge theory.  It would be very interesting to explore nonperturbative effects of the quantum field theory in this setting, particularly space-time curvature corrections."
    },
    "0410/hep-th0410223_arXiv.txt": {
        "abstract": "It has recently been suggested that Planck scale physics may effect the evolution of cosmological fluctuations in the early stages of cosmological inflation in a non-trivial way, leading to an excited state for modes whose wavelength is super-Planck but sub-Hubble. In this case, the issue of how this excited state back-reacts on the background space-time arises. In fact, it has been suggested that such back-reaction effects may lead to tight constraints on the magnitude of possible deviations from the usual predictions of inflation. In this note we discuss some subtle aspects of this back-reaction issue and point out that rather than preventing inflation, the back-reaction of ultraviolet fluctuations may simply lead to a renormalization of the cosmological constant driving inflation. ",
        "introduction": "The most important success of the inflationary Universe scenario \\cite{Guth} is that it provides a causal mechanism for the origin of the observed density fluctuations and microwave background anisotropies~\\cite{Mukhanov} (see also Refs.~\\cite{Press,Lukash}). Key to this success is the fact that the physical wavelength corresponding to a fixed comoving scale is exponentially stretched during the period of inflation. Thus, provided that the period of inflation lasts sufficiently long, fluctuations on scales of cosmological interest today originate on sub-Hubble scales during inflation. Since inflation red-shifts all initial classical fluctuations, it is reasonable to assume that matter starts out in a quantum vacuum state (in the frame set by the background cosmology, see e.g. Ref.~\\cite{RHB84} for a discussion). Each fluctuation mode thus starts out in its vacuum state at the time that the initial conditions are set up (e.g. the beginning of the period of inflation), it undergoes quantum vacuum oscillations while the wavelength is smaller than the Hubble radius, but freezes out when the wavelength equals the Hubble radius (see e.g. Ref.~\\cite{MFB} for comprehensive reviews of the theory of cosmological fluctuations). Subsequently, the quantum state of the fluctuations undergoes squeezing on super-Hubble scales, and re-enters the Hubble radius during the post-inflationary Friedman-Lema\\^\\i tre-Robertson-Walker (FLRW) phase as a highly squeezed and effectively classical state (see e.g. Ref.~\\cite{Polarski} for a discussion of the classicalization of the state).  \\par  However, as first pointed out in Ref.~\\cite{RHBrev}, this success of inflationary cosmology leads to an important conceptual problem, the {\\it trans-Planckian problem}. Since the period of inflation in typical scalar-field-driven inflationary models is very long (see e.g. Ref.~\\cite{Lindebook} for a review), the scales of cosmological interest today are not only sub-Hubble, but in fact sub-Planck at the beginning of inflation. Thus, the formalism used to calculate the evolution of fluctuations is in fact not justified. It is possible that the unknown trans-Planckian physics will lead to an evolution of the fluctuations on sub-Planckian scales which from the point of view of free scalar field theory coupled to General Relativity looks non-adiabatic.  In Ref.~\\cite{MB1} (see also Ref.~\\cite{Niemeyer}), toy models for such an evolution were constructed making use of modified dispersion relations which were assumed to describe the physics on sub-Planckian scales. Since the time interval spent in the trans-Planckian domain may depend on the wavelength, such models may lead to changes in the spectral index of the fluctuations.  \\par  Subsequently, other approaches to the trans-Planckian problem were suggested, e.g. analysis based on space-space non-commutativity~\\cite{Columbia}, space-time non-commutativity~\\cite{Ho}, minimal trans-Planckian assumptions (starting each mode in some vacuum state at the time when its wavelength equals the Planck length)~\\cite{Ulf,Ven}. These analysis typically give that trans-Planckian corrections to the predictions for cosmological fluctuations are proportional to $(H_{\\rm inf}/\\mC)^n$ where $H_{\\rm inf}$ is the Hubble parameter during inflation, $\\mC$ a new scale at which non standard physical effects show up and $n$ a number which depends on the initial state assumed at the time of ``creation''. On the other hand, analysis based on modified dispersion relations give a correction proportional to the time spent by the physical modes in the region where adiabaticity is violated, see e.g.~\\cite{MB2} for a recent review.  \\par  Tanaka~\\cite{Tanaka} and Starobinsky~\\cite{Starob} (see also Ref.~\\cite{later}) have, however, raised an important concern regarding the possible amplitude of trans-Planckian corrections (see also Ref.~\\cite{TS} for unrelated concerns): if trans-Planckian physics leads to an excited state for fluctuation modes on sub-Hubble but super-Planck scale during the period of inflation, the back-reaction of these excitations on the background must be considered. In the case where the trans-Planckian effects are modeled by a modified dispersion relation, a simple estimate of the energy density carried in these ultraviolet modes \\begin{equation} \\label{estimate} \\langle \\rho \\rangle _{_{\\rm UV}} = \\int_{k_{\\rm phys} = H_{\\rm inf}} ^{k_{\\rm phys} = \\mC} {\\rm d}^3k_{\\rm phys} \\omega_{\\rm phys}(k_{\\rm phys}) n_{k_{\\rm phys}}  \\, , \\end{equation} where $k_{\\rm phys}$ is the physical wavenumber, $n_{k_{\\rm phys}}$ is the occupation number, and $\\omega_{\\rm phys}$ is the frequency of the mode, leads to the conclusion that $\\langle \\rho \\rangle _{_{\\rm UV}}$ will exceed the background density unless $n_{k_{\\rm phys}}$ is smaller than $(H_{\\rm inf}/m_{_{\\rm Pl}})^2$ [see also Eq.~(\\ref{consbeta}) above], thus constraining the possible effects of trans-Planckian physics on the spectrum of fluctuations (in the above, we have assumed that the ultraviolet cutoff is the usual Planck scale $m_{_{\\rm Pl}}$ which is not mandatory at all).  \\par  In this article, we point out some subtleties with the above back-reaction argument which may change the conclusions dramatically. If we assume, as is conventionally done in analyzing quantum fields in curved space-time, that the ultraviolet cutoff scale is time-independent in terms of physical length, then in an exponentially expanding background geometry, the contribution of ultraviolet (i.e. sub-Hubble) modes to the energy density is constant in time, as, in fact, follows directly from the time translation invariance of the physics. Moreover, the corresponding equation of state, due to the fact that the dispersion relation is modified, can strongly differ from that of ultra-relativistic particles and, as we demonstrate below, tends to that of the vacuum. Hence, our main conclusion is that, instead of preventing inflation, the ultraviolet modes may in fact simply renormalize the value of the cosmological constant driving inflation.  \\par  This article is organized as follows. In the next section, Sec.~II, we describe the arguments that have been put forward to claim that there is a back-reaction problem and we criticize them. Then, in Sec.~III, we present an explicit calculation of the equation of state of a scalar field with a modified dispersion relation. We show that the ultraviolet modes possess an equation of state which is almost that of a cosmological constant. Finally, in Sec.~IV, we point out problems with our approach, indicate directions for further investigations and present our general conclusions. ",
        "conclusions": "Now that we have determined the energy density and pressure of the ultraviolet modes, we will study their back-reaction on the background space-time and matter. We work in the context of the toy model studied in the previous section. We assume that at early times all modes with wavenumber larger than $\\Lambda_2$ start out in their adiabatic vacuum. Thus, after subtraction of the quantum vacuum terms, these modes do not contribute to the energy-momentum tensor, and thus there will be no remaining ultraviolet divergences.  \\par  We now follow a mode with fixed comoving wavelength which starts out deep in the ultraviolet region ($k_{\\rm phys}>\\Lambda_2$) in its adiabatic vacuum. The mode will then spend a finite time interval in the intermediate frequency range $\\Lambda_1<k_{\\rm phys} < \\Lambda_2$ during which the adiabaticity condition for mode evolution is violated, the state gets squeezed, and, from the point of view of the vacuum state for $k_{\\rm phys}<\\Lambda_1$, a non-vanishing occupation number $n_k$ is generated. This occupation number remains constant when $k_0<k_{\\rm phys} < \\Lambda_1$, since in this frequency range the adiabaticity condition is restored.  \\par  Given this setup, we consider the energy density $\\rho_{UV}$ of ultraviolet modes (modes with wavelength smaller than the Hubble radius) which is the sum of~(\\ref{nuven}) and~(\\ref{fuven}), see Eq.~(\\ref{totalrho}). By time-translation invariance of the background, both terms are independent of time in a de Sitter background in which $H_{\\rm inf}$ is constant. Thus, \\begin{equation} \\label{constancy} \\frac{{\\rm d}\\rho_{_{\\rm UV}}}{{\\rm d}t} \\, = \\, 0 \\, . \\end{equation} This seems to indicate that the ultraviolet energy density evolves like a cosmological constant, and its back-reaction, rather than preventing inflation, will simply lead to a renormalization of the cosmological constant.  \\par  Naive intuition, namely treating the equation of state as radiative, \\ie $p_{_{\\rm UV}}=\\rho_{_{\\rm UV}}/3$, would have led to a problem with this conclusion, namely an extreme non-conservation of the energy-momentum tensor of the ultraviolet modes, a non-conservation on the energy density scale of $\\Lambda _1^4$ (in this context the difference between $\\Lambda_1$ and $\\Lambda_2$ is not relevant). However, our analysis of the previous section has shown that in fact the energy-momentum tensor of the ultraviolet modes is that of a cosmological constant, up to correction terms which are suppressed by a factor of $H_{\\rm inf}/\\mC$ compared to the dominant terms.  \\par  Let us take a look at the equations of back-reaction. Following the method discussed in Refs.~\\cite{Abramo1,Abramo2} in the context of the problem of the back-reaction of infrared cosmological fluctuations. The back-reaction effect of interest here is the fact that linear cosmological fluctuations effect the background metric and matter if one works out the Einstein equations to quadratic order in the amplitude of the primordial perturbations. To be specific, in the following we shall consider a homogeneous, isotropic and spatially flat background, and will take matter to be a scalar field $\\varphi$ with a quadratic potential given by the scalar field mass $m$. For the fluctuations we parameterize the equation of state as \\begin{equation} \\label{eos2} p_{_{\\rm UV}}\\equiv (-1+\\alpha )\\rho_{_{\\rm UV}} \\, , \\end{equation} where (by the above discussion) $\\alpha $ is a positive constant expected to be of the order $H_{\\rm inf}/\\mC$.  \\par  Our starting point consists of taking the FLRW equations \\begin{eqnarray} \\label{FRW} H^2 &=& {\\kappa \\over 3} \\rho \\, , \\quad \\frac{{\\rm d}\\rho}{{\\rm d}t} = - 3 H (\\rho + p) \\, , \\end{eqnarray} where $\\kappa \\equiv 8\\pi /\\mP^2$. The presence of the fluctuations produces a back-reaction effect on the background which is quadratic in the amplitude of the fluctuations~\\cite{MAB,ABM} and leads to a correction of both metric and matter, \\ie to corrections $\\delta H$ of the Hubble expansion rate and $\\delta \\varphi $ of the scalar field \\begin{eqnarray} \\label{bransatz} H= H_{\\rm inf}+ \\delta H \\, ,\\quad \\varphi = \\varphi_{\\rm inf} + \\delta \\varphi  \\, , \\end{eqnarray} where the subscripts ``${\\rm inf}$'' stand for the quantities evaluated in the unperturbed background. The back-reaction of the linear fluctuations on the background is described by contributions $\\rho_{\\rm br}=\\rho _{_{\\rm UV}}$ and $p_{\\rm br}=p_{_{\\rm UV}}$ (and we use the ultraviolet energy densities and pressures as the back-reaction quantities) to the energy density and pressure. If we insert the back-reaction ansatz~(\\ref{bransatz}) into the equations~(\\ref{FRW}) (written in the slow-roll approximation) and linearize in $\\delta H$ and $\\delta \\varphi $, we obtain the following equations \\begin{widetext} \\begin{eqnarray} \\label{lin1} 2 H_{\\rm inf}\\, \\delta H &=& {\\kappa \\over 3}\\left( \\rho_{_{\\rm UV}} + m^2 \\varphi_{\\rm inf}\\delta \\varphi \\right)\\, ,\\quad m^2\\left(\\frac{{\\rm d}\\varphi_{\\rm inf}}{{\\rm d}t}\\right)\\delta \\varphi +m^2\\varphi_{\\rm inf}\\left(\\frac{{\\rm d}\\delta \\varphi }{{\\rm d}t}\\right) = -3 H_{\\rm inf}\\left( \\rho_{_{\\rm UV}} + p_{_{\\rm UV}} \\right) -\\frac{{\\rm d}\\rho_{_{\\rm UV}}}{{\\rm d}t} \\, . \\end{eqnarray} \\end{widetext} We have assumed that the background model has a massive potential, \\ie $V(\\varphi)=m^2\\varphi ^2/2$. The back-reaction of infrared modes (modes with wavelength larger than the Hubble radius) was analyzed in \\cite{MAB,ABM} (see also Ref.~\\cite{TW}) and was shown to correspond to a negative cosmological constant whose absolute value increases in time (see Ref.~\\cite{RBrev2} for resulting speculations on how this effect might be used to address the cosmological constant problem). Here, we will study the back-reaction effects of the ultraviolet modes for which ${\\rm d}\\rho_{\\rm br}/{\\rm d}t = {\\rm d}\\rho_{_{\\rm UV }}/{\\rm d}t =0$ and for which the equation of state is given by (\\ref{eos2}). In this case, the linearized Klein-Gordon equation can be simplified to yield \\begin{equation} \\label{lin3} \\frac{{\\rm d}\\delta \\varphi }{{\\rm d}t}+\\frac{1}{\\varphi _{\\rm inf}}\\left( \\frac{{\\rm d}\\varphi_{\\rm inf}}{{\\rm d}t}\\right) \\delta \\varphi = - {{3 \\alpha H_{\\rm inf}} \\over {m^2 \\varphi_{\\rm inf}}} \\rho_{_{\\rm UV}} \\, . \\end{equation} The solution of Eqs.~(\\ref{lin3}) and (\\ref{lin1}) is \\begin{eqnarray} \\delta \\varphi &=& -3\\sqrt{\\frac{\\kappa }{6}}\\left(\\frac{\\rho _{_{\\rm UV}}}{m}\\right) \\alpha t \\label{br1} \\, ,\\\\ \\delta H &=& \\frac{\\kappa }{6} \\frac{\\rho_{_{\\rm UV}}}{H_{\\rm inf}} \\left(1 - 3\\alpha H_{\\rm inf}t\\right) \\, . \\label{br2} \\end{eqnarray} As expected $\\delta \\varphi $ and $\\delta H$ vanish if $\\rho _{_{\\rm UV}}=0$. As stressed in \\cite{Unruh} and later analyzed in detail in \\cite{Ghazal}, it is important to express the result in terms of physical observables instead of in terms of the non-measurable background time coordinate $t$.  The obvious clock in our simple system is the scalar field $\\varphi \\equiv \\varphi _{\\rm inf}+\\delta \\varphi $ itself.  After a simple calculation it follows that the back-reaction effect $\\delta H$ measured in terms of $\\varphi$ is given by \\begin{equation} \\label{br3} H_{\\rm inf}+\\delta H=\\sqrt{\\frac{\\kappa }{6}} \\left(m\\varphi +\\sqrt{\\frac{\\kappa }{6}} \\frac{\\rho _{_{\\rm UV}}}{H_{\\rm inf}}\\right)\\, . \\end{equation} Thus, the effect of the ultraviolet back-reaction terms in a de Sitter background corresponds to a positive re-normalization of the cosmological constant. This result does not depend on the value of $\\alpha$.  \\par  So far, we have shown that a large ultraviolet back-reaction does not prevent inflation, in contrast to naive expectations. Instead, it leads to a re-normalization of the cosmological constant. On the other hand, as is evident from Eq.~(\\ref{br2}), the back-reaction terms can lead to a faster rolling of the scalar field $\\varphi_{\\rm inf}$. However, for the small values of $\\alpha $ of the order of $H_{\\rm inf}/\\mC$ indicated by our analysis, the increase in the rolling speed does not prevent a phase of inflation of sufficient length (\\ie having $|\\delta \\varphi | < |\\varphi _{\\rm inf}|$ for a time interval $t\\simeq H_{\\rm inf}^{-1}$) as long as \\begin{equation} n_k \\lsim {H _{\\rm inf}\\over \\mP} \\, , \\end{equation} which is a less stringent constraint than the one derived by Tanaka, and still allows observable effects on cosmic microwave fluctuations from trans-Planckian physics.    We have presented an attempt to study the effects of back-reaction on the trans-Planckian problem of inflationary cosmology. The question initially posed in Refs.~\\cite{Tanaka} and~\\cite{Starob} is whether the back-reaction of a state which consists of excited modes during the inflationary phase on scales smaller than the Hubble radius will prevent inflation. We have seen that the analysis is more subtle than it initially appears from the above works. We have shown that the back-reaction of an excited state preserving time-translation invariance does not prevent inflation, but simply leads to a renormalization of the Hubble constant. However, back-reaction will lead to a slightly faster rolling of the scalar field. As long as the occupation numbers $n_k$ are smaller than $H_{\\rm inf}/\\mP$, the rolling will be consistent with the standard inflationary paradigm. Such occupation numbers can lead to observable effects on the cosmic microwave background.  \\par  To render the analysis well defined, we have considered a dispersion relation for which the violation of adiabaticity is concentrated in a finite range of physical wavenumbers $\\Lambda _1< k_{\\rm phys} < \\Lambda _2$. We have assumed that all modes start off in their adiabatic vacuum state when $k_{\\rm phys} > \\Lambda_2$. They are squeezed while $\\Lambda_1< k_{\\rm phys} < \\Lambda _2$, and then emerge as excited states when $k_{\\rm phys} < \\Lambda_1$. Let us also notice that by making the wavelength interval $\\Lambda _1< k_{\\rm phys} < \\Lambda _2$ small, the energy density in the far ultraviolet modes can always be made small compared to the cutoff energy density $\\Lambda _1^4$ (or $\\Lambda _2^4$). Thus, in our approach the Planck energy density problem recently discussed in Ref.~\\cite{Jacob} in an approach to quantum field theory on a growing lattice, in which the number of fundamental field modes is increasing in time, does not arise (the analysis of Ref.~\\cite{Jacob} finds that the continual creation of modes at the cutoff scale yields an energy density which is of Planck scale). In our analysis, the Hilbert space of modes is time-independent, but as time proceeds an increasing subset of this space gets populated. Our mechanism for continual excitation of new modes appears to be more smooth than the lattice approach of Ref.~\\cite{Jacob}.  \\par  There are several important deficiencies in our analysis. First, we have used an ad-hoc regularization and renormalization prescription which consists of imposing effectively an abrupt cutoff in momentum space, and of subtracting the ground state energy of each field Fourier mode. This procedure is not covariant. It would be of interest to study our problem using a mathematically more rigorous regularization prescription, such as adiabatic regularization (see e.g. Ref.~\\cite{Gianni} and references therein).  \\par  Another serious concern is that we have considered matter fluctuations without taking into account the induced metric fluctuations. It is well known that the inclusion of metric fluctuations leads to dramatically different results for super-Hubble-scale perturbations~\\cite{MAB,ABM}. Thus, one might expect that the gravitational fluctuations could play an important role in the far ultraviolet region where $\\omega_{\\rm phys} < H_{\\rm inf}$. However, at the present time we are not able to study this issue because the non-standard dispersion relation has been set up for the matter sector only.  It is a challenge for future research to include the presence of non-standard dispersion relations consistently in both the gravitational and matter sector.   \\vskip0.5cm \\centerline{\\bf Acknowledgments}  We wish to thank Ted Jacobson and Martin Lemoine for interesting discussions. During the course of this work, R.~B. was supported by the Perimeter Institute for Theoretical Physics, and (at Brown) by the US Department of Energy under Contract DE-FG02-91ER40688, TASK~A. J.~M.  thanks the High Energy Group of Brown University and the Perimeter Institute for Theoretical Physics for warm hospitality."
    },
    "0410/astro-ph0410564_arXiv.txt": {
        "abstract": "This paper describes a concept for the high-accuracy absolute calibration of the instrumental polarization introduced by the primary mirror of a large-aperture telescope. This procedure requires a small aperture with polarization calibration optics (e.g., mounted on the dome) followed by a lens that opens the beam to illuminate the entire surface of the mirror. The Jones matrix corresponding to this calibration setup (with a diverging incident beam) is related to that of the normal observing setup (with a collimated incident beam) by an approximate correction term. Numerical models of parabolic on-axis and off-axis mirrors with surface imperfections are used to explore its accuracy. ",
        "introduction": "\\label{sec:intro}  Astronomical polarimetry requires frequent calibration operations to remove the instrumental polarization introduced by the various optical elements encountered by the light beam along its path. This contamination is usually determined by placing calibration optics early in the light path, which is used to feed light in a known state of polarization into the instrument. By measuring the polarization of the light that comes out of the system, it is possible to characterize it in terms of its Jones matrix (or, alternatively, the Mueller matrix if one is using the Stokes formalism).  Obviously, we can only characterize and calibrate the optical elements that the beam encounters {\\it after} the polarization calibration optics. Therefore, calibrating a telescope requires placing such optics at the telescope entrance, before the first reflection on the primary mirror (M1) occurs. This approach\\cite{SLMP+97} is being successfully employed for the Dunn Solar Telescope (at the Sacramento Peak Observatory, managed by the National Solar Observatory) and the German VTT on the island of Tenerife (at the Observatorio del Teide of the Instituto de Astrof\\' \\i sica de Canarias). In both cases, an array of linear polarizers and retarders are slided in the light path, on top of the entrance window, for calibration operations. Separate calibrations are obtained for the telescope and the instrument, so that the former does not need to be done as frequently. The Jones matrix of the complete system is then obtained as the product of the telescope and instrument matrices.  Unfortunately, entrance window polarizers are not practical for apertures larger than $\\sim$1~m diameter. In the past this has not been a major concern because: a)solar telescopes have apertures that do not exceed 1~m; and b)large astronomical telescopes have not been used for polarization measurements. However, this scenario is starting to change. Polarimetry is proving to be a very powerful tool to explore a broad range of Astrophysical problems\\cite{WS02}, resulting in a rapidly increasing interest to develop polarimeters for existing large telescopes. Second, the development of the Advanced Technology Solar Telescope (ATST)\\cite{KRH+02,KRK+03} demands a reliable method to calibrate a large telescope for high accuracy spectro-polarimetry.  Solar telescopes have been calibrated in the past by observing magnetic structures and making assumptions on the underlying physics. This poses important challenges, however, especially when pushing the envolope towards new observational domains. Consider for example the ATST, which is intended to do polarimetry at the $\\sim$10$^{-4}$ accuracy level. One of the common assumptions that is usually made in solar polarimetry is to consider that the continuum radiation is unpolarized. However, scattering processes can polarize the continuum and generate signals of the order of $\\sim$1\\% in the blue side of the visible spectrum. It has been stated \\cite{G03} that ``the direct observation of the polarisation of the continuous radiation still is a major outstanding observational challenge''. Considerations on the symmetry properties of Stokes profiles are not appropriate either. Gradients in the line-of-sight velocity or magnetic field introduce spurious asymmetries that can invalidate these assumptions. Furthermore, some physical processes operating in the atoms are known to induce Stokes asymmetries even in the absence of gradients. Spectral lines forming in the incomplete Paschen-Back regime become asymmetric\\cite{SNTBLdI04}. Moreover, the alignment-to-orientation conversion mechanism\\cite{KMN84} may also cause profile asymmetries.  In summary, it is very important to have an absolute calibration that does not rely on preconceived ideas on the objects under study. This is specially true when the instrumentation is being used to explore new scientific realms. The present work has been motivated by the challenge of calibrating the 4-meter primary mirror of the ATST to meet its very stringet polarization requirements. It might also be possible to use the calibration method proposed here in other existing large-aperture telescopes. However, the actual design of a practical implementation is beyond the scope of this paper. The main point of this work is to show that a calibration setup with inclined beam incidence can be used to measure the polarization properties of a large-aperture primary mirror. Geometrical effects can be calculated and removed from the measured Jones or Mueller matrices resulting in a good approximation to such matrices in normal observing conditions. ",
        "conclusions": "\\label{sec:conclusions}  This paper introduces a new concept to calibrate telescopes for astronomical polarimetry. The proposed method is particularly useful for modern large-aperture telescopes, for which it is probably the only practical procedure (at least for purely instrumental calibration). An accurate absolute calibration will be crucial for the new weak-signal science that the ATST will open. Existing night-time telescopes may also take advantage of this calibration procedure.  The Jones (and Mueller) matrix of an on-axis mirror is almost unaffected by the non-collimated incidence of the beam in the CS. For the particular configuration considered in \\S\\ref{sec:onaxis}, the Jones matrix obtained from the calibration is good to almost $10^{-5}$ with no need for any additional correction (see Eq~[\\ref{onMCS}]). This is in spite of the relatively large mirror imperfections in the simulation, which generate polarizing terms of the order of 6\\%.  An off-axis mirror suffers more instrumental polarization due to the asymmetric configuration. A mirror that produces instrumental polarization of a few percent can be calibrated to reach the $10^{-4}$ level (see Table~\\ref{table:offaxis2}). In addition to measuring the Jones matrix of the calibration setup ($M^{CS}$), it is also necessary to calculate the correction term $\\Delta M$. This calculation is straightforward, though, and $\\Delta M$ does not need to be recalculated unless the mirror is recoated or it degrades to a point where its {\\it average} refraction index changes significantly. Note that the $10^{-4}$ calibration accuracy includes some uncertainty on the average refraction index of the mirror.  \\appendix"
    },
    "0410/astro-ph0410087_arXiv.txt": {
        "abstract": "\\noindent In this talk I review how to identify cosmic ray accelerators that are high energy neutrino emitters. I also delineate the prospects for a new multi-particle astronomy: neutrons as directional pointers + antineutrinos as inheritors of directionality. ",
        "introduction": "Conventional astronomy spans about 18 decades in photon wavelengths, from $10^4$~cm radio-waves to $10^{-14}$~cm $\\gamma$-rays of GeV energy. Because the universe is opaque to photons of TeV energy and above (see Fig.~\\ref{gamf})~\\cite{Learned:2000sw}, present studies focus on hadrons, neutrinos, and gravitational waves as messengers probing the high energy universe. The best candidates to serve as messengers in a new astronomy of the high energy behavior of distant sources are neutral particles. This is because the orbit of a charged cosmic ray can be substantially bent, both by extragalactic magnetic fields~\\cite{Sigl:2004gi} and by the ambient magnetic field of our own Galaxy~\\cite{Stanev:1996qj}; destroying the possibility of locating the source. The most promising messenger is the neutrino: it can be copiously produced in cosmic beam dumps and can traverse unscathed dense astrophysical environments. In this talk I delineate the prospects to identify  high energy neutrino emitters.   \\begin{figure} \\label{gamf} \\centering \\leavevmode \\epsfysize=5.5cm \\epsfbox{gamf.eps} \\leavevmode \\epsfysize=5.4cm \\epsfbox{azukar.eps}\\\\ \\caption{Left: Mean interaction length for photons on the ultraviolet, visible, infrared, and microwave backgrounds~\\cite{Learned:2000sw}. Right: Upper end of the cosmic ray energy spectrum as observed by 5 different experiments (for details see main text).} \\end{figure} ",
        "conclusions": "We have analyzed the possibility of detecting the neutrino counterparts of TeV $\\gamma$-ray observations in a model where $\\gamma$-rays originate through $\\pi^0$-decay at the source. We have found that IceCube will attain sensitivity to observed neutrinos from the Crab Nebula. The discussion presented here can be easily generalized to other sources (see {\\it e.g.,}~\\cite{Anchordoqui:2004eu,Alvarez-Muniz:2002tn}). Indeed the AMANDA-II experiment has already achieved the sensitivity to probe the blazar Mrk 501 in its 1997 flaring state if the neutrino and $\\gamma$-ray fluxes are equal~\\cite{Ahrens:2003pv}. We have also estimated the ``essentially guaranteed'' $\\overline \\nu$ flux  originated in the decay of neutrons emitted by the Cygnus OB association. The expected event rate above 1 TeV at IceCube is $dN/dt|_{\\rm signal} \\sim 20$~yr$^{-1}$ antineutrino showers (all flavors) and a $1^\\circ$ directional signal of 4 $\\overline \\nu_\\mu$ events, well above atmospheric background.  The detection of (anti)neutrinos pointing towards TeV $\\gamma$-ray sources or to the $10^{9}$~GeV intriguing directional signals would not only provide a final answer to the origin of these particles, but also confirm the acceleration of protons/nuclei to ultrahigh energies."
    },
    "0410/astro-ph0410278_arXiv.txt": {
        "abstract": "We derive simple analytic expressions for the continuum lightcurves and spectra of flaring and flickering events that occur over a wide range of astrophysical systems. We compare these results to data taken from the cataclysmic variable SS Cygni and also with SN~1987A, deriving physical parameters for the material involved. Fits to the data indicate a nearly time-independent photospheric temperature arising from the strong temperature dependence of opacity when hydrogen is partially ionized. ",
        "introduction": "\\label{sec:intro}  In a recent paper \\citep{pearson03}, we explained the unusual flaring activity of the cataclysmic variable (CV) system AE~Aqr in terms of the aftermath of the collision between two gas clouds. We modelled the resulting fireball numerically, comparing the results to analytic approximations for the optical lightcurves and continuum spectra and to observed lightcurves and spectra. We set out here an improved analytic calculation for an expanding fireball with an LTE ionization structure and compare our results to more detailed numerical simulations. We fit our results to multiwavelength lightcurves of two very different systems; deriving values for the physical conditions in them.  Flickering and flaring occurs across the whole range of astrophysical systems from stars to active galaxies. It appears to be a recurrent feature of accreting systems and, in particular, those where an accretion disk is present. This flickering process may well be associated with the anomalously high viscosity present in these systems and represent the effect of magnetic reconnections such as that from the viscosity mechanisms of \\citet{hawley91}. Such a sudden localised deposition of energy would raise the gas temperature and the consequent overpressure would give rise to a local expansion of the disk material. With material cooling adiabatically, this expansion would rapidly become supersonic.  Observations by \\citet{bruch00} suggest that the flickering in CVs is not uniformly distributed across the disk but instead is associated with the stream-disk impact point and the innermost boundary layer region of the disk. Similarly \\citet{patterson81} found that the flickering in HT Cas was associated with the inner part of the accretion disk. These findings, however, are contradicted by \\citet{welsh95} for HT Cas, by \\citet{baptista04} for V2051 Oph and \\citet{baptista02} for the Low-Mass X-ray Binary (LMXB) X1822-371. All these studies show flickering to arise from a range of disk radii.  The terms ``flickering'' and ``flaring'' are often used interchangably when describing the stochastic variability of accreting sources and, even when defined in a particular field, are not necessarily used consistently. \\citet{warner95} and \\citet{baptista04} describe flickering in CVs as the continuous, random brightness fluctuations of 0.01--1~mag on timescales of seconds to dozens of minutes although the exact numerical ranges differ between the two. In contrast the unusal CV AE~Aqr is universally described as a ``flaring'' source since it has rarer by brighter events that often occur in batches. These events can raise the total optical luminosity of the system by factors of 2-3 contrasting with the 5-20\\% typical of CV flickering \\citep{bruch92} but on a similar timescale. Consequently, we adopt the convention of describing the small amplitude, continuous variations as ``flickering'' and reserve the term ``flaring'' for larger scale events.  The underlying model used to reproduce the AE Aqr lightcurves was based on consideration of the flux emitted by a hot, spherically symmetric ball of gas as it expanded and cooled. While we saw that the ``expanding fireball'' situation in AE~Aqr could arise from collisional heating of gas blobs, the same situation might arise from the very different causes outlined above. As a result, these ``fireball'' models may have a much wider range of applicablility than simply the unusual behaviour of AE Aqr. Accordingly, in this paper, we will compare analytic expressions derived for the lightcurves of such expansions to observations of SS Cygni and SN~1987A.  SS Cyg is a member of the dwarf nova subclass of cataclysmic variables. It consists of a $1.2 M_\\odot$ white dwarf accreting material through an accretion disk from a $0.71 M_\\odot$ K5V secondary that loses material via Roche lobe overflow \\citep{ritter03}. SS~Cyg was the second member of this subclass identified \\citep{wells96} and has had a near continuously monitored lightcurve for over a century \\citep{warner95}. Like all dwarf novae, SS Cyg exhibits optical variability over a large range of timescales e.g. outbursts ($\\Delta m\\approx3.5$ mag, $t\\approx40$ days), orbital modulations ($\\Delta m\\approx0.5$ mag, $P=0.275130$ days), flickering ($\\Delta m=0.01-0.2$ mag, $t\\sim1$ min) and Dwarf Nova Oscillations ($\\Delta m\\approx0.02$ mag, $t\\approx2.8-10.9$~s) \\citep{mauche01,teeseling97,honey89,warner95,warner04}. We shall concentrate our modelling on the lightcurve from a flickering event.  SN~1987A was the brightest supernova observed since Kepler's in 1604. The lightcurve can be broken into three phases: a ``flash'' lasting a few hours, a subsequent ``bump'' lasting around 4 months and a final exponential ``tail'' powered by radioactive decay of unstable nuclei in the ejecta. It was a type II supernova with an identified progenitor (Sk-$69^{o}202$) that was a blue supergiant. It is  believed that it lasted $\\sim10^{7}$~y as a 20-$30 M_{\\odot}$ main sequence blue supergiant star, before becoming a red supergiant for $\\sim10^{6}$~y during which phase it lost 3-$6 M_{\\odot}$ before finally becoming a blue supergiant again for a few thousand years \\citep{mccray93}. The rebrightening ``bump'' phases arises from a photosphere moving out with expelled material and eventually reversing as the material becomes optically thin again. It is interesting to note that the photosphere maintained a roughly constant temperature $\\sim6000$~K close to the recombination temperature for hydrogen at the expected densities in the expanding shell. The models of AE~Aqr flares in \\citet{pearson03} showed a similar isothermal behaviour. ",
        "conclusions": ""
    },
    "0410/astro-ph0410422_arXiv.txt": {
        "abstract": "We describe a new, distance-independent method for calculating the magnetic activity strength in low mass dwarfs, L$_{\\rm{H}\\alpha}$/L$_{\\rm{bol}}$. Using a well-observed sample of nearby stars and cool standards spanning spectral type M0.5 to L0, we compute ``$\\chi$'', the ratio between the continuum flux near H$\\alpha$ and the bolometric flux, f$_{\\lambda6560}$/f$_{\\rm{bol}}$. This ratio may be multiplied by the measured equivalent width of the H$\\alpha$ emission line to yield L$_{\\rm{H}\\alpha}$/L$_{\\rm{bol}}$. We provide $\\chi$ values for all objects in our sample, as well as fits to $\\chi$ as a function of color and average values by spectral type. This method was used by \\citet{we04} to examine trends in magnetic activity strength in low mass stars. ",
        "introduction": "Low mass stars are the most populous objects in our Galaxy, comprising roughly 70$\\%$ of the stars in the Solar Neighborhood and almost half of the Galactic stellar mass [\\citet{he97}]. Until the last decade, the low luminosities of these stars prevented their study by large scale spectroscopic surveys. However, efforts such as the Palomar/MSU Nearby Star Spectroscopic Survey [\\citet{re95a}, \\citet{ha96}, \\citet{gi02}, \\citet{re02}], 2MASS followup [\\citet{gi00b};  \\citet{gi00a}], and NLTT followup [\\citet{re04}, \\citet{le03}] have characterized global properties of nearby M dwarfs with samples numbering $\\gtrsim$ 1000 objects. Recently, the Sloan Digital Sky Survey (SDSS) has made it possible to study spectroscopic samples approaching 10,000 low mass stars [\\citet{we04}].  In particular, these surveys have provided insight into the magnetic behavior of low mass stars. There are many diagnostics by which a star may be deemed magnetically active, including X-ray and FUV emission lines, but in optical wavelengths, H$\\alpha$ is the most readily available signature of activity. Although H$\\alpha$ in either emission or absorption may indicate the presence of a chromosphere, most large-scale optical studies have used the more restrictive definition of H$\\alpha$ in emission. This practice is due to the difficulty in breaking the degeneracy between intermediate activity and truly weak chromospheres, both of which present as weak H$\\alpha$ absorption (see \\citet{gia82} and \\citet{gi02} for a discussion of this phenomenon). In this paper we use ``active'' to mean the presence of H$\\alpha$ in emission.  Increasingly large studies of active M dwarfs have built on one another to create our current picture of activity at the bottom of the main sequence. Originally, \\citet{jo74} showed a rise in the fraction of active M dwarfs towards later spectral types, with 100$\\%$ of M5 stars showing activity. The PMSU investigation of \\citet{ha96} extended the trend with much better statistics, and found that the fraction of active stars peaked at  $\\sim$60$\\%$ at spectral type M5.5. \\citet{gi00b} observed a sample of later type stars from 2MASS, and found that the incidence of activity rose to 100$\\%$ at type M7, but declined at later types. \\citet{bu00} added additional late M and L dwarfs to the sample and confirmed that the active fraction decreases past type $\\sim$M7 - M8. Most recently, \\citet{we04} used a much larger sample of low mass dwarfs from SDSS to further refine the trend: these data reveal a curiously elegant bell curve distribution for the incidence of magnetic activity in low mass stars. The fraction peaks at $\\sim$73$\\%$ of M8 stars being active.  Several of these studies consider the strength of activity, which is quantified as the ratio of the luminosity in H${\\alpha}$ to the bolometric luminosity. The activity strength remains relatively constant (though with large scatter) at log(L$_{\\rm{H}\\alpha}$/L$_{\\rm{bol}}$) $\\sim$-3.7 through the early to mid-  M spectral types, and appears to decline at types later than M6, while the active fraction is still rising. The increasing fraction of active stars may be an age effect (lower mass stars retain their activity longer) on the rising branch (M0 - M8).The activity strength results indicate that some damping mechanism apparently sets in around type M6. At types later than M8, this damping mechanism [possibly related to the atmospheric temperature, see \\citet{mo02}] overwhelms the age effect and results in fewer active dwarfs at lower atmospheric temperature. The activity fraction and activity strength results, taken together, provide important constraints on models of magnetic dynamos and magnetic field production in M and L dwarfs, particularly as these objects span the mass range where the stars become fully convective and develop degenerate cores.  Quantifying the strength of activity as the ratio between H${\\alpha}$ luminosity and bolometric luminosity, while informative, is prone to systematic error. Determining luminosities depends on having well-known distances, and as samples get larger (e.g. from SDSS) the stars are typically well beyond the immediate Solar Neighborhood (i.e. $>$ 25pc from the Sun). Thus, precise trigonometric parallax measurements are generally not available as in previous local surveys. In principle, one should be able to calculate the ratio L$_{\\rm{H}\\alpha}$/L$_{\\rm{bol}}$ and expect the distance dependencies to cancel. However, in order to get L$_{\\rm{bol}}$, one must have determined M$_{\\rm{bol}}$, which relies on knowing M$_V$, having excellent color measurements, and a reliable bolometric correction. M$_V$ generally comes from distance estimates based on spectroscopic and/or photometric parallax determinations, which are not always reliable or in agreement, due primarily to a paucity of calibration data at the bottom of the main sequence. Calculating L$_{\\rm{H}\\alpha}$, on the other hand, relies on knowing the flux in the continuum near H$\\alpha$. This is often obtained from the photometric colors [\\citet{re95b}]. Alternatively, it may be measured directly from spectrophotometrically calibrated data. Neither method is necessarily congruous with that used to obtain L$_{\\rm{bol}}$. Thus, in practice, the ratio usually contains two distances, typically determined in different ways, leading to possible systematic error.  An alternative method is to determine the ratio of the flux in the continuum near H${\\alpha}$ and the bolometric flux. This ratio, which we call $\\chi$, can then be multiplied by the equivalent width in H$\\alpha$ for a given star to give the ratio L$_{\\rm{H}\\alpha}$/L$_{\\rm{bol}}$. If one has excellent quality data, the ratio f$_{\\lambda6560}$/f$_{\\rm{bol}}$ may be calculated for each of the stars observed. In this paper, we use a well-observed sample of nearby M dwarfs and early L dwarfs to determine $\\chi$ as a function of several colors (I$_C$-K$_s$, V-I$_C$, $r - i$, $i - z$) and of spectral type, in effect providing an improved ``bolometric correction'' for low mass dwarfs. ",
        "conclusions": "In order to study the strength of magnetic activity as a function of spectral type, we require a distance-independent way of determining the ratio L$_{\\rm{H}\\alpha}$/L$_{\\rm{bol}}$. We used a sample of nearby stars and cool standards from 2MASS, spanning spectral types M0.5 to L0, with ranges in color 1.89 $\\leq$ I$_C-K_s$ $\\leq$ 5.12, and 1.91 $\\leq$ V$-I_C$ $\\leq$ 4.40. For the 8 parsec sample stars, we obtained the bolometric flux from photometry combined with bolometric corrections. For the 2MASS standards, optical photometry was obtained by convolving the spectra with filter response, then bolometric corrections were applied. The flux in the continuum near H$\\alpha$ was obtained directly from the spectra. In both cases, we computed SDSS photometry for the sample using color transformations, provided in Section 3. We then found the ratio $\\chi$ = f$_{\\lambda6560}$/f$_{\\rm{bol}}$. In Section 3, we provided fits to $\\chi$ as a function of color (V - I$_C$, I$_C$ - K$_s$, $r - i$, $i - z$). Table 1 gives average $\\chi$ values at each spectral type in the event photometry is not available, while Table 2 provides $\\chi$ values for all stars in our sample. $\\chi$ may be multiplied by the measured equivalent width in H$\\alpha$ for a given star to yield the activity strength ratio, L$_{\\rm{H}\\alpha}$/L$_{\\rm{bol}}$. If the reader possesses excellent quality data,  the method described here may also be used to calculate f$_{\\lambda6560}$/f$_{\\rm{bol}}$ for individual program stars.  This method has already yielded cleaner results for activity strength calculations in \\citet{we04}, better illuminating trends in magnetic behavior at the end of the Main Sequence. Improved color transformations and a larger sample of stars with well-calibrated photometry and spectra are required to reduce the scatter in these relations."
    },
    "0410/astro-ph0410614_arXiv.txt": {
        "abstract": "The redshift of the galaxy lensing HE~0435$-$1223 is 0.4546 $\\pm$ 0.0002, based on observations obtained with the Low Dispersion Survey Spectrograph 2 (LDSS2) on the Magellan Consortium's 6.5~m Clay telescope.  {\\it Hubble Space Telescope}/ACS observations of the system also reveal a spiral-rich group of 10 galaxies within 40\\arcsec\\ of the elliptical lensing galaxy.  The redshifts for two of these galaxies were measured to be in the foreground (at $z=0.419$) with respect to the lens, thus at least some of the nearby galaxies are not part of the same physical group as the lensing galaxy. Mass models of the system (assuming same-plane deflectors) that take the local group environment into account do better at explaining the observed emission-line flux ratios (which are presumably unaffected by microlensing) than single halo models, but the match is still not perfect.  In particular, component A (a minimum of the light travel time) is observed to be 0.20 mag brighter than predicted and component C (also a minimum image) is observed to be 0.16 mag fainter than predicted.  Mass models for the system predict an A--D time delay of either 15.8 or 17.6 days ($H_{o} = 72$ km s$^{-1}$ Mpc$^{-1}$) depending on the details of the local galaxy environment. ",
        "introduction": "For a variety of reasons, follow-up observations of gravitationally lensed quasars have not kept pace with the lens discovery rate. For example, both lens and source redshifts are required to convert a measured time-delay into an estimate of the Hubble constant (Refsdal 1964) or to measure the mass-to-light evolution of lensing galaxies (Rusin et al. 2003), but only half of the $\\sim$80 known systems have complete redshift information.  Poor knowledge of the lens galaxy environment also hinders lensing applications.  Nearby groups or clusters can bias estimates of the Hubble constant by contributing some fraction of the lensing mass convergence (Gorenstein, Shapiro \\& Falco 1988; Saha 2000), and such structures are likely responsible for the large ($\\sim10\\%$) shears required to model quadruple systems (Holder \\& Schechter 2003) and for explaining the observed quad-to-double ratio (Keeton \\& Zabludoff 2004).  It is therefore important to understand the characteristics of each lens system as thoroughly as possible.  In this paper, we report our measurement of the lens redshift for the quadruple quasar HE~0435$-$1223 and explore the system's local galaxy environment from {\\it HST}/ACS imaging.  \\begin{figure*}[t] \\plotone{NickMorgan.fig1.ps} \\caption{{\\it R}-band image of HE0435$-$1223 and field obtained with the du Pont 2.5-m telescope at Las Campanas. The panel spans 7.1\\arcmin\\ across.  Targets observed with LDSS-2 are labeled (S1, S2, lens, G1-17).  {\\it Inset}: Close-up of HE~0435$-$1223 from Figure~2 of Wisotzki et al. (2002).  The parallel lines show the slit orientation and width used for LDSS2 spectroscopy.} \\label{fig:fc_dupont} \\end{figure*}  HE~0435$-$1223 was discovered to be a $z~=~1.689$ quasar during the Hamburg/ESO (HES) survey for bright quasars in the Southern Hemisphere (Wisotzki~et~al.~2000).  The quasar was later found to be gravitationally lensed by Wisotzki et al. (2002) during snapshot followups of HES quasars with the 6.5~m Baade telescope at the Las Campanas Observatory (LCO).  The image morphology is a symmetric quad with image separations of 2\\farcs3 and 2\\farcs6 along the long axes of the configuration, somewhat reminiscent of the Northern Hemisphere ``Einstein Cross'' Q~2237+0305 (Huchra~et~al.~1985).  The lensing galaxy was prominently detected in the Baade discovery images, with its light profile and $gri$ aperture colors suggestive of a $0.3<z<0.5$ elliptical galaxy.  The system holds high promise for measuring a differential time delay and estimating the Hubble constant.  Wisotzki~et~al.~(2002) cite evidence for both short-term (20\\% over $\\sim$2 months) and long-term (1 magnitude over $\\sim$ 12 years) variability in the total quasar flux.  The symmetric image configuration also implies relatively short differential time delays (on the order of several weeks), which ought to make it straightforward to distinguish the time-delay signature from longer-timescale microlensing variability.  Wisotzki et al. (2003) provided a refined galaxy redshift estimate of $z=0.44 \\pm 0.02$ based on integral field observations of the system. However, their estimate was obtained from low-resolution ($\\Delta\\lambda = 300$\\AA) spectral rebinning that traced the galaxy's overall spectral energy distribution rather than from the detection of stellar absorption features.   In \\S\\ref{ldss2}, we report a precise spectroscopic measurement of the lensing galaxy redshift of $z=0.4546\\pm0.0002$ using the Low-Dispersion Survey Spectrograph 2 (LDSS2) on the Clay 6.5~m telescope. In \\S\\ref{hstacs}, we make use recent {\\it HST}/ACS observations of the system to explore the HE~0435$-$1223 galaxy environment and find evidence for a spiral-rich group of at least 10 galaxies within 40\\arcsec, at least two of which are in the foreground (at $z=0.419$) with respect to the lens. Overall the data suggest a complex lensing environment.  We explore several lens models for the system using the {\\it HST}/ACS astrometry and emission-line flux ratios from Wisotzki et al. (2003) in \\S\\ref{models}, and discuss the model implications and time-delay predictions in \\S\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  \\begin{figure}[t] \\plotone{NickMorgan.fig6.ps} \\caption{Diagram showing relative galaxy positions (filled circles). Dot sizes are proportional to the respective galaxy's $I$-band flux within a 2\\arcsec\\ aperture.  The diagonal lines show the shear directions for the IEx (dotted) and IEISx (solid) lens models.  The star marks the flux-weighted center of the G20-G21-G12 group of galaxies.} \\label{fig:shear} \\end{figure}  The improved flux predictions and closer alignment between the shear direction and the NW group are arguments for favoring the multiple halo model, but the shear still does not point to any obvious perturber.  It points between the faint G20 and G21 galaxies as seen from the lens.  This is about +20\\degr\\ from the G20-G21-G12 flux-weighted center and about +30\\degr\\ from G12, the brightest of all neighboring galaxies within 40\\arcsec\\ of the lens (see Figure~6).  We {\\it know} that G12 is at essentially the same redshift as the lensing galaxy, so if the only significant perturbers were G12 to the NW and G22 to the SW, then we would expect the IEISx shear to point directly toward G12.  It does not, which suggests that if the flux ratios are to be explained solely by the macromodel, then the lensing potential must be even more complicated than the double halo model considered here.  Alternatively, the macromodel could be essentially correct and a separate effect could be perturbing the flux ratios.  In general, there are four mechanisms that could alter the image fluxes: microlensing, quasar variability, differential dust extinction and substructure lensing. Presumably microlensing is not an issue since we are working with the emission-line fluxes.  Quasar variability is also not a factor since the A$-$C time delay is less than a day in all models, much too short compared to quasar variability timescales (typically weeks to months).  Wisotzki et al. (2002) already commented that extinction is probably not very important for this system since the lensing galaxy is an elliptical and the Magellan colors of the four quasars agreed to $\\lesssim$0.05 mag. In general, Falco et al. (1999) found median differential extinctions among a subsample of optically-selected lenses of $\\Delta E(B-V)=0.04$ mag.  Only 10\\% of their sightlines through elliptical galaxies have $0.15<\\Delta E(B-V)<0.25$, so the $\\sim$0.2 mag of extinction needed to reconcile component C's relative faintness is unlikely but not impossible.   However, the PMAS resolved spectral observations of the four quasars obtained by Wisotzki et al. (2003) show the quasar continuum slopes to be virtually identical (L. Wisotzki, private communication).  This would be a strong coincidence if any significant dust extinction were present, and of course rules out significant differential extinction.  The remaining explanation is substructure lensing.  Millilensing by low-mass satellites (typically modeled as NFW profiles or truncated isothermal halos) in the lens galaxy can either brighten or dim image fluxes by several tenths of a magnitude.  The effect has been modeled both for individual systems (Metcalf \\& Zhao 2002) and in a statistical sense (Dalal \\& Kochanek 2002), and typically requires the subclumps to comprise a few percent by mass of the primary galaxy's dark-matter halo.  Although we do not consider substructure models here, such an effect remains the most plausible explanation of the HE~0435$-$1223 emission-line fluxes given the present data.  We do note that substructure lensing by strictly isothermal halos, either embedded in the lensing galaxy or along the line of sight, cannot account for the observed flux ratios.  Isothermal mass clumps always brighten positive-parity images (Keeton 2003), but components A and C, both positive-parity images, are respectively brighter and fainter than predicted. Of course the unlensed quasar flux is a free parameter, so one could dim the source such that both A and C became brighter than predicted, but then images B and D (both negative-parity images) would be $\\sim$0.2 mag brighter than predicted as well.  Millilensing (and microlensing as well) tends to dim saddle-point images (Keeton 2003, Schechter \\& Wambsganss 2002), so the situation using isothermal halos becomes somewhat contrived.  We can predict the differential time delays for the system using the measured lens redshift of $z=0.4546\\pm0.0002$ from \\S\\ref{ldss2}.  The A$-$D delay is the longest delay and therefore the most interesting to measure.  The ISx, IEx, and IEISx models predict A$-$D delays of 12.21, 15.80, and 17.54 days, respectively.  The elliptical halo (IEx) is formally preferred over the spherical halo (ISx) with $\\chi^2/dof$ of 3.6 compared to 2.4, where the improvement comes wholly from the quasar and galaxy positions.  The double halo (IEISx) is preferred over the single halo (IEx) with $\\chi^2/dof$ of 2.4 compared to 1.9, where the improvement comes mostly from the images fluxes.  The last two models yield 10\\% differences in the A$-$D time delay, which gives an estimate of the error that one can expect from uncertainties in the local galaxy environment.  There are several ways to improve the time delay prediction.  Redshift information for the remaining 8 of 10 neighboring galaxies, especially for G22, would significantly improve the model confidence.  In the least, it would point to which galaxies might share a common dark matter halo and thus guide the choice of parametric models.  The {\\it HST}/ACS images already show a partial Einstein ring for the system in $I$-band, and one would suspect a complete Einstein ring to be visible in the infrared.  The extra constraints offered by modeling the ring may be able to distinguish between the single and double halo models without recourse to the image flux ratios. For example, it would be interesting to see if the shear direction for a ring-constrained IEISx model points more in line with the center of the NW group than found here."
    },
    "0410/astro-ph0410491_arXiv.txt": {
        "abstract": "We present thermal-infrared images of three extreme bipolar objects, M~2-9, Mz~3, and He~2-104.  They are bipolar planetary nebulae with bright central stars and are thought to be powered by symbiotic binary systems.  The mid-infrared images spatially resolve the spectral energy distributions of the central engines from the surrounding nebulae.  A warm dust component of several hundred degrees can account for the core emission, while a cooler component of $\\sim$100 K produces the more extended emission from the bipolar lobes.  In every case, the dust mass for the unresolved core region is orders of magnitude less than that in the extended lobes, raising doubts that the hypothetical disks in the core could have been responsible for pinching the waists of the nebulae.  We find total masses of roughly 0.5--1 M$_{\\sun}$ in the nebulae of M~2-9 and Mz~3, requiring that this material was donated by intermediate-mass progenitor stars.  The mass of He~2-104's nebula is much lower, and any extended emission is too faint to detect in our images.  Extended dust emission is detected around both M~2-9 and Mz~3, in both cases resembling the distribution of ionized gas.  Our images of Mz~3 have the highest signal-to-noise in the extended polar lobes, and we show that the fairly uniform color temperature derived from our images can explain the 110~K dust component that dominates the far-infrared spectral energy distribution.  In the case of Mz~3, most of the mass traced by dust is concentrated at high latitudes, and we note possible evidence for grain destruction in shocks indicated by an anticorrelation between [Fe~{\\sc ii}] and dust emission.  Except for these regions with enhanced [Fe~{\\sc ii}] emission, the dust continuum closely resembles the distribution of ionized gas. ",
        "introduction": "The formation of bipolar structure in planetary nebulae (PNe) is one of the enduring challenges to understanding the late stages of stellar evolution.  No simple model has yet emerged that can account for how a diffuse slowly-rotating asymptotic giant branch star develops strong bipolarity (not to mention point symmetry, multipolar outflows, jets, etc.) as it becomes a PN, although confining circumstellar disks, disk/jet precession, core rotation, or magnetic fields may play important roles (reviews of shaping mechanisms are given by Balick \\& Frank 2002; Frank 1999; Mellema \\& Frank 1995; and contributions in Meixner et al.\\ 2004).  One can alleviate the mental strain required to invent self-consistent single-star models by invoking binary systems to provide an axis of symmetry and angular momentum, which are necessary for bipolar phenomena.  The ``Butterfly'' (M~2-9), the ``Ant'' (Mz~3), and the ``Southern Crab'' (He~2-104) are among the most dramatic bipolar PNe with tightly-pinched waists seen in {\\it Hubble Space Telescope} ({\\it HST}) images (Balick \\& Frank 2002; Balick 1999; Balick et al.\\ 1997; Corradi et al.\\ 2001).  It is interesting that each of these is thought to be powered by a symbiotic binary system (Whitelock 1987; Balick 1989; Corradi \\& Schwarz 1993; Corradi et al.\\ 2000; Schmeja \\& Kimeswenger 2001; Smith 2003).  These objects have been studied extensively with spectra at visual and near-infrared (IR) wavelengths. They have bright but highly-reddened stellar cores with rich emission-line spectra, accompanied by faint nebular spectra from lower density regions in their polar lobes (Evans 1959; Cohen et al.\\ 1978; Allen \\& Swings 1972; Swings \\& Andrillat 1979; Balick 1989; Lutz et al.\\ 1989; Goodrich 1991; Hora \\& Latter 1994; de Freitas Pacheco \\& Costa 1996; Phillips \\& Cuesta 1999; Smith 2003); M~2-9 and Mz~3 in particular are spectroscopic twins at visual and near-IR wavelengths except for the lack of H$_2$ in Mz~3 (Smith 2003).  Each object requires an ionization source with a temperature of order 30,000 K, and the nebulae seem to have moderately-enhanced abundances of He and N, suggesting intermediate-mass progenitors.  Published distances to each object vary widely, scattered around 1-2 kpc; in this study we quote numerical values like M and L for a nominal distance of 1 kpc.  By comparison, little is known about the nature of dust in these sources.  The spatial distribution of thermal-IR emission can constrain the dust temperature and mass in the core regions and polar lobes separately, in order to then estimate important quantities like the total mass of the nebulae and the mass of putative circumstellar disks.  Low-resolution 8-13 $\\micron$ spectra of M~2-9 and Mz~3 show similar, nearly flat continuum emission in both sources, except for the presence of [Ne~{\\sc ii}] 12.8 $\\micron$ in Mz~3 (Aitken \\& Roche 1982), but these observations did not include any spatial information. Although each object was observed with the {\\it Infrared Astronomical Satellite} ({\\it IRAS}), the low spatial resolution of those data cannot separate the relative contributions of the core and polar lobes.  Ground-based images in the thermal-IR with arcsecond-scale spatial resolution have only been published for Mz~3 (Quinn et al.\\ 1996), but that was one image in a single broad filter, making it difficult to constrain the dust temperature or mass.  In an effort to clarify the spatial distribution of mid-IR properties of M~2-9, Mz~3, and He~2-104, we obtained 8-25 $\\micron$ images of these sources with MIRLIN on the IRTF and TIMMI2 on the ESO 3.6m telescope. ",
        "conclusions": "We presented thermal-IR images of the bipolar planetary nebulae M~2-9, Mz~3, and He~2-104, obtained at the IRTF and at ESO.  The distribution of dust in these systems is of interest as each object is strongly bipolar with a bright central star, and each is thought to be powered by a symbiotic binary system.  Our mid-IR images allowed us to construct SEDs of the bright, unresolved central engines separate from the surrounding nebulae.  In each case, we find that a warm dust component of several hundred degrees can account for the core emission (260, 320, and 320 K in M~2-9, Mz~3, and He~2-104, respectively), while a cooler component produces the more extended emission from the bipolar lobes (85, 110, and 95 K in M~2-9, Mz~3, and He~2-104, respectively).  Hot dust at temperatures up to $\\sim$900 K is needed to fit the near-IR SEDs of each source, but these additional components contribute negligible mass.  In every case, the dust mass for the unresolved core region is orders of magnitude less than the mass of the extended lobes, raising doubts that the hypothetical disks in the core could have been responsible for pinching the waists of the bipolar nebulae.  Assuming a gas:dust mass ratio of 230 appropriate for graphite grains and solar composition, we find masses of roughly 0.8 and 0.6 M$_{\\sun}$ in the nebulae of M~2-9 and Mz~3, respectively.  These relatively high masses require that this material was donated by intermediate-mass progenitor stars.  The mass of He~2-104's nebula is much lower, in accordance with the fact that we detect no extended dust emission from its much fainter nebula.  We detected extended dust emission around both M~2-9 and Mz~3, and in both cases it resembled the distribution of ionized gas seen in narrowband emission-line images.  Our images of Mz~3 have the highest signal-to-noise in the extended polar lobes, and we show that the fairly uniform color temperature derived from these images can explain the 110~K dust component that dominates the far-IR spectral energy distribution.  Thus, the mass derived from the cool 110~K component traces material in the polar lobes.  In the case of Mz~3, the dust is concentrated at high latitudes.  We also note possible evidence for grain destruction in shocks, indicated by an anticorrelation between dust seen in our images and [Fe~{\\sc ii}] emission in previously-published near-IR images.  While dust appears to be absent or weak in these regions (i.e. the polar ``blisters'' of Mz~3), 12.8 $\\micron$ [Ne~{\\sc ii}] emission is detected.  Except for these regions with enhanced [Fe~{\\sc ii}] emission, the thermal-IR continuum from dust closely resembles the distribution of ionzed gas.  Given the dust temperatures of the as-yet unresolved central sources, where dust presumably resides in a circumstellar disk, each of these bright objects may prove to be worthwhile targets for mid-IR interferometric observations to resolve the disks."
    },
    "0410/astro-ph0410458_arXiv.txt": {
        "abstract": "{ We present optical integral field observations of the H II region containing the ultraluminous X-ray source Holmberg II X-1. We confirm the existence of an X-ray ionized nebula as the counterpart of the source due to the detection of an extended He II $\\lambda4686$ region (21$\\times$47 pc) at the {\\it Chandra} ACIS-S position. An extended blue objects with a size of 11$\\times$14 pc is coincident with the X-ray/He II $\\lambda4686$ region, which could indicate either a young stellar complex or a cluster. We have derived an X-ray to optical luminosity ratio of L$_{X}/$L$_{B}\\ge170$, and presumable it is L$_{X}/$L$_{B}\\sim300-400$ using the recent HST ACS data. We find a complex velocity dispersion at the position of the ULX. In addition, there is a radial velocity variation in the X-ray ionized region found in the He II emission of $\\pm50$ km s$^{-1}$ on spatial scales of 2-3$^{\\prime\\prime}$. We believe  that the putative black hole not only ionizes the surrounding HII gas, but also perturbs it dynamically (via jets or the accretion disk wind). The spatial analysis of the public {\\it Chandra} ACIS-S data reveals a point-like X-ray source and marginal indication of an extended component ($\\ll 15$ \\% of the total flux). The {\\it XMM-Newton} EPIC-PN spectrum of HoII X-1 is best fitted with an absorbed power-law in addition to either a thermal thick plasma or a thermal thin plasma or a multi-colour disk black body (MCD). In all cases, the thermal component shows relatively low temperatures (kT$\\sim0.14-0.22$ keV). Finally we discuss the optical/X-ray properties of HoII X-1 with regards to the possible nature of the source. The existence of an X-ray ionized nebula coincident with the ULX and the soft X-ray component with a cool accretion disk favours the interpretation of an intermediate-mass black hole (IMBH). However the complex velocity behaviour at the position of the ULX indicate a dynamical influence of the black hole onto the local HII gas. } \\titlerunning{Integral field spectroscopy of HoII X-1} ",
        "introduction": "The two main classes of discrete galactic X-ray sources, X-ray binaries and supernova remnants (SNR), have been known and relatively well understood for decades. The third class of  galactic X-ray sources was detected with Einstein due to its high spatial resolution and its large collecting area (Fabbiano \\cite{Fab89}). These objects have become known as Ultraluminous X-ray sources (ULX) and they have 0.5-10 keV luminosities of 10$^{39-41}$ erg s$^{-1}$, generally higher than black hole binaries such as Cyg X-1 and SMC X-1, but lower than that of active galactic nuclei (AGN). Assuming Eddington luminosities, this corresponds to accretion onto black holes of masses between ten and several hundred solar masses suggesting intermediate-mass black holes (IMBHs; see Colbert \\& Mushotzky \\cite{Col99}, Miller \\& Colbert \\cite{Mil04} for a review). ULX are not located in the dynamical center of their host galaxies and thus they are not caused by sub-Eddington accretion onto a central AGN-type super-massive black hole. The identification of the optical counterparts of ULX is one of the most important issues to determine the nature of these objects. The number of optically identified ULX is still limited, only a small number of reliable optical counterparts are known to date. Several ULX seem to be associated with H II regions or nebula (e.g., Pakull \\& Mironi \\cite{Pak01}, Foschini et al. \\cite{Fos02}, and Wang \\cite{Wa02}). However, some ULX have been associated with an accreting black hole in a globular cluster (see Angelini et al. \\cite{Ang01}, Wu et al. \\cite{Wu02}). The optical counterpart to an extremely luminous X-ray source near Holmberg IX is a shock-heated nebula, which is associated with a optically faint non-stellar source (Miller \\cite{Mil95}).  The discovery of the intense He II $\\lambda4686$ nebular recombination line in Holmberg II X-1 (hereafter; HoII X-1) indicates that the interstellar medium probably reprocesses part of the X-ray luminosity of $\\sim 10^{40}$ erg s$^{-1}$.  Assuming quasi-isotropic emission (Pakull \\& Mironi \\cite{Pak01}) and the distance of $D=3.2$ Mpc, the X-ray luminosity of HoII X-1 corresponds to the Eddington mass of $\\approx$ 80 M$_{\\sun}$, which is considered a rather strict lower limit to the mass of accreted compact object if the X-ray emission is isotropic.  A more accurate distance to HoII of 3.39 Mpc was recently determined by Karachentsev et al. (\\cite{Kar02}). The new distance estimate would result in only a minimal  difference in the size, flux, and luminosity values we have calculated assuming  a value of $D=3.2$ Mpc.  The ROSAT HRI and PSPC data of HoII X-1 were presented in detail by Zezas et al. (\\cite{Ze99}), revealing a point-like, variable source (on scales of days and years) at the edge of the compact H II region \\#70 (Hodge, Strobel \\& Kennicut \\cite{Ho94}). The ROSAT PSPC spectrum was best described by either a steep power-law with $\\Gamma=2.7$ or a thermal plasma with $kT\\sim 0.8$ keV. Miyaji, Lehmann \\& Hasinger (\\cite{Mi01}; hereafter MLH01) found that the ASCA spectrum extends to harder energies. The hard part of the spectrum is best fitted with a flatter power-law with $\\Gamma\\sim1.9$ and intrinsic absorption above that of our galaxy. A multi-colour disk black body model (MCD, Mitsuda et al. \\cite{Mit84}) did not fit the ASCA spectrum of HoII X-1, unlike some other ULX's of similar luminosities. A joint PSPC-ASCA spectral analysis showed a soft excess component above the power law component. The soft excess could be described by either a MCD model with $kT_{in}=0.17$ keV, or a thin thermal plasma with $kT=0.3$ keV. MLH01 disfavored the MCD interpretation for the soft excess based on the large discrepancy between the black hole mass estimated from $kT_{in}$ and that estimated from the normalization. Furthermore, the spatial analysis of the ROSAT HRI image indicates an extended component.  In this paper we report on optical integral field and long-slit spectroscopic observations of the ultraluminous X-ray source X-1 in Holmberg II, as well as on results from the spatial and spectral analysis of the public {\\it XMM-Newton} and {\\it Chandra} data.  The outline of this paper is as follows. The optical imaging, and the long-slit and integral field spectroscopy of HoII X-1 are presented in Sect. 2. The H II region associated with the X-ray source is presented in Sect. 3. The optical properties (e.g., emission line flux maps, velocity dispersion, and radial velocities) of the H II region are described in Sect. 4. In Sect. 5 we present the X-ray spectral and spatial analysis based on public {\\it XMM-Newton} EPIC-PN and {\\it Chandra} ACIS-S data. The implications of our results with respect to the nature of the ULX HoII X-1 are discussed in Sect. 6.  \\begin{table*}[t] \\begin{center} \\caption[]{Instruments and observations} \\label{instr} \\begin{tabular}{|l|llll|} \\hline \\noalign{\\smallskip} & \\multicolumn{2}{c}{{\\bf integral field spectroscopy}} & {\\bf long-slit spectroscopy} & {\\bf imaging}\\\\ \\hline \\noalign{\\smallskip} telescope           & Calar Alto 3.5m     & SAO 6m       & SAO 6m & CFHT 3.6m\\\\ instrument          & PMAS                 & MPFS          & LSS  & OSIS \\\\ \\hline \\noalign{\\smallskip} date of observation & 28.10.2001           & 15.03.2002    & 14-15.01.2002 & 11.03.2000\\\\ exposure time       & 4$\\times$900s mosaic,- offsets: 4$^{\\prime\\prime}$       & 2$\\times$900s  & 1800s (N1-3) & 2$\\times$1200s (H$\\alpha$) \\\\ &                 & & 2700s (N4-5)    &2$\\times$300s ($R$) \\\\ &                 &           &   &2$\\times$500s ($B$) \\\\ seeing              & 1.0$^{\\prime\\prime}$                &  1.5--2.0$^{\\prime\\prime}$        & 1.5$^{\\prime\\prime}$/1.1$^{\\prime\\prime}$ & $\\sim0.7^{\\prime\\prime}$\\\\ field of view       & 8$^{\\prime\\prime}$$\\times$8$^{\\prime\\prime}$, 12$^{\\prime\\prime}$$\\times$12$^{\\prime\\prime}$ (mosaic)            & 16$^{\\prime\\prime}$$\\times$16$^{\\prime\\prime}$     & 2$^{\\prime\\prime}\\times140^{\\prime\\prime}$ & 3$^{\\prime}$$\\times$3$^{\\prime}$\\\\ pixel scale         & 0.5$^{\\prime\\prime}$/pix             & 1.0$^{\\prime\\prime}$/pix      & 0.41$^{\\prime\\prime}$/pix& 0.088$^{\\prime\\prime}$/pix\\\\ image size in pixels& 2048$\\times$4096 (binned 2$\\times$2)    & 1024$\\times$1024 & 1024$\\times$1024 & 2048$\\times$2048\\\\ spectral coverage   & 4450-5140 \\AA        & 4210-6820 \\AA & 4330-6750 \\AA  & -\\\\ dispersion          & 0.76 \\AA/binned pix   & 2.7 \\AA/pix    & 2.4 \\AA/pix & -\\\\ spectral resolution & 1.5 \\AA              & 7.0 \\AA        & 7.2 \\AA &-\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "Before we discuss the nature of the ultraluminous X-ray source HoII X-1 we summarize our results:  \\begin{enumerate} \\item{An extended He II $\\lambda4686$ region (21$\\times47$ pc) inside an H II region is positionally coincident with the HoII X-1 source. The existence of such an extended He~II emission region is very rare in H II regions. The He II luminosity agrees well with those published by Pakull \\& Mirioni (\\cite{Pak01}) and Kaaret et al. (\\cite{Kaa04}).} \\item{Neither the emission line fluxes nor the line flux ratios of the lines other than He II $\\lambda4861$ are correlated with HoII X-1. The emission line flux ratios agree well with those of normal H II regions. This result excludes a SNR nature of the radio source inside the Holmberg II region \\#70, which is probably not coincident with the X-ray source. } \\item{We find a blue extended counterpart for the ULX with a size of 11$\\times14$ pc, which is most probably a young stellar complex or a cluster. We have derived an X-ray to optical luminosity ratio of L$_{X}/$L$_{B}\\ge170$ for the blue stellar complex. The recently published {\\it HST} ACS images of HoII X-1 (Karret et al. \\cite{Kaa04}) confirm this, however they suggest a ratio of L$_{X}/$L$_{B}=300-400$.} \\item{The increased FWHM of $\\sim50$ km s$^{-1}$ in a scale of $\\sim20-30$ pc at the position of HoII X-1 is most likely not directly gravitational related  to the rotation around a black hole, as the resulting black hole mass would be $\\sim10^{7}$M$_{\\sun}$. Further, there is a radial velocity variations inside the He II region of $\\pm50$ km s$^{-1}$ on spatial scales of $\\sim2^{\\prime\\prime}$ ($\\pm30$ pc in projection). The complex velocity field at the location of the ULX suggests that the putative black hole ionizes the surrounding HII gas and perturbs it dynamically (via jets or the accretion disk wind).} \\item{The spatial analysis of the {\\it Chandra} ACIS-S image gives marginal indications for an extended emission component with an extended flux fraction well below 15 \\%. This could be consistent with the thin thermal plasma component in the {\\it XMM} spectrum. } \\item{The {\\it XMM} EPIC-PN spectrum is well fitted  with an absorbed power-law ($\\Gamma\\sim2.6$) for the hard component and a thermal emission for the soft component, which can be due to a thick plasma (bbody), a thin plasma (mekal), or a multicolor black body disk with relatively low inner disk temperature ($kT_{in}\\sim0.16$      keV). The intrinsic absorption of N$_{H}\\sim1.6 \\times 10^{21}$ cm$^{-2}$ is greater than the galactic absorption. } \\end{enumerate}  There are three models for ULX extensively discussed  in the literature:  i) ULX may be black holes of ''normal'' stellar masses ($\\sim 10$M$_{\\sun}$) in binaries, which accrete gas in a supercritical regime. They are SS\\,433-like objects or microquasars in their transient activity states, whose hard radiation can be collimated (and beamed) along jets and accretion disks axes (Fabrika \\& Mescheryakov, \\cite{Fa01}, King et al., \\cite{Ki01}).  ii) ULX may be black holes with few tens of solar mass ($\\sim 10$M$_{\\sun}$) and that their X-ray emissions is from the disk shining at super-Eddington luminosities (Begelman \\cite{Beg02}). The slim disk model (Abramowicz et al. \\cite{Abra88}, Ebisawa et al. \\cite{Ebi03}) allows to explain the observed super-Eddington luminosity, hard X-ray spectra and spectral variations.  iii) The ULX may be intermediate mass black holes (IMBHs), $\\sim 10^3$M$_{\\sun}$ (Colbert \\& Mushotzky \\cite{Col99}), which were formed from the very first stars (Madau \\& Rees \\cite{Mad01}) or in globular clusters (Miller \\& Hamilton \\cite{Mil02}). The IMBH could accrete gas from a close companion or even from the interstellar medium and become bright X--ray sources when they are in dense gas environments. Spectra of some ULX (Miller et al. \\cite{Mil03}) support the idea that they are IMBHs.  Substantial extended He II~$\\lambda4886$ emission has been confirmed inside the compact H II region \\#70 at the position of HoII X-1. Such strong He II~$\\lambda4886$ emission is produced by reprocessing of the X-ray source located inside the nebula. Pakull \\& Mirioni (\\cite{Pak01}) have concluded that the optical radiation is consistent with a quasi-isotrophic X-ray source and no significant beaming is required. In this case  the source is really so bright in X--rays (L$_{X}\\sim 10^{40}$\\,erg/s), which supports an IMBH interpretation ($iii$). The extended blue counterpart of HoII X-1 suggests that the IMBH is located in a young stellar complex or cluster.   The IMBH may disturb the velocity field (via the accretion disk wind or jets) of the H II region at its position, where we have found an increase of the velocity dispersion, and a radial velocity variation accross the He II~$\\lambda4686$ region. The amplitude of the radial velocity gradient accross the He II~$\\lambda4686$ region is $\\pm50$ km s$^{-1}$ on spatial scales of $\\pm2^{\\prime\\prime}$. This means that the ULX could not have only ionized the surrounding gas, but also has interacted with it dynamically.   The X-ray properties of HoII X-1 seem to agree with the IMBH interpretation ($iii$).  The soft component of the X-ray spectrum of HoII X-1 can be described by a MCD model with a inner disk temperature of $kT\\sim0.16$ keV. Such a cool accretion disk indicates intermediate-mass black holes with a black hole mass of $M_{BH}\\simeq10^{3}$M$_{\\sun}$ (Miller et al. \\cite{Mil03}). A further indication for a IMBH is the long-term variability of HoII X-1 (see MLH01 and Dewangan et al. \\cite{Dew04}).  The optical/X-ray properties of HoII X-1 seem to favour  the IMBH nature of the object.  High angular resolution integral field observation with HST or ground based near-infrared integral field spectroscopy using adaptive optics  at an 8-m class telescope are needed to prove the IMBH nature of HoII X-1, and would allow to put tighter constrains on the velocity structure of the gas and the stars inside the cluster around HoII X-1, and could enable us to measure the mass of the IMBH."
    },
    "0410/astro-ph0410034_arXiv.txt": {
        "abstract": "{ We present current results from our ongoing project to study the parsec-scale relativistic jet kinematics of sources in the U.S. Naval Observatory's Radio Reference Frame Image Database (RRFID).  The RRFID consists of snapshot observations using the VLBA plus up to 9 additional antennas at 8 and 2 GHz, and is intended to allow monitoring of these sources for variability or structural changes so they can be evaluated for continued suitability as radio reference frame objects.  The Image Database currently contains about 3000 images of 450 sources from 1994 to 2004, with some sources having images at 20 epochs or more.  We have now completed analysis of the 8 GHz images for all sources observed at 3 or more epochs from 1994 to 1998.  The completed analysis comprises 966 images of 87 sources, or an average of 11 epochs per source.  Apparent jet speeds have been measured for these sources, and the resulting speed distribution has been compared with results obtained by other large VLBI surveys. The measured apparent speed distribution agrees with those found by the 2~cm survey and Caltech-Jodrell Bank (CJ) survey; however, when a source-by-source comparison is done with the 2~cm survey results, significant disagreement is found in the apparent speed measurements for a number of sources. This disagreement can be traced in most cases to either an insufficient time baseline for the current RRFID results, or to apparent component mis-identification in the 2~cm survey results caused by insufficient time sampling.  These results emphasize the need for long time baselines and dense time sampling for multi-epoch monitoring of relativistic jets. }  \\titlerunning{Current Results from the RRFID Kinematic Survey} ",
        "introduction": " ",
        "conclusions": "Although about 50 epochs have been observed for the RRFID program over the last ten years, only about the first 20 epochs from 1994-1998 have been fully reduced. If the additional 30 epochs obtained from 1999-2004 are eventually added to those already reduced for this paper, then the RRFID Kinematic Survey will be by far the largest survey of jets (in terms of number of images) yet made, and will have the advantages of both a long time baseline (10 years), and dense time sampling ($\\approx 5$ epochs/year).  This survey also has the obvious advantage that the VLBI observations comprising it have already been made, so it remains only to reduce the data.  Meanwhile, work is continuing on the current data set, including searches for non-linear motions of jet components, and analysis of parsec to kiloparsec misalignment angles."
    },
    "0410/astro-ph0410202_arXiv.txt": {
        "abstract": "{In this paper we present deep VLT multi wavelength imaging observations of the field around the radio galaxy \\object{PKS~1138$-$262} aimed at detecting and studying a (potential) proto-cluster centered at this radio source.  \\rg\\ is a massive galaxy at $z = 2.16$, located in a dense environment as indicated by optical, X-ray and radio observations.  We had already found an over-density of \\lya\\ emitting galaxies in this field, consistent with a proto-cluster structure associated with the radio galaxy.  In addition, we find 40 candidate \\ha\\ emitters that have nominal rest frame equivalent width $> 25$ \\AA\\ within 1.8 Mpc and 2000 \\kms\\ of the radio galaxy. Furthermore, we find 44 objects with $I - K > 4.3$.  This number of extremely red objects (EROs) is about twice the number found in blank field ERO surveys, suggesting that some EROs in this field are part of the proto-cluster.  The density of \\ha\\ emitters and extremely red objects increases towards the radio galaxy, indicating a physical association.  From comparisons with other $K$ band, ERO, \\ha\\ and \\lya\\ surveys, we conclude that \\rg\\ is located in a density peak which will evolve into a cluster of galaxies. ",
        "introduction": "The search for clusters at high redshift has two main incentives: distant clusters can be used to constrain cosmological models and they provide a reservoir of high redshift galaxies, which can be used to study galaxy formation and evolution.  According to hierarchical clustering theories, clusters form by the gravitational amplification of primordial density fluctuations. In a low density universe, fluctuations cease to grow after a redshift $z \\sim (\\Omega_0^{-1} -1)$ \\citep{pee80}, resulting in a cluster population that evolves very slowly at low redshift. In an $\\Omega_0 = 1$ universe density fluctuations continue to grow even at the present epoch, implying that the cluster population would still be evolving rapidly \\citep{eke96}. The detection of even a single distant massive cluster, such as MS 1054$-$03 at $z = 0.83$, constrains the parameters ($\\Omega_0, \\sigma_8$) of cosmological models \\citep{bah98, don98}.  The study of galaxies in nearby and distant clusters provides strong constraints on their evolution and formation. It has been shown that clusters at high redshift can contain a different galaxy population mix than those at low redshift.  \\citet{but84} found that at $0.1 < z < 0.5$, compact clusters have significant numbers of blue galaxies, the fraction increasing with redshift, while at $z < 0.1$ cluster cores are essentially devoid of these. This effect can be explained if we assume a process in which spirals lose their gas and ability to form stars or are converted into early type galaxies in the course of their evolution. Massive ellipticals, however, dominating cluster cores up to $z \\sim 1$ seem to form a homogeneous population.  Even at $z = 1.27$, \\citet{dok01a} observe a scatter of the colour-magnitude relation of cluster galaxies similar to that in lower redshift clusters.  The tight colour-magnitude relation observed out to high redshifts can be explained by several galaxy formation scenarios.  \\citet{pro99} propose a passive luminosity evolution model, where galaxies form all their stars in a single burst at $z = 2$ after a monolithic collapse \\citep[e.g.][]{egg62}.  \\citet{dok01b} show that the observations are also consistent with a scenario in which early-type galaxies are continuously transformed from spiral galaxies, causing the progenitors of the youngest, low redshift early-type galaxies to drop out of the sample at higher redshifts (\\emph{progenitor bias}). Models by \\citet{dok01b} show that about half of the early-type galaxies were morphologically transformed at $z < 1$ and their progenitors may have had roughly constant star formation rates before transformation. These models indicate a mean luminosity weighted formation of stars in early-type galaxies of $z = 2.0$ (for $\\Omega_{\\rm M}$ = 0.3 and $\\Omega_\\Lambda$ = 0.7), consistent with the currently favoured hierarchical galaxy formation models predicting the merging of smaller galaxies at high redshift \\citep{kau96}.  A third explanation is given by metallicity differences in bright and faint ellipticals.  A model by \\citet{kau98} which includes hierarchical formation of ellipticals out of disc galaxies which have formed stars at modest rates and allows for the ejection of metals out of discs by supernova explosions predicts the establishment of a mass-metallicity relation among both late- and early-type galaxies.  In this model, large ellipticals are more metal-rich because they are formed from the mergers of larger discs.  The strongest observational constraints on these models come from the highest redshift data.  Both for the study of cluster and galaxy evolution a sample of clusters at high redshift is therefore desired.  \\medskip  \\noindent In recent years, much effort has been invested in the search for distant clusters, using both optical and X-ray observations. At $z > 0.5$, it becomes difficult to identify the projected two-dimensional over-density produced by cluster galaxies, because large numbers of foreground and background galaxies reduce the density contrast in the optical wavelength regime.  However, the $J-K$ colour of nearly all galaxies out to $z = 2$ is a simple function of redshift, because their near infrared light is dominated by evolved giant stars. The highest redshift cluster found to date \\citep[CIG J0848+4453 at $z = 1.27$]{sta97} has been discovered in a near infrared field survey as a high density region of objects with very red $J-K$ colours. Optical spectroscopy confirmed the redshifts of eight members and a 4.5$\\sigma$ ROSAT X-ray detection confirms the cluster's existence.  Although the detectability of hot cluster gas in X-rays is severely reduced by cosmological surface brightness dimming, \\citet{ros98} has found within the ROSAT Deep Cluster Survey a cluster at $z = 1.11$ \\citep[RDCS J0910+5422]{sta02} with 9 spectroscopically confirmed cluster members and a cluster at $z = 1.26$ \\citep[RX J0848.9+4452]{ros99} with 6 cluster members confirmed. Most of the confirmed galaxies have red colours, consistent with passively evolved ellipticals formed at high redshift ($z \\sim 5$).  The latter cluster is very close to CIG J0848+4453, with which it might form a superstructure and is possibly in the process of merging.  The recently started XMM Large Scale Structure Survey \\citep{ref02}, covering over 64 square degrees of sky, should be able to detect clusters with X-ray luminosities of $2 \\times 10^{44}$ erg s$^{-1}$ out to $z = 2$.  Despite the success of the near infrared and X-ray techniques, it is difficult to push these methods to find $z \\gg 1$ over-densities.  \\medskip  \\noindent A practical way to find clusters and groups of galaxies at high redshift is to study fields containing luminous radio galaxies. These can be observed up to the epoch of galaxy formation and efficiently selected by their steep spectrum in the radio regime \\citep{rot94,lac94,bre00}. The most distant radio galaxy found to date has a redshift of 5.2 \\citep{bre99b}. The host galaxies of powerful radio sources are amongst the most massive at any redshift \\citep{jar01,bre02} and are associated with $\\sim 10^9$ M$_\\odot$ BHs \\citep{lac01,mcl02}.  There has long been evidence that powerful radio galaxies at high redshift (HzRGs, $z > 2$) are located in the center of (forming) clusters of galaxies.  \\citet{yat89} and \\citet{hil91} find that the average environment of 70 powerful classical double radio sources at $0.15 < z < 0.82$ is that of an Abell 0 cluster, with some in environments as rich as Abell class 1.  At $z \\sim 1$ and higher, there is also evidence for galaxy over-densities associated with radio galaxies. \\citet{bes00} presents an analysis of the environments of 28 3CR radio galaxies at $0.6 < z < 1.8$. The density of $K$-band galaxies in these field, their angular cross correlation amplitude and near infrared colours correspond to the properties of low redshift Abell richness class 0 to 1 clusters. The author concludes that many, but not all, powerful radio galaxies at $z \\sim 1$ lie in cluster environments.  Furthermore, \\citet{nak01} have applied a photometric redshift technique based on five optical and near infrared images of the field of 3C324 at $z = 1.2$ and identified 35 objects as plausible cluster members.  The evidence extends even to $z = 3.8$, where observations in the field of radio galaxy 4C41.17 by \\citet{ivi00} tentatively reveal a number of luminous submm galaxies over-dense by an order of magnitude as compared to typical fields.  There is also evidence that the environments of HzRGs are dense in terms of ambient gas, as expected for the centers of (forming) clusters.  Radio continuum observations of $\\sim 70$ radio galaxies at $z \\sim 2$ \\citep{car97,pen00b} show that 20 -- 30\\% have large ($\\geq 1000$ rad m$^{-2}$) radio rotation measures (RMs). These RMs are most probably due to magnetized, ionized gas local to the radio sources and are comparable to the RMs of lower redshift radio galaxies which lie in the centers of dense, X-ray emitting cluster atmospheres \\citep{tay94}.     \\medskip  \\noindent Although convincing evidence for high density environments associated with radio galaxies at $z > 1$ has been demonstrated, the redshifts of possible cluster members have not been confirmed by spectroscopy and it is therefore impossible to provide a velocity dispersion for these structures.  In a program which is currently being carried out with the VLT, we are targeting the fields of luminous radio galaxies at $z > 2$ and observe these with the aim of detecting line emitting galaxies in the associated cluster.  At $z > 2$ the \\lya\\ line is redshifted into the optical wavelength region, where we can use narrow band filters ($\\sim$ 1\\%) to isolate its flux from the sky background. We have selected ten luminous radio galaxies at $2.2 < z < 5.2$, for which we will carry out both imaging and multi object spectroscopy. The survey is progressing very well and has already produced the discovery of the most distant structure of galaxies known \\citep[at $z = 4.1$,][]{ven02}.  We consider the \\lya\\ imaging and spectroscopy as a first step to the characterization of the cluster properties. The \\lya\\ emitters in a cluster form only a fraction of the galaxies present and might not be representative for some of the cluster or galaxy properties. The program will therefore be followed up by e.g.\\ broad band imaging in several colours. In this paper, we present new optical and near infrared observations of the field of \\rg\\ with the aim of uncovering populations of \\ha\\ emitting galaxies and EROs in the structure.  The radio galaxy \\rg\\ at a redshift of 2.16, was selected from a compendium of more than 150 $z > 2$ radio galaxies as the optimum object for beginning a high redshift cluster search. It combines most of the above mentioned cluster indications with a redshift suitable for both \\lya\\ and \\ha\\ imaging. The magnitude of \\rgs\\ is the brightest of all known radio galaxies close to $z = 2$. After correction for possible non stellar components, the $K$ band magnitude is 16.8, from which a stellar mass of $10^{12}$ M$_\\odot$ was inferred \\citep{pen97}. The radio galaxy possesses a giant ($\\sim$ 120 kpc) and luminous \\lya\\ nebula, with a wealth of structure: a bright region associated with the radio jet and filaments extending over $>$ 40 kpc \\citep{pen97,kur00b}. The optical counterpart of the radio galaxy is extremely clumpy and resolved into many components by the HST \\citep{pen98}. These clumps have properties similar to LBGs. The morphology of the system is consistent with hierarchical models of galaxy formation in which the LBG building blocks will merge into a single massive system, such as the massive galaxies observed at the centers of some rich clusters. The extremely distorted radio morphology \\citep{car97} is strong evidence that the jets have been deflected from their original direction by a dense and clumpy medium. The observed rotation measures of the radio emission (6200 rad m$^{-1}$, the largest in a sample of 70 HzRGs, see \\citealt{car97,pen00b}) and its steep gradient over the radio galaxy components also testify that the radio source is embedded in a dense magnetized medium. Additional evidence for a dense surrounding medium comes from \\emph{Chandra} X-ray observations, which reveal thermal emission from shocked gas \\citep{car02}. The pressure of this hot gas is adequate to confine the radio source.  Narrow band imaging of redshifted \\lya\\ emission of a 7\\arcmin$\\times$7\\arcmin\\ region around the radio galaxy \\citep[][Paper I]{kur00} and subsequent \\lya\\ spectroscopy \\citep[][Paper II]{pen00a} revealed 14 \\lya\\ emitting galaxies and one QSO. The galaxies have redshifts in the range 2.16$\\pm$0.02 with a velocity dispersion substantially smaller than expected for a random sample of galaxies selected by the narrow band filter. In addition, the \\emph{Chandra} X-ray observations of the field of \\rgs\\ have revealed at least five AGN at $z \\sim 2.16$ \\citep{pen02}.  On the basis of the evidence from the radio galaxy properties, the \\lya\\ halo and the galaxy over-density, we concluded that the structure of galaxies surrounding \\rg\\ is (the progenitor of) a cluster.  \\medskip  \\noindent The new observations of \\rgs\\ are reported in Sect.~\\ref{obs}. Detection and photometry of objects in the field of \\rgs\\ are presented in Sect.~\\ref{cat}. Subsequently, the selection from these objects of $K$ band galaxies, EROs, candidate \\ha\\ emitters and candidate \\lya\\ emitters is presented in Sect.~\\ref{can}. The properties of the EROs and candidates are analyzed in Sect.~\\ref{prop}.  A discussion of the implications of these results for the nature of the structure can be found in Sect.~\\ref{clus}, which is followed by a summary of the results and conclusions in Sect.~\\ref{sum}.  Throughout this article, we adopt a Hubble constant of $H_0 = 65$ km s$^{-1}$Mpc$^{-1}$ and a $\\Lambda$ dominated cosmology: $\\Omega_{\\rm M}$ = 0.3 and $\\Omega_\\Lambda$ = 0.7.  The over-densities of galaxies at high redshift, which have not yet reached virialization and/or a colour-magnitude relation with a red sequence, but will later form clusters, will be called \\emph{proto-clusters} here. ",
        "conclusions": "\\label{sum}  We have presented new near infrared and optical imaging of the field of radio galaxy \\rg\\ at $z = 2.16$. From these data and published optical imaging and spectroscopy, we find several galaxy populations associated with the radio galaxy.  $\\bullet$ In the 12.5 square arcminute field covered by the $K_s$ band images, we detect 550 objects in $K_s$. The number density of $K$ band sources is a about 1.5 times higher than blank field counts and increases towards the radio galaxy.  $\\bullet$ In the same area, we find 44 objects with $I - K > 4.3$. The surface density of EROs is about a factor two higher than the mean density found in blank fields. Because the density of EROs increases towards the radio galaxy, we conclude that some of the EROs are galaxies at $z \\sim 2.2$ and associated with the radio galaxy. Most of these are star forming galaxies reddened by dust, although up to five may be passive ellipticals.  $\\bullet$ Also in the same area, we find 40 candidate \\ha\\ emitters. The density of the brightest emitters is comparable to the density in other AGN fields and, again, increases towards the radio galaxy. The upper limit to the \\lya/\\ha\\ ratio of these sources is 1.1, but for most candidates the ratio is much smaller, meaning that these objects contain some dust or neutral hydrogen. Two \\ha\\ candidates are also identified as EROs. These must be dusty starbursts at $z \\sim 2.2$.  $\\bullet$ In the effective 43.6 square arcminute area of the optical images, we find 40 candidate \\lya\\ emitters. The surface density of these candidates does not show an increase towards the radio galaxy. The \\lya/\\ha\\ ratio of the 26 \\lya\\ candidates in the area covered by infrared imaging is in the range $2 - 7$. The sight lines through which we observe the \\lya\\ radiation are therefore relatively dust free.  $\\bullet$ The SFRs derived from rest frame UV flux of the sample of \\ha\\ and \\lya\\ candidates are comparable, but the galaxy masses derived from the flux around $\\sim 7000$\\AA\\ are on average larger for the \\ha\\ candidates than for the \\lya\\ candidates. From the spatial distribution of the two samples, we deduce that the \\ha\\ emitters are more settled into the potential well of the proto-cluster than the \\lya\\ emitters.     $\\bullet$ From the spatial distribution, number density and redshift distribution of objects detected in the field of \\rg, we conclude that the radio galaxy is located in a density peak which will evolve into a cluster of galaxies at the present day.  The bound mass at $z = 2.16$ is $10^{15}$ M$_\\odot$ at most.   \\smallskip The structure around \\rg\\ is one of the few proto-clusters known at $z > 2$ and deserves further study. X-ray observations of this field have been carried out by \\citet{car02} and an analysis of the X-ray point sources and AGN content of the proto-cluster is presented in \\citet{pen02}.  Infrared spectroscopy of a small sample of the candidate \\ha\\ emitters has been carried out and will be reported in Kurk et al.\\ (in preparation)."
    },
    "0410/astro-ph0410728_arXiv.txt": {
        "abstract": "We propose a possible unified model of short and long gamma-ray bursts (GRBs), X-ray rich GRBs, and X-ray flashes. The jet of a GRB is assumed to consist of multiple sub-jets or sub-shells (i.e., an inhomogeneous jet model). The multiplicity of the sub-jets along a line of sight $n_s$ is an important parameter. If $n_s$ is large ($\\gg 1$) the event looks like a long GRB, while if $n_s=1$, the event looks like a short GRB. Finally, when $n_s=0$, the event looks like an X-ray flash or an X-ray rich GRB. Furthermore, our model may also explain the bimodal distributions of $T_{90}$ duration of BATSE-GRBs. ",
        "introduction": "While the association with the supernova is almost established or strongly suggested for the long gamma-ray bursts (GRBs) \\citep{stanek2003,hjorth03}, the origins of short GRBs and X-ray flashes (XRFs) remain unclear. The observed event rate of  short GRBs is about a third of the long GRBs while the observed event rate of  XRFs is also about a third. Although there may be a possible bias effect to these statistics, these numbers are, in an astrophysical sense, the same or  comparable. If these  phenomena arise from essentially different origins, the similar number of events is just by chance. On the other hand, if they are related, the similar number of  events is natural and the ratio of the event rate tells us something about the geometry of the central engine. We propose a unified model in which the central engine of short GRBs, long GRBs and  XRFs is the same and the apparent differences come essentially from  different viewing angles. For details, see \\citet{yin04b} and \\citet{toma2004}. ",
        "conclusions": "Let $\\Delta\\theta_{\\rm sub}$, $\\vartheta_c$ and $\\bar{n}_s$ be the typical opening half-angle of the sub-jet, the core size of the whole jet and the mean multiplicity in the core. Then the total number of the sub-jets ($N_\\tot$) is estimated as $N_\\tot=\\bar{n}_s(\\vartheta_c/\\Delta\\theta_{\\rm sub})^2\\sim 10^3$, so that the total energy of each sub-jet is $\\sim 10^{48}$ erg. In our model, the event rate of long GRBs is in proportion to $\\vartheta_c^2$. Let $M$ be the number of sub-jets in the envelope of the core with a  multiplicity $n_s=1$. Then the event rate of short GRBs is in proportion to $M\\Delta\\theta_{\\rm sub}^2$, so that $M\\sim 10$ is enough to explain the event rate of short GRBs.   Since the core may be regarded as a uniform jet, our model for the XRFs is analogous to the off-axis uniform jet model \\citep{yin02,yin03b,yin04a}. However the afterglow could have a different behavior between the core-envelope sub-jet model and the uniform jet model. In the uniform jet model, the afterglows of XRFs should resemble the orphan afterglows that initially have a rising light curve \\citep{yin03a,g02}. An orphan afterglow may be actually observed in XRF 030723 \\citep{f04}, but the light curve may peak too early \\citep{z03}. The optical afterglow of XRF 020903 is not observed initially ($<0.9$ days) but may not be consistent with the orphan afterglow \\citep{so03}. These problems could be overcome by introducing a Gaussian tail with a high Lorentz factor around the uniform jet \\citep{z03} since the energy redistribution effects may bring the rising light curve to earlier times \\citep{z03,kg03}. The afterglow of a short GRB is difficult to predict since it could resemble both the orphan and normal afterglow depending on the sub-jet configuration within the envelope.   It is also found that our model can reproduce the bimodal distribution of $T_{90}$ duration of GRBs observed by BATSE \\citep{toma2004}. In our model, The duration of $n_s=1$ burst is determined by the angular spreading time of one sub-jet emission, while that of $n_s \\geq 2$ burst is determined by the time interval between the observed first pulse and the last one. These two different time scales naturally lead a division of the burst $T_{90}$ durations into the short and long ones. It has commonly been said that the observed bimodal distribution of $T_{90}$ durations of BATSE bursts shows the different origins of short and long GRBs. However, the bimodal distribution is also available as a natural consequence of our unified model of short and long GRBs. The clear prediction of our unified model is that short GRBs should be associated with energetic SNe. Indeed, one of short GRBs shows the possible association with a SN \\citep{germany00}. Even if the SNe are not identified with short GRBs due to some observational reasons we predict that the spatial distribution of short GRBs in host galaxies should be similar to that of the long GRBs. Another prediction is that short GRBs should have the same total kinetic energies as long GRBs, which might be confirmed by radio calorimetry.   Interestingly, our model has predicted short XRFs or short X-ray rich GRBs \\citep{yin04b}. They are observed when  isolated sub-jets are viewed slightly off-axis. The observed short XRF~040924 may be a kind of these bursts \\citep{fenimore04}."
    },
    "0410/gr-qc0410027.txt": {
        "abstract": "In this paper we establish dualities between inflationary, cyclic/ekpyrotic, and phantom cosmologies within the patch formalism approximating high-energy effects in scenarios with extra dimensions. The exact dualities relating the four-dimensional spectra are broken in favour of their braneworld counterparts; the dual solutions display new interesting features because of the modification of the effective Friedmann equation on the brane. We then address some qualitative issues about phantomlike cosmologies without phantom matter. ",
        "introduction": "The ultimate theory of everything, if any, is a long-living mirage that physicists and mathematicians have been pursuing for years in the attempt to solve many fundamental problems rooted in our modern view of the Universe. One of the open issues is how to reconcile general relativity and quantum physics, two separate branches that experiments and observations has widely accepted as meaningful descriptions of natural phenomena, at least each in its own range of influence. The marriage between the two would require deep modifications of both and, although great progress has been made in this direction thanks to string theory, a happy ending to the story is still missing. In particular, the  mostly successful big bang model of cosmological evolution, which manages to glue gravitation and microphysics together in a very nontrivial way, sits on the paradox of the initial singularity: the original point from which all came decrees the failure of general relativity as a self-contained framework, since the relevant cosmological quantities diverge by definition when going back to the first instant of the past. At present we know that quantum effects can resolve such a point into a finite speck and smooth out the worried infinities \\cite{strin}.  From a philosophical perspective, the big bang has raised many questions about the nature of time and its birth, leading to the (indeed not new) hypothesis that the Universe may experience a cyclic succession of expansions and contractions in which the big bang singularity is just a transitory phase (a bounce) in a wider process of evolution; see \\cite{old} for old attempts to implement this idea. At a semiclassical level, the structure of the perturbations generated through the bounce can be more complicated than the standard one in a monotonically expanding universe; for example, vector modes cannot be neglected during the contracting phase in contrast to their decaying behaviour in the post big bang phase \\cite{BB}. General phenomenology of cyclic models and bouncing cosmological perturbations have been studied, e.g., in \\cite{PP1,CDC,PPG,KST,EWST,AW}; the case of a bouncing closed universe has been investigated in \\cite{close}.  Moreover, bouncing flat cosmologies may require a violation of the null energy condition $\\rho+p\\geq 0$, where $\\rho$ and $p$ are the energy density and pressure of a perfect fluid describing the matter content of the early Universe \\cite{BM,PP1,AW,PP2,PPG}. The field associated to an equation of state $p=w\\rho$ with $w<-1$ is called ``phantom'' \\cite{pha1} and its nonconventional properties, including superacceleration ($\\ddot{a}/a>H^2$), can give rise to a new sort of singularity (the big rip) as well as to an explanation of current observations of dark energy. For a reference list on the subject, see \\cite{cal3} to which we add \\cite{ACL,phant}.  An interesting singularity-free setup, alternative to inflation and motivated by string theory, is the ekpyrotic scenario, which explains the large-scale small anisotropies of the cosmic microwave background via a collision between wrinkling branes \\cite{ekpy1,KOSST,ekpy2}. A general-relativistic treatment of ekpyrotic/cyclic scenarios predicts a scale-invariant scalar spectrum (with scalar index $n_s-1\\approx 0$) and a blue-tilted tensor spectrum $n_t\\approx 2$, while standard inflation generates almost scale-invariant spectra. In the latter case, this is a consequence of the slow-roll (SR) approximation, stating that both the parameter \\be \\label{epsilon} \\epsilon \\equiv -\\frac{d \\ln H}{d \\ln a}=-\\frac{\\dot{H}}{H^2}\\,, \\ee and its time derivative must be sufficiently small.\\footnote{In the following we will refer to $\\epsilon$ as the ``SR parameter'' even when the slow-roll approximation $\\epsilon \\ll 1$ is not applied.} Conversely, the cyclic model achieves scale invariance when $\\bar{\\epsilon}\\equiv \\epsilon_\\text{cyclic} \\gg 1$. Recently, two remarkable dualities were discovered in flat cosmology, one relating inflationary to ekpyrotic/cyclic spectra \\cite{GKST,dual,lid04} and the other connecting inflationary to phantom spectra  \\cite{ACL,phdu,lid04}. More precisely, given an inflationary model there exist both cyclic and phantom cosmologies with the same spectra and such that \\bs\\label{4Dduality}\\ba \\bar{\\epsilon} &=& 1/\\epsilon\\,,\\label{4Ddual}\\\\ \\hat{\\epsilon} &=& -\\epsilon\\,,\\label{4Dphdual}\\\\ \\bar{\\epsilon} &=& -1/\\hat{\\epsilon}\\,, \\ea\\es where $\\hat{\\epsilon}\\equiv \\epsilon_\\text{phantom}$. In four dimensions, this triality is exact for arbitrary (even varying) $\\epsilon$ \\cite{lid04}. Other dualities can be found in \\cite{altdu}.  The search of viable bouncing mechanisms has led to explore several possibilities that involve, for instance, varying couplings \\cite{BKM}, noncommutative geometry \\cite{MMMZ}, quantum gravity and cosmology \\cite{quang} (see also \\cite{sri04}). In particular, a Randall-Sundrum (RS) modification of the Friedmann equations has been considered \\cite{SS}, in which a phantom component may help to tear apart black holes during the bounce \\cite{BFK,PZ}.  In braneworld scenarios the visible universe is confined into a (3+1)-dimensional variety (a brane) embedded in a larger noncompact spacetime (the bulk). One of the first problems one has to deal with when constructing such models is how to stabilize the extra dimension. This can be achieved in a number of ways; in the RS example, Goldberger and Wise have provided a mechanism according to which a 5D massive scalar is put into the bulk with a potential of the same order of the brane tension $\\lambda$ \\cite{GW}. If the energy density $\\rho$ on the brane is smaller than the characteristic energy of the scalar potential, $\\rho/V \\sim \\rho/\\lambda \\ll 1$, then the radion is stabilized and one gets the standard Friedmann equation $H^2 \\propto \\rho$ on the brane. On the contrary, if the brane energy density is comparable with the stabilization potential, $\\rho/\\lambda \\gtrsim 1$, the bulk backreacts because it feels the presence of the brane matter, the minimum of the potential is shifted, and the well-known quadratic corrections to the Friedmann equation arise \\cite{quco}.  Here we shall follow the second alternative and consider nonstandard cosmological evolutions on the brane, extending the RS discussion to arbitrary scenarios we dubbed ``patch cosmologies'' \\cite{cal3}, with \\be \\label{FRW} H^2=\\beta_q^2 \\rho^q\\,. \\ee Here $H$ is the effective Hubble rate experienced by an observer on the brane, $\\beta_q$ is a constant we will assume to be positive, and the exponent $q$ is equal to 1 in the pure 4D (radion-stabilized) regime, $q=2$ in the high-energy limit of the RS braneworld, and $q=2/3$ in the high-energy limit of the Gauss-Bonnet (GB) scenario;\\footnote{Gauss-Bonnet cosmologies have been studied, e.g., in \\cite{cal3,gabo} and references therein.} if braneworld corrections are important in the early Universe, one can follow the cosmological evolution through each energy patch in a given time interval where the patch approximation is valid.  In this paper (Secs. \\ref{setup} and \\ref{patdua}) we will investigate the above-mentioned triality for general $q$ and show that Eq. (\\ref{4Dduality}) no longer realizes exact correspondences between cyclic, inflationary, and phantom patches. According to the new dualities we will establish, any expanding universe is mapped to either a contracting or phantom universe which no longer display exactly the same scalar perturbations; our results are in agreement with previous investigations \\cite{lid04}. Since braneworld spectra are broken under duality, and mapped into a quantitatively different contractinglike or phantomlike spectra, these transformations are not symmetries in the strict meaning of the word.  We will define the dual transformations both in a given patch by imposing $\\bar{q}=q$ and between different patches, which we will call cross dualities. In addition, it will turn out that the generalized version of the 4D contracting (phantom) mapping gives rise to a phantom (contracting) dual solution when flipping the sign of $q$.  At last, inspired by a modified version of the phantom duality we shall outline some proposals for (i) a new bouncing scenario, (ii) the generation of features in the power spectrum breaking scale invariance, and (iii) an alternative to standard inflation. A full detailed implementation of these topics in a rigorous theoretical framework is beyond the scope of this article and will be left for the future; rather, we shall provide some preliminary comments in Sec. \\ref{qbounce}. Conclusions are in Sec. \\ref{concl}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{concl}  In this paper we have generalized the four-dimensional triality between inflationary, cyclic, and phantom cosmologies to the case of a braneworld scenario with a modified effective Friedmann evolution $H^2=\\rho^q$. The exact, simple relations Eqs. (\\ref{4Dduality}) between the parameters $\\epsilon=\\dot{R}_H$ of models with an ordinary scalar field are broken and extended consequently. The self-dual solutions and the duals of power-law inflation have been provided for three scenarios (General Relativity, Randall-Sundrum, and Gauss-Bonnet braneworlds) in the presence of either a normal scalar field or a Born-Infeld tachyon. Finally, starting from a new version of the ``phantom'' duality, we have set some remarks on cosmologies with $q<0$.  The structure of the triality is deeply modified: The cosmologies dual to inflation either display singularities within finite time intervals or are nonsingular at the origin. This last feature is appealing as regards the construction of nonsingular bounces and would deserve further attention. Hopefully, the cross dualities we have found will help in the understanding of high-energy phantom and bouncing models of the early Universe.  Many topics remain to explore. For example, one should consider the contribution of the nonlocal physics of the bulk in order to set a truly consistent picture of braneworld cosmologies and dualities. In the typical inflationary context, Eq. (\\ref{FRW}) encodes the most part of the braneworld effective evolution; in fact, the simplest contribution of the projected Weyl tensor is $\\propto a^{-4}$ and is damped away during the accelerated expansion. However, the dark radiation term is no longer negligible in a contracting universe and should be taken into account.  In parallel, it would be interesting to explore two other directions. The first, most important issue should be to motivate $q<0$ scenarios within string theory, super or quantum gravity, since at this stage they are rather speculative. The second direction goes towards a study of the cosmological perturbations through the bounce, by further modeling the too simple step-function transitions we presented. A more concrete model would try to provide a smooth big crunch/big bang phase and allow a nonsharp transition in $q$.  In order to fully resolve the singular bounce we should rely on a description more general than classical gravity. To find reasonable solutions of the big bang singularity and embed a bouncing picture in a well-established (string-) theoretical framework will perhaps be one of the most promising lines of research in the following years, not only for the immediate cosmological implications (observability of preinflationary physics and comprehension of the high-energy early Universe) but also because it might lead to a better understanding of the still controversial but intriguing \\emph{landscape} of vacua \\cite{ban04}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0410/astro-ph0410124_arXiv.txt": {
        "abstract": "We report the first measurement of the X-ray spectrum of the z = 5.99 quasar SDSSp~J130608.26+035626.3 from a 120 ks observation by the \\axaf\\ ACIS-S instrument.  Between 0.5 and 7 keV, corresponding to 3.5--49 keV in the quasar rest frame, we find an energy index of 0.86 $\\pm$ 0.2, consistent with the typical indices found for radio quiet quasars at lower redshifts, and inconsistent with the index required to match the diffuse X-ray background.  We have a weak  indication of a redshifted Fe-K line.  In comparing the counting rate between an earlier, short observation and the longer observation reported here, we find evidence for source variability at the 99.9\\% confidence level. We note that other nearby X-ray sources would bias the measured $\\alpha_{\\rm ox}$ = 1.70 by $-$0.09 if the X-ray flux were determined from within a 60\\arcsec\\ extraction circle. Our results for the energy index and the $\\alpha_{\\rm ox}$ are consistent with no strong evolution in the active galactic nucleus emission mechanism with redshift out to z $\\approx$ 6, and therefore with the picture that massive black holes have already formed less than 1 Gyr after the big bang. ",
        "introduction": "The Sloan Digital Sky Survey has revealed the most distant quasars known \\citep{Fan01, Fan04}.  \\axaf\\ observations on 2002 January 29 easily detected three quasars at redshifts 5.82, 6.28, and 5.99 in brief, $\\approx$ 5--8 ks, ACIS-S observations \\citep{Schwartz02, Brandt02, Wilkes02, Bechtold02}. The latter, SDSSp J130608.26+035626.3 (hereafter, J1306), was followed up with a 120 ks \\axaf\\ observation, based on the claim by \\citet{Schwartz02} of an associated jet. \\citet{Ivanov02} subsequently identified the X-rays attributed to a jet with a faint galaxy, and the present observation clearly confirms the galaxy identification.  From the optical observations, J1306 is a broad line quasar, whose spectrum contains Ly$\\alpha$+NV emission, and various intervening absorption systems \\citep{Fan01, Becker01}. This is a general characteristic of the 12 quasars now known at z $\\approx$ 6 \\citep{Fan04}. From the optical absolute magnitude M$_{1450}$ = $-$26.93 \\citep{Fan01}, the central black hole mass must be at least 10$^{{\\rm 9}}$ M$_{\\sun}$, unless it is emitting at much greater than the Eddington luminosity.  We report here a study of the X-ray variability and spectrum of J1306. The X-ray spectrum of distant quasars is of interest to search for possible evolution of the active galactic nucleus (AGN) emission mechanism, which might give clues to the history of the formation and the fueling of the central massive black hole.  With  AGNs widely inferred to make up the bulk of the diffuse X-ray background, it remains to define the model of quasar evolution and of the relative numbers of absorbed and unabsorbed emitters that will reproduce the shape of the 2--10 keV diffuse X-ray background to the precision measured by \\citet{Marshall80}. It remains possible that a high fidelity spectral fit must also consider a contribution from quasars with concave curvature in their spectrum from 5 to 40 keV, as suggested by \\citet{Schwartz88}.  Previous measurements of radio-quiet quasars in the redshift range 2.64--4.35 gave photon spectral indices $\\Gamma \\approx$ 2.0 \\citep{Ferrero03,Brocksopp04,Grupe04}, with the exception of an index 1.4 reported by \\citet{Brocksopp04} for RX J122135.6+280613 at z=3.305. \\citet{Yuan98} report a mean index $\\Gamma \\approx \\rm{ 2.23\\pm 0.15}$ for 56 \\emph{ROSAT} quasars with z $>$ 2.0. The result reported here provides an example at z=5.99 with $\\Gamma$= 1.86. ",
        "conclusions": "In this Letter, we present the first \\axaf\\ X-ray spectrum of a z $\\sim$ 6 quasar.  The values for the energy index and the $\\alpha_{\\mathrm{ox}}$ derived from this \\axaf\\ spectrum of J1306 are consistent with the values found for lower redshift radio-quiet quasars \\citep{Brandt02}. Specifically,  our result for $\\alpha_{\\mathrm{ox}}$ is consistent with no strong evolution in $\\alpha_{\\mathrm{ox}}$ for optically selected radio-quiet quasars out to z $\\approx$ 6.  A preprint by \\citet{Farrah04} gives an \\emph{XMM} spectrum of SDSS J1030+0524, a quasar at z=6.30. They reach very similar conclusions as to the spectrum and luminosity, and implications for lack of evolution.  While \\emph{XMM} affords higher statistics, the 45\\arcsec ~extraction radius gives less certainty on the precise origin of the flux. Thus the two observations are somewhat complementary.  X-ray studies of high-redshift quasars reveal the conditions in the immediate vicinity of their supermassive black holes. Our spectrum of J1306 hints at the presence of an Fe-K emission line. Our poorly determined EW of 114 eV is consistent with the scatter in compilations by \\citet{Nandra97} and \\citet{Page04} which would extrapolate to predict a $\\approx$70 eV EW for a radio-quiet quasar of the 2$\\times$10$^{45}$ ergs s$^{-1}$ luminosity. Recent near-infrared spectroscopy of three other high redshift SDSS quasars (z $\\geq$ 5.7) with the \\emph{Hubble Space Telescope} \\citep{Freudling03} clearly detected Fe II  emission lines at 2200-2600 \\AA ~(rest frame) and the Keck-II observations of \\citet{Barth03} have detected Si IV, C IV, Mg II, and the 2900-3000 \\AA ~Fe line complex. Both results indicate that the strength of the line complex, as well as the Fe II/Mg II ratio, is comparable to that of lower redshift quasars. The results of near-infrared studies such as these as well as the results of X-ray studies such as the one presented here and elsewhere \\citep{Brandt02, Vignali03} are beginning to indicate that there apparently is no strong evolution in the broadband AGN emission mechanism with redshift out to z $\\approx$ 6.  This conclusion, if confirmed with a larger sample, implies that formation of the central black hole engines can be completed in less than 1 Gyr after the big bang, with the inference that a massive galaxy in which they are imbedded has also formed. Therefore, in order to address the question of which, if any, of the observed quasar properties evolve with time, we must probe deeper still to identify quasars at even higher redshifts.  This might best be done with the \\axaf\\ deep surveys, e.g., by identifying the objects with extreme X-ray-to-optical flux ratios found by \\citet{Koekemoer04}."
    },
    "0410/astro-ph0410312_arXiv.txt": {
        "abstract": "I discuss the empirical properties of diffuse X-ray emitting gas in the halos of both nearby starburst galaxies and normal spiral galaxies, based on high resolution X-ray spectral imaging with the Chandra X-ray Observatory. Diffuse thermal X-ray emission can provide us with unique observational probes of outflow and accretion processes occurring in star-forming galaxies, and their interaction with the inter-galactic medium. I consider both the  spatial distribution of the diffuse X-ray emission in and around edge-on starburst galaxies with superwinds (\\eg surface brightness profiles, distribution with respect to H$\\alpha$ and radio emission), and its spectral properties (\\eg thermal or non-thermal nature, abundance ratios, temperatures and soft and hard X-ray luminosities). These results are discussed in the context of current theoretical models of supernova-driven superwinds, and compared to the more limited data on extra-planar hot gas around edge-on normal galaxies. ",
        "introduction": "\\label{sec:dks:intro}  Galaxy formation and evolution, in particular that of star-forming galaxies, can not understood in terms of isolated galaxies evolving independently of their environment. Even in the field galaxy population mass accretion, mass outflow and galaxy/galaxy interactions and mergers play a fundamental role. Merging of stellar systems is important, but it not the only significant process, and accretion from the inter-galactic medium (IGM) and gas and metal loss to the IGM can and probably does occur via purely gaseous processes even outside dense galaxy group and clusters \\citep[see \\eg][]{toft02,tremonti04}. To observe and quantify these processes in action we must observe the disk/halo interfaces, and halos (the galaxy/IGM interface) of local star-forming galaxies.  The diffuse thermal X-ray emission from gas with temperatures between $10^{6}$ and $10^{8}$ K is a particularly important probe of the conditions within the halos of star-forming galaxies.  \\begin{itemize} \\item In the \\citet{chevclegg} model for galactic superwinds driven by multiple core-collapse supernovae (SN) and massive star stellar winds, the merged, metal-enriched, SN and wind ejecta has a temperature within the starburst region of $T\\sim  10^{8} \\epsilon/\\beta K$, and thus would be a source of \\emph{faint} hard thermal X-ray emission ($E>2$ keV). The fraction of SN mechanical energy thermalized is $\\epsilon$ ($\\la 1$), while $\\beta$ ($\\ga 1$) is the ratio of total mass added to the hot gas compared to that from SN and stellar wind ejecta. \\item The warm neutral and ionized gas in superwinds is observed to have velocities typically in the range $200$ -- 1000 km s$^{-1}$ \\citep{ham90,heckman2000}. This material is believed to be embedded within or at the boundary of a hot, metal-enriched, wind fluid that has even higher velocities \\citep[\\eg][]{ss2000}. Various processes will thus create soft thermal X-ray emission ($E<2$ keV), such as strong shocks ($T\\sim 3.5\\times10^{6} \\, [v_{\\rm shock}/500 {\\rm km\\,s}^{-1}]^{2}$) or thermal conduction at interfaces between the warm gas and the hot wind fluid. \\item The depth of Milky-Way-like galaxy gravitational potential wells corresponds to a Virial temperature of $T_{\\rm vir} \\sim 2\\times10^{6} \\, (v_{\\rm circ}/230 {\\rm km\\,s}^{-1})^{2}$ K. If gas inflowing from the IGM passes through an accretion shock, it will be heated to soft X-ray emitting temperatures. \\end{itemize} In non-active galaxies this emission is purely collisionally excited, and thus diffuse X-ray emission on galactic scales is only created by mechanical energy return from SNe or stellar winds, or accretion into deep gravitational potential wells. The complicated processes of photo-ionization and radiative feedback from massive stars, so important at IR, optical and UV wavelengths, are negligible at X-ray wavelengths.  The diffuse thermal X-ray emission from star-forming galaxies is a mixture of continuum (bremsstrahlung and recombination) and line emission processes. For soft thermal plasmas ($6 \\la \\log T (K) \\la 7$) line emission dominates the net emission for abundances $Z \\ga 0.2 Z_{\\odot}$. As newly synthesized metals are believed to initially enter the hot phases of the ISM, X-ray observations will be vital in capturing galactic chemical evolution and IGM enrichment at work. Unfortunately the CCD-spectrometers on Chandra and XMM-Newton have too low a spectral resolution to resolve the strong line emission complexes from O, Fe, Ne, Mg and Si from the continuum over the energy range 0.5 -- 2.0 keV, which complicates absolute element abundance measurements for such soft metal-bearing plasmas. Nevertheless, some useful spectral elemental diagnostics are possible with existing instrumentation, as I shall discuss.  Much of this contribution is based on our recent Chandra-based survey of 7 edge-on starbursts and 3 edge-on normal spirals \\citep{strickland04a,strickland04b}, along with a preview of more recent work on diffuse hard X-ray emission from M82. Our findings on the starburst systems and their superwinds are presented in \\S~\\ref{sec:dks:starbursts:hard} and \\ref{sec:dks:starbursts:soft} Their properties are compared empirically to the diffuse X-ray properties of the normal spirals in \\S~\\ref{sec:dks:both}. Conclusions drawn from this work, and some further questions raised by it, are summarized in \\S~\\ref{sec:dks:conclusions} ",
        "conclusions": "\\label{sec:dks:conclusions}  Chandra observations of starburst galaxies have answered some of the pressing questions regarding X-ray emitting gas in superwinds. The very hot plasma predicted by the \\citet{chevclegg} model has probably been detected in M82 at expected flux levels in the 6.7 keV Fe line, along with non-thermal continuum emission. Soft X-ray emission in winds is of low volume filling factor, as is associated with some form of wind fluid interaction with cooler denser ambient disk and halo gas. Strongly mass-loaded wind models can be ruled out.  The extra-planar X-ray luminosity of the approximately edge-on star-forming disk galaxies in our sample is directly proportional to the host galaxy's SF rate. The vertical extent of the diffuse X-ray emission is proportional to the host galaxy's size. We definitely need more X-ray observations of edge-on \\emph{normal} star-forming disk galaxies, especially if we are to constrain accretion models. What we can say with the existing observations is that the properties of the diffuse soft X-ray-emitting plasma in normal galaxies appears very similar to that in starbursts, scaled proportionally to lower SF rates. The existing non-detections of extra-planar hot gas do not prove that there is no hot halo gas.  With the aim of being provocative, I will close by asking some questions inspired by these results. How similar can the physical conditions in a normal galactic fountain be to a superwind? Can warm and hot gas filling factors be low in the halo of NGC 891, despite the apparent smoothness of the emission? Are the majority of non-detections of extra-planar warm ionized gas, hot ionized gas, radio emission, and dust significant? In other words, is there no extra-planar plasma, or is it just too faint to detect?"
    },
    "0410/astro-ph0410638_arXiv.txt": {
        "abstract": "Here we discuss the properties of star cluster winds in the supercritical, catastrophic cooling regime. We demonstrate that catastrophic cooling inhibits superwinds and after a rapid phase of accumulation of the ejected material within the star-forming volume a new stationary isothermal regime,  supported by the ionizing radiation from the central cluster, is established. The expected appearance of this core/halo supernebula in the visible line regime and possible late evolutionary tracks for super-star cluster winds,  in the absence of ionizing radiation, are thoroughly discussed. ",
        "introduction": "Within the volume occupied by a star cluster (SC), the energy injected by stellar winds and supernova explosions is fully thermalized via random interactions. This generates the large central overpressure that continuously accelerates the ejected gas and eventually blows it out of the star cluster volume to compose a superwind. In the adiabatic solution of Chevalier \\& Clegg (1985; hereafter referred to as CC85; see also Canto, et al. 2000 and Raga et al. 2001) temperature and density present  almost homogeneous values within the central volume, whereas the expansion velocity grows almost linearly from 0 km s$^{-1}$ at the center, to the sound speed ($c$) at the cluster radius $r = R_{SC}$. There is then a rapid evolution as matter streams away from the star cluster and the wind parameters (velocity, density, temperature and pressure) soon approach their asymptotic values: $V_w \\to $ $V_{A\\infty} = (2 L_{SC}/{\\dot M}_{SC})^{1/2} \\sim 2c(R_{SC})$, $\\rho_w \\sim $ $r^{-2}$, $T_w \\sim $ $r^{-4/3}$ and thus $P_w \\sim r^{-10/3}$, where $L_{SC}$ and ${\\dot M}_{SC}$ are energy and mass deposition rates.  Recent results on the outflows expected from SCs have led us to realize that the adiabatic steady wind solution proposed by CC85 and followers, becomes inapplicable in the case of massive and concentrated clusters. Radiative cooling strongly modifies first the temperature distribution predicted for adiabatic stationary winds ($T_w \\sim r^{-4/3}$) bringing suddenly and within  a small radius, the temperature down to 10$^4$ K, restricting then the X-ray emissivity of the winds to a volume much smaller than previously thought (see Silich et al. 2003 and 2004; hereafter referred to as Papers I and II). Also, as shown in Paper II, for more energetic clusters,  strong radiative cooling  promotes the sudden leakage of thermal energy right within the star cluster volume itself, and for the cases in which the radiative losses there exceed 30$\\%$ of the stellar energy deposition rate, when cooling becomes catastrophic, then the stationary superwind solution is totally inhibited. In this latter case, the rapid drop in temperature within the SC volume, leads to a sudden drop in central pressure and particularly to a sudden drop in sound speed. This inhibits the fast acceleration predicted in the adiabatic solution as required by  the flow to reach its adiabatic terminal speed ($v_{A\\infty} \\sim 10^3$ km s$^{-1}$). The sudden drop in sound speed upsets also  the balance, demanded by the stationary solution, between the stellar mass input rate and the rate at which matter can flow away from the cluster, ie: $\\dot M_{SC} = 4 \\pi R_{SC}^2 \\rho_{SC}(R_{SC}) c(R_{SC})$, and this inevitably leads to mass accumulation.   Here we show how the metallicity attained by  the ejected matter from a coeval cluster, strongly affects the limits found in Paper II bringing the energy input rate further down to much smaller values and thus affecting even lower mass clusters. We also show how after the thermal stationary superwinds are inhibited, and after an inevitably short phase of matter accumulation within the SC volume, a second stationary solution can be found. This is only possible while the stellar UV flux ionizes the deposited matter to compose a dense supernebula with an isothermal ($T \\sim 10^4$ K), slow ($v_{\\infty} \\sim 50 $ km s$^{-1}$), stationary wind (see section 2). Finally, as the evolution continues and the stars producing the UV photon output evolve into supernovae, a third quasi-adiabatic stationary stage could become possible after a second phase of matter accumulation within the star cluster radius. This could in principle evolve into a massive, cold, neutral outflow, however, its maximum velocity is well below the escape speed and thus the ejected matter, unable to escape the gravitational potential of the star cluster, will sooner or latter recolapse to cause a second major burst of star formation. Our conclusions and the observational properties of the various evolutionary phases are given in section 3. ",
        "conclusions": ""
    },
    "0410/astro-ph0410228.txt": {
        "abstract": "{ We report on a complete CCD imaging survey of 226 elliptical galaxies in the North-East quadrant of the Virgo cluster, representative of the properties of giant and dwarf elliptical galaxies in this cluster. We fit their radial light profiles with the Sersic $r^{1/n}$ model of light distribution. We confirm the result of Graham \\& Guzman (2003) that the apparent dichotomy between E and dE galaxies in the luminosity--$<\\mu>_e$ plane no longer appears when other structural parameters are considered and can be entirely attributed to the onset of \"core\" galaxies at $B_T \\sim -20.5$ mag. When \"core\" galaxies are not considered, E and dE form a unique family with $n$ linearly increasing with the luminosity. \\\\ For 90 galaxies we analyze the B-I color indices, both in the nuclear and in the outer regions. Both indices are bluer toward fainter luminosities. We find also that the outer color gradients do not show any significant correlation with the luminosity. The scatter in all color indicators increases significantly toward lower luminosities, e.g. galaxies fainter than $B_T \\sim -15$ have a B-I spread $>$0.5 mag. ",
        "introduction": "The widespread belief that dwarf (dE) and giant (E) elliptical galaxies form two distinct families brought unanimous consent that separate mechanisms are responsible for their formation and evolution (see Ferguson \\& Binggeli 1994 for a review). Alleged evidence for a structural dE/E dichotomy includes: i) the light profiles of dEs follow exponential laws (Binggeli et al. 1984), as opposed to their giant counterparts that follow the $r^{1/4}$ law (de Vaucouleurs 1948); ii) scaling relations of the structural parameters $r_e$, the effective radius, and $<\\mu>_e$, the mean surface brightness within $r_e$, with the total luminosity indicate that dEs of increasing luminosity have greater surface brightness while giant Es show the reverse trend (Binggeli \\& Cameron 1991, Binggeli et al. 1984). Moreover they occupy distinct loci in the $r_e$ vs. $<\\mu>_e$ plane: giants have dimmer surface brightness with increasing radius (Kormendy \\& Djorgowski 1989), while dEs do not.\\\\ Hydrodynamical simulations, including mass loss from stars and gas heating and cooling, show that gravitational instability of primeval density fluctuations (with or without dark matter) results in the formation of elliptical galaxies with brighter $<\\mu>_e$ at high luminosity, as observed in the dE regime, without showing any dichotomy, i.e. the bright objects result as scaled-up versions of the less luminous systems (Athanassoula 1993). \\\\ Various mechanisms involving merging have been proposed for the formation of E galaxies that are consistent with the alleged dE/E dichotomy. Dwarf Es in clusters might derive from the harassment of dwarf spiral galaxies (Moore et al. 1996; Moore, Lake, \\& Katz 1998; Mao \\& Mo 1998; Mayer et al. 2001), while giant E galaxies might arise from the merger of two spiral galaxies or from the multiple merging of dEs (Kauffmann et al. 1993).\\\\ There is however some observational evidence arguing for a continuity rather than a dichotomy between dEs and Es: i) the U-V (Caldwell 1983) and B-H (Scodeggio et al. 2002) colors and the metallicity are found to smoothly increase with luminosity; ii) the Near-IR light concentration index (Scodeggio et al. 2002) and the central surface brightness (Caldwell \\& Bothun 1987; Jerjen \\& Binggeli 1997) increase monotonously with the absolute magnitude; iii) no clear dichotomy can be assessed between giants and dwarfs in the Near-IR $<\\mu>_e-R_e$ and fundamental plane relations (Zibetti et al. 2002). \\begin{figure*} \\centering \\includegraphics[width=16.0cm]{0571f01.ps} \\caption{The N-E region of the Virgo cluster including galaxies with $m_p\\leq19.0$. The dashed line represents the boundary of the VCC catalog. The rectangles include the studied WFH and WFV areas, covering approximately 26 $deg^2$. The location of M87 is indicated.} \\label{sample} \\end{figure*} Whether there is kinematic evidence for the dE/E dichotomy remains a controversial issue: the preliminary conclusion of Geha et al. (2001) that dE and E are both not rotation-supported systems has been subsequently questioned by the same authors (Geha et al. 2003) who concluded, in agreement with Pedraz et al. (2002) and van Zee et al. (2004), that dE galaxies can have significant rotation.\\\\ Moreover the very existence of the aforementioned dE/E dichotomy has been recently questioned by Graham \\& Guzman (2003; GG03 hereafter). These authors argued that the dichotomy is apparent, as it derives from the use of exponential and de Vaucouleurs laws to decompose the light profiles of dE and E galaxies respectively, combined with the use of $<\\mu>_e$ as representative of the surface brightness. When $\\mu_o$ is analyzed, as obtained from the inner extrapolation of Sersic $r^{1/n}$ laws (Sersic 1968), no dichotomy remains.\\\\ Stimulated by the conclusions of GG03 we undertook this project aimed at re-assessing the relations involving the size, the surface brightness and the luminosity of dE-E galaxies, all ingredients entering in the determination of the fundamental plane. The present analysis is carried out using a sample which by selection should be representative of the properties of dE-E galaxies in one cluster.\\\\ The Virgo cluster, owing to its relatively small distance from us (17 Mpc), represents  an appropriate testbed for such a study, although its significant line-of-sight depth (Gavazzi et al. 1999; Solanes et al. 2002) and sub-clustering makes it not ideal. In spite of their low surface brightness, dwarf galaxies as faint as 19 mag are within reach of mid-size telescopes and seeing limitations are less severe when studying galaxies at the distance of Virgo than objects 5 times further away such as the Coma galaxies that require the superior resolution of HST. Since their luminosity function was determined ( e.g. Sandage et al. 1985), early systematic studies of dwarf galaxies in Virgo were based on photographic material (Binggeli \\& Cameron 1991, 1993,  Young \\& Currie (1998)). Modern photometry obtained with panoramic detectors exists today for many Virgo galaxies (e.g. Caon et al. 1993; Gavazzi et al. 2003; Barazza et al. 2003) and many have been observed with the HST (e.g. Faber et al. 1997; Cote et al. 2004). A large stretch of the Virgo cluster was mosaiced with the Isaac Newton Telescope in what is known as the \"Virgo Wide Field Survey\" (Sabatini et al. 2003). These data are used in the present analysis to re-assess the issue of the dE/E dichotomy.\\\\ The analyzed sample is described in some detail in Section 2.  Sections 3 and 4 describe the method used to derive the light profiles and their fitting with Sersic models. The B-I color analysis and the analysis of the various structural B-band parameters are carried out in Section 5 and the conclusions of our work are briefly summarized in Section 6. ",
        "conclusions": "We carried out a CCD survey of the North-East quadrant of the Virgo cluster containing $\\sim$ 30\\% of the galaxies in this cluster. We analyzed images for 226 and 90 galaxies respectively in the B- and I-band, representative of the properties of giant and dwarf elliptical galaxies in this cluster. \\\\ We fit the galaxies' radial light profiles with the Sersic $r^{1/n}$ model of light distribution obtaining the structural parameters $n$, $<\\mu>_e$, $r_e$, $B_T$, $C_{31}$ and $\\mu_o$. We find that the Sersic model provides an adequate representation of the observed light profiles of dE and E galaxies spanning 9 magnitudes.\\\\ We confirm the result of Graham \\& Guzman (2003) that the apparent dichotomy between E and dE galaxies in the luminosity--$<\\mu>_e$ plane is due to the onset of ``core\" formation in elliptical galaxies at $B_T \\sim -20.5$ mag. The segregation among dE and E no longer appears when other structural parameters are considered.\\\\ For 90 galaxies in the WFV we analyze the B-I color indices, both in the nuclear and in the outer regions. Both indices are bluer toward fainter luminosities, however the scatter in these relations increases significantly toward lower luminosities. Moreover we find that the outer color gradients do not show any significant correlation with the luminosity."
    },
    "0410/astro-ph0410581_arXiv.txt": {
        "abstract": "The HI distribution  at the epoch of reionization (EOR) is largely determined by the sizes and distribution of the ionized regions. In the scenario where the ionized regions have comoving sizes of the order of a few Mpc, the large scale statistical properties of the HI distribution are dominated by the Poisson noise of the discrete ionized regions,  and it is highly non-Gaussian. We investigate the possibility of probing reionization by studying these non-Gaussian features using  future  radio interferometric observations of redshifted 21 cm HI radiation. We develop a formalism relating correlations between the visibilities measured at  three different  baselines and frequencies to the bispectrum of HI fluctuations. For visibilities at the same frequency, this signal is found to be of the same order as the two visibility  correlation  which probes the HI power spectrum.  For visibilities at different frequencies, we find that the  correlations decay within a frequency difference of $\\sim 1 \\, {\\rm MHz}$. This  implies  that it is, in principle, straightforward to extract this  HI signal from various contaminants which are believed to have  a continuum spectra and are  expected to be correlated even at large frequency separations. ",
        "introduction": "There has recently been a lot of interest in understanding exactly how and when the universe was reionized. There now are significant observational constraints mainly from three different kinds of observations.  The observation of quasars at redshift $ z \\sim 6$ which show strong HI absorption  \\citep{becker} indicates that at least $1\\%$  of the total hydrogen mass   at $ z \\sim 6 $  is neutral \\citep{fan}, and the  neutral mass fraction decreases rapidly   at lower redshifts. This is a strong indication that the epoch of reionization ended at $ z\\sim 6 $ . Observations of the CMBR polarization, generated through Thomson scattering of  CMBR photons by free electrons along the line of sight, indicates that the reionization began at a redshift $z > 14$. On the other hand, the observed anisotropies of the CMBR indicate that the total optical  depth of the Thomson scattering is not extremely high, suggesting that reionization could not have started at redshift much higher than about  30 (\\citealt{kogut}; \\citealt{spergel}). A third constraint comes from determinations of the IGM temperature from observations of the ${\\rm Ly} \\alpha$  forest in the $z$ range $2$ to $4$ which indicates a complex reionization history with there possibly being an order unity change in the neutral hydrogen fraction at $z \\le 10$ (\\citealt{theuns}; \\citealt{hui}).   Mapping the HI distribution at high redshifts using radio observations  of the redshifted 21 cm radiation (\\citealt{madau};  \\citealt{scott};  \\citealt{kumar} ) holds the possibility of probing  the transition from a largely neutral to a largely ionized universe    at a level of detail surpassing any other techniques.     \\citet{zald} (hereafter ZFH) have developed a statistical technique based on  the angular power spectrum, on lines similar to the analysis of CMBR anisotropies, for analysing the HI signal from the epoch of reionization (EOR) in radio interferometric observations. Extracting the HI signal from various Galactic and extra Galactic contaminants (eg.  \\citealt{cooray};  \\citealt{dimat}; \\citealt{gnedin2}; \\citealt{oh};  \\citealt{dimat1}; \\citealt{shaver}; ) is one of the most important challenges. Most of the known contaminants are expected to have continuum spectra, and ZFH show that it should in principle be possible to extract the HI signal using the fact that, unlike the contaminants,  it   will be uncorrelated at two slightly different frequencies. The frequency dependence of the angular power spectrum of the HI signal and foregrounds has recently  been  analysed in detail by \\citet{santos}.  An alternative statistical technique for analysing the HI signal is to study  the correlations between the complex visibilities measured at different baselines and frequencies in radio-interferometric observations. This has been developed in the context of observing HI from $z < 6$  (\\citealt{bharad2}; \\citealt{bharad3}; \\citealt{bharad4}) and later generalized to the  EOR signal in \\citealt{bharad6} (hereafter BA). The possibility of using visibility correlations to quantify the EOR signal  has also been proposed by \\citet{morales}   who further discuss how  the different frequency signatures of the contaminants and the HI signal can be used to distinguish between the two. Recently  \\citep{morales1} has   addressed the issue of the power spectrum sensitivity   of the EOR HI signal.  Various investigations (eg. ZFH, BA) show that the power spectrum of HI fluctuations at EOR has contributions from  mainly  two distinct effects, the clustering of the hydrogen which, on large scales,  is assumed to follow the dark matter distribution and the fluctuations arising from the presence of discrete  regions of ionized gas surrounding the sources responsible for reionizing the universe. The details of the reionization process  are not very well understood (eg. \\citealt{barkana}), and the shape, size and distribution of these ionized regions is one of the very important issues which will be probed by  21 cm HI observations. There has recently been progress in analytically modeling the growth of the ionized regions  \\citep{furlanetto1} (hereafter FZH)  based on the findings of simulations   (\\citealt{ciardi};  \\citealt{sokasiana}; \\citealt{sokasianb}; \\citealt{nusser}; \\citealt{benson}; \\citealt{gnedin}) which show that  there will not be a large number of small HII regions around individual ionizing sources, rather there will be a few large ionized regions centered on places  where the ionized sources are clustered. The size of these ionized regions are expected to be  around a few Mpc (comoving) or possibly larger at EOR. In such a  scenarios, on scales larger than the size of the individual ionized  regions,  the HI signal will be dominated by the Poisson noise arising from the discrete nature of the ionized regions (eg. ZFH, BA, FZH, \\citealt{furlanetto2}). Further, the HI signal is expected to be highly non-Gaussian .    Nearly all of the work on quantifying    the  EOR HI signal expected in radio interferometric observations has focused on the two point statistics namely the angular power spectrum and  the correlations between pairs of visibilities. Both these quantities are actually equivalent  and they  basically probe the power spectrum of HI fluctuations at EOR. The power spectrum completely quantifies a Gaussian random field, but the higher order statistics would contain independent information if the HI fluctuations at EOR were not a Gaussian random field.  FZH  have used the pixel distribution function, a one-point statistics, to quantify non-Gaussian features in the HI distribution. \\citet{He} have studied the non-Gaussian features that arise in the HI distribution in  the log-normal model.  In this paper we address the issue of quantifying the non-Gaussian features of  the HI signal expected in radio interferometric observations. In particular, we focus on the correlation between three visibilities. This is expected to be zero if the signal were a Gaussian random field, and deviations from zero are a clear signature of the non-Gaussian properties of the HI distribution. Here we derive the relation between the three visibility correlation and the   bispectrum of the HI  fluctuations. The bispectrum quantifies correlations between three Fourier modes, and this is non-zero only when there are phase correlations between different modes. The three visibility correlation, as we show, is comparable to the correlations between two visibilities and this leads us to speculate that  this  will play an important role in detecting the HI signal. Further, the higher order  correlations  contains independent information, and observing these would throw independent light on  the topology   and morphology of the   HI distribution  at EOR.  Finally, an outline of the paper. In Section 2. we present the formalism relating the three visibility correlation to the HI bispectrum.  In Section 3. we introduce a simple model for the HI distribution at reionization and calculate its bispectrum. In Section 4. we present results  for the  three visibility correlation expected from HI at reionization and discuss some consequences.  It may also be noted that we use the values $(\\Omega_{m0},\\Omega_{\\lambda0},  \\Omega_b h^2,h)=(0.3,0.7,0.02,0.7)$ for the cosmological parameters thoughout. ",
        "conclusions": "In this section we present results for the visibility correlations expected from HI during  the epoch of reionization. Our aim being to illustrate the non-Gaussian nature of the expected signal and its dependence on the   ionized regions, we show results centered on only at a single frequency namely $125 \\, {\\rm MHz}$.  This corresponds to a  redshift $z=10.4$ when  the mean neutral fraction is   $\\bar{x}_{\\HI}=0.67$  ({\\it ie.} $f_V =0.33$). We choose this particular frequency as $\\bar{x}_{\\HI}$ is quite  close to $0.5$ where the HI signal is expected to be maximum, simultaneously  ensuring that the volume fraction where the HI density  predicted by our model becomes  negative is small ($\\sim 5 \\%$). Further, the HI signal is expected to be dominated by discrete ionized regions  and hence we  anticipate significant non-Gaussian features.  We have used eqs. (\\ref{eq:a9}) and (\\ref{eq:s2}) to calculate the expected correlations between two and three visibilities respectively. For this it is necessary to specify   a value for $\\theta_0$, the   beam size of the   individual antennas in the array. Further, it may be noted that $\\theta_0  \\approx 0.6 \\times \\theta_{\\rm   FWHM}$. The value of $\\theta_0$ will depend on the physical dimensions of the antennas and the wavelength of observation.  For the GMRT $\\theta_0=1^{\\circ}$ at $325 \\, {\\rm MHz}$. We  scale this  using $\\theta_0 \\propto  \\lambda$ to  obtain $\\theta_0=2.6^{\\circ}$    at $125 \\, {\\rm MHz}$  which we use here. The HI signal predicted here is  for observations using the GMRT,  and they  can be directly compared to  those in BA.     Both $\\C_2$ and $\\C_3$  scale as $\\theta_0^2$, and it is straightforward to scale the results presented here to make  visibility correlation predictions for  other radio telescopes.  The comoving radius   of the ionized spheres  $R$  is a free parameter in  our model. Investigations on the growth of the ionized spheres (FZH) show that these will be at least a few Mpc in radius (possibly larger) at the redshift of interest. We have considered three possible values $R=(1, \\,  3 , \\, 5) \\, {h^{-1}   \\, Mpc}$ for which  the respective values of $\\bar{n}_{\\HI}$ are $(78, \\, 2.9 , \\, 0.63 ) \\, \\times 10^{-3} \\, h^{3} {\\rm Mpc}^{-3}$.  For ease of graphical presentation, we have restricted our analysis of $\\C_3$ to equilateral triangles for which the size of the baseline $U$ completely specifies the triangle, and we have $\\C_3(U,\\Delta \\nu)$. Further, we first consider the correlations at the same frequency {\\it ie. } $\\Delta \\nu=0$   Figures (\\ref{fig:1}) and ({\\ref{fig:2}) show the results  $[\\C_2(U)]^{1/2}$ and $[-\\C_3(U)]^{1/3}$ respectively.  \\begin{figure} \\includegraphics[width=84mm]{c2n.eps} \\caption{This shows the visibility correlation $\\C_2(U,\\Delta \\nu)^{1/2}$ as   a function of $U$ for $\\Delta \\nu=0$, for different values of $R$, the comoving radius of the ionized  spheres. These predictions are for observations centered at $125 \\, {\\rm MHz}$. } \\label{fig:1} \\end{figure} \\begin{figure} \\includegraphics[width=84mm]{c3n.eps} \\caption{This shows the visibility correlation $[-\\C_3(U,\\Delta \\nu)]^{1/3}$ as   a function of $U$ for $\\Delta \\nu=0$, for different values of $R$, the comoving radius of the ionized  spheres. These predictions are for observations centered at $125 \\, {\\rm MHz}$. } \\label{fig:2} \\end{figure} We find that at small $U$,   $[\\C_2(U)]^{1/2}$ is  more or less constant with a value  of the   order of $\\sim 0.2 \\  {\\rm mJy}$ for $R=3 \\, h^{-1}\\, {\\rm Mpc}$. The   signal is proportional to $R^{3/2}$ and its magnitude  increases as the ionized spheres become larger.  Each baseline $U$ can be associated with a comoving length-scale  $\\rn/(2 \\pi U)$ at the comoving distance  where the HI radiation originated.   The signal from the ionized spheres  is constant across the   baselines for which $\\rn/(2 \\pi U)$ is larger than the size of the  spheres, and the signal falls at baselines for which $U > \\rn/(2 \\pi R)$.  Each    baselines resolves out features larger than $\\rn/(2 \\pi U)$, and the  presence of discrete ionized regions make very little contribution to  the signal at the large baselines.    Comparing the results for $[\\C_2(U)]^{1/2}$ presented here with those  presented in BA  which also includes the effects of gravitational clustering, we note that the gravitational clustering signal is also of the order $\\sim 0.1 \\,  {\\rm    mJy}$ at small $U$.  The gravitational clustering signal  also falls with increasing $U$,  and the combined signal would depend critically on the size of the bubbles.  For example, the signal from discrete ionized sources would dominate over the gravitational clustering signal at baselines $U > 500$ if the  ionized spheres had  comoving radius $R= 5 \\, h^{-1} \\, {\\rm Mpc}$,  whereas the gravitational clustering signal would possibly dominate throughout for  $R= 1 \\, h^{-1} \\, {\\rm Mpc}$.  Turning our attention next to $\\C_3$ (Figure \\ref{fig:2}), the first point to note is that this is negative. The shape of $\\C_3$ as a function of $U$ is very similar to that of $\\C_2$, and its magnitude is  around  $[-\\C_3(U)]^{1/3} \\, \\sim \\, 0.06 \\, {\\rm mJy}$ at $R=3 \\, h^{-1} \\, {\\rm Mpc}$, which is around three times smaller than $[\\C_2(U)]^{1/2}$. At small $U$,   $[-\\C_3(U)]^{1/3}$ is  more or less  constant. Although our results are restricted to equilateral triangles, we expect the correlations to be nearly constant for triangles of all shapes provided all the baselines satisfy $U < \\rn/(2 \\pi R)$.  The   signal is proportional to $R^2$ and its magnitude  increases a little faster than that of $[\\C_2(U)]^{1/2}$  as $R$ is increased. We expect the dark matter density fluctuations at $z > 10$  to be well in the  linear regime  on comoving  length-scales $\\sim 10 \\, h^{-1}\\,  {\\rm Mpc}$ or larger, and the contribution to $\\C_3$ from  non-linear gravitational  clustering  is expected to be very small on these scales.   It may be noted that the contribution to $\\C_3$ from linear gravitational clustering is exactly zero in the standard scenario where the initial density fluctuations are a Gaussian random field. Further, the  comoving length-scale      $10 \\, h^{-1}\\,  {\\rm   Mpc}$ corresponds to the baseline  $U \\sim 100$ and we expect the contribution from individual ionized spheres considered here to be the dominant signal at these baselines. \\begin{figure} \\includegraphics[width=84mm]{c2ndelnu.eps} \\caption{This shows the visibility correlation $\\C_2(U,\\Delta \\nu)^{1/2}$ as   a function of $\\Delta \\nu$ for the three different values of  $U$ shown in the figure.  The  comoving radius of the ionized  spheres is assumed to be $R=3 \\, h^{-1} {\\rm Mpc}$ and the predictions are for observations centered at $125 \\, {\\rm MHz}$. } \\label{fig:3} \\end{figure} \\begin{figure} \\includegraphics[width=84mm]{c3ndelnu.eps} \\caption{This shows the visibility correlation $[-\\C_3(U,\\Delta \\nu)]^{1/3}$ as   a function of $\\Delta \\nu$ for the three different values of  $U$ shown in the figure.  The  comoving radius of the ionized  spheres is assumed to be $R=3 \\, h^{-1} {\\rm Mpc}$ and the predictions are for observations centered at $125 \\, {\\rm MHz}$. } \\label{fig:4} \\end{figure}  We next consider the correlations between the visibilities at different frequencies. Again, for ease of graphical presentation we have restricted our analysis of $\\C_3(U,\\Delta \\nu_1, \\Delta \\nu_2)$ to equilateral triangles with the added restriction that $\\Delta \\nu_1 = \\Delta \\nu_2 = \\Delta \\nu$, so we have $\\C_3(U,\\Delta \\nu)$. We have shown results only for $R=3 \\, h^{-1} \\, {\\rm Mpc}$, but a similar behaviour is expected for other values also. We find that both $\\C_2(U,\\Delta \\nu)$ and $\\C_3(U,\\Delta \\nu)$ fall rapidly, in nearly the same fashion independent of $U$, and are very close to zero by $\\Delta \\nu \\approx  0.5 \\, {\\rm MHz}$.  One of the main challenges in observing cosmological HI is to  extract it  from  various contaminants which are expected to swamp this signal. The contaminants include  Galactic synchrotron emission, free-free  emission from ionizing halos  \\citep{oh}, faint radio loud quasars \\citep{dimat1} and synchrotron emission from low redshift galaxy clusters \\citep{dimat}. Fortunately, all of these foregrounds have smooth continuum spectra and we expect their contribution to the visibilities to be correlated over large $\\Delta \\nu$, whereas the HI contribution is  uncorrelated beyond   $1 \\, {\\rm MHz}$ or less. It is, in principle, straightforward to fit the visibility correlations $\\C_2$ and $\\C_3$  at large $\\Delta \\nu$ and remove any slowly  varying component thereby separating the contaminants from the HI signal. We also use this opportunity to note that this is a major advantage of using visibility correlations as compared to the  angular power spectrum  which exhibits substantial correlations even at two frequencies separated by $\\sim 10 \\, {\\rm MHz}$ \\citep{santos}.  An important fact which emerges from our analysis is that the HI signal  in the correlation between three visibilities is of the same order as  the correlation between two  visibilities, the former being around three times smaller. This is a generic feature of the EOR HI signal, valid   if the ionized regions are bubbles of the size $R=1 h^{-1} \\, {\\rm Mpc}$ or larger. This signal arises from the Poisson noise of the discrete ionized regions, and it is  enhanced if the  size of the bubbles is increased. The fact that there is a substantial $\\C_3$ tells us  that there  are large  phase correlations between the visibilities. This is a consequence of  the  fact that there are only a few coherent features (the ionized regions) which dominate the whole HI signal.   Investigations on the growth of ionized regions (FZH) show that there will be a spread in the  sizes of  the ionized regions at any given epoch. This will smoothen some of the sharp features seen in Figures (\\ref{fig:1}) and (\\ref{fig:2}). The ringing seen in the these figures is an artifact of there being only a single value of $R$ and we do not expect this feature to be there if we have a spread in $R$. Further, the gravitational clustering signal not shown here may also dominate at large $U$. Despite all these limitations, we can still expect substantial correlations between three visibilities in a a more realistic analysis,  this being a robust signature  of the fact that reionization occurs through a few, large ($R \\sim $ a few Mpc) bubbles of ionized gas and the HI signal is dominated by  Poisson noise.   We next briefly discuss  the noise levels and the integration times required to observe the HI signal, particularly addressing  the question  whether $\\C_3$ can be detected with integration times comparable to those needed for  $\\C_2$. We consider an array of $N$ antennas, the observations lasting a time duration $t$, with frequency channels of width $\\delta \\nu$ spanning a total bandwidth $B$.  It should be noted that the effect of a finite channel width $\\delta \\nu$ has not been included in our calculation which assumes infinite frequency resolution. This effect can be easily included by convolving our results for the visibility correlation with the frequency response function of a single channel. Preferably, $\\delta \\nu$ should be much smaller than the frequency separation at which the visibility correlation become  uncorrelated.  We use $S$ to denote the frequency separation within which the visibilities are correlated,  and beyond which they become uncorrelated.  We use $\\N_2$ and $\\N_3$ to denote the rms. noise in $\\C_2$ and $\\C_3$ respectively. It is well known that  $\\N_2 = \\left( \\frac{2 k_B T_{SYS}}{A_{ef}}   \\right)^2 \\frac{1}{\\delta \\nu \\, t}$ \\citep{thomp}, and we have $\\N_3\\sim \\left( \\frac{2 k_B T_{SYS}}{A_{ef}}   \\right)^{3} \\frac{1}{(\\delta \\nu \\, t)^{3/2}}$ assuming that we have Gaussian random noise, where $T_{SYS}$ is the system temperature and $A_{ef}$ is the effective area of a single antenna. The noise contributions  will be reduced by a factor $1/\\sqrt{N_o}$ if we  combine $N_o$ independent samples of the  visibility correlation.  A possible observational strategy for  a preliminary detection of the HI signal would be to combine the visibility correlations at all baselines and frequency separations where there is a reasonable amount of signal. This gives $N_o=[N(N-1)/2] \\,  (B/\\delta \\nu) \\, (S/\\delta \\nu)$ for the two visibility correlation and $N_o=[N(N-1)(N-2)/6] \\,  (B/\\delta \\nu) \\, (S/\\delta \\nu)^2$ for the three visibility correlations.  It should be noted that we have used the fact that the $\\C_3$ is  non-zero only for the baselines between triplets of antennas. Combining all of this we have $[\\N_2]^{1/2} \\sim  \\left( \\frac{2 k_B T_{SYS}}{A_{ef}}   \\right) \\left[ \\frac{2}{N (N-1) B S} \\right]^{1/4} \\frac{1}{t^{1/2}}$ and $[\\N_3]^{1/3} \\sim  \\left( \\frac{2 k_B   T_{SYS}}{A_{ef}}   \\right) \\left[ \\frac{6}{N (N-1)(N-2) B S^2 } \\right]^{1/6} \\frac{1}{t^{1/2}}$.  The ratio $[\\N_3]^{1/3}/[\\N_2]^{1/2} \\sim [N (N-1) B/(N-2)^2 S]^{1/12}$ has a very weak dependence on $N$, $B$ and $S$ for a reasonable choice of values, and is of order unity.  We thus see that, for a given integration time,  we will achieve comparable noise levels in both the two and three visibility correlations.  Estimates of the integration time to detect $\\C_2$ (or equivalently the angular power spectrum) (BA, ZFH) indicate this to be around a few hundred  hours. We find that it should be possible to also detect $\\C_3$ in a comparable integration time."
    },
    "0410/astro-ph0410062_arXiv.txt": {
        "abstract": "{ We present a numerical  and analytical study of the thermal fragmentation of a turbulent  flow of interstellar hydrogen.  We first present  the different  dynamical  processes and  the  large range  of spatial (and  temporal) scales that need to  be adequately represented in numerical simulations.  Next, we present bidimensionnal simulations of turbulent converging flows which induce  the dynamical condensation of the  warm neutral phase into  the cold phase.  We  then analyse the cold structures and the fraction  of unstable gas in each simulation, paying  particular  attention  to  the  influence  of  the  degree  of turbulence.  When the  flow is very turbulent a  large fraction of the gas remains in the thermally unstable domain.  This unstable gas forms a  filamentary  network.   We  show  that the  fraction  of  thermally unstable gas  is strongly correlated  with the level of  turbulence of the  flow.  We  then develop  a semi-analytical  model to  explain the origin of this  unstable gas.  This simple model  is able to reproduce quantitatively  the   fraction  of   unstable  gas  observed   in  the simulations and  its correlation  with turbulence. Finally,  we stress the fact that even when the flow is very turbulent and in spite of the fact that a large fraction of the gas is maintained dynamically in the thermally unstable  domain, the classical picture of  a 2-phase medium with stiff thermal fronts and  local pressure equilibrium turns out to be   still  relevant  in   the  vicinity   of  the   cold  structures. ",
        "introduction": "The  dynamics  of the  interstellar  gas  is  of great  importance  to understand the formation of structures such as molecular clouds and it is therefore determinant in the star formation process.  In this respect, the neutral atomic  hydrogen (HI) is crucial since it represents more than 50$\\%$ of the  ISM and since the molecular clouds form through  the  condensation of HI  gas.   Because of the numerical stiffness of the    problem, previous numerical models  attempting  to simulate the ISM at a scale of about 1kpc and  to form molecular clouds self-consistently  have  not  considered heating and  cooling functions that  allow   thermal bistability  between  100   and   8000  K  (e.g V\\'azquez-Semadeni  et    al.    1995, 1996,   Korpi    et al.   1999, Ballesteros-Paredes et al.  1999) or have a numerical resolution which is not appropriate  to  adequately describe  this physics down  to the scale of  the CNM  structures (Gazol  et al.  2001).  Therefore  it is currently  unclear and indeed  almost unexplored  to  what extent  the physics of HI may  or  may not have  a  significant influence  on  the formation and the evolution of molecular clouds.  From the  theoretical point   of  view (Field  et  al. 1969,  McKee \\& Ostriker 1977, Wolfire et al. 1995)  as well as from the observational one (Low et al. 1984, Boulanger  \\& P\\'erault 1988, Kulkarni \\& Heiles 1987, Troland \\& Heiles 2003), it is now well established that HI is a thermally bistable  medium   which at thermal  equilibrium   and for a thermal pressure close  to about (in the  vicinity of the sun)  4000 K cm$^{-3}$, can   be in two different   thermodynamical states, namely a warm and diffuse phase (WNM) and a cold and dense one (CNM) roughly in pressure equilibrium.  The linear stability analysis  (Field 1965), the  quasi-static thermal front propagation  (Zel'dovich  \\&  Pikel'ner 1969, Penston   \\& Brown 1970)  and  more generally   the non-linear  development of  a  single structure (see Meerson 1996 for a  review and S\\'anchez-Salcedo et al. 2002 for a recent study) have been under investigation for a long time and are reasonably well understood.  However, it is only recently that the behaviour of  a thermally bistable  flow in the fully dynamical or turbulent regime has been investigated.   \\subsection{Previous work} In  the    context  of solar   chromosphere    and corona, Dahlburg et al. (1987)  and Karpen et  al. (1988) perform  2D simulations of a gas that  may undergo thermal   instability.  They have considered  random velocity perturbations   in  an initially thermally   unstable gas and study the development   of  cold  structures.  They  note  that  large amplitude  velocity   perturbations can    prevent  the formation   of condensed structures.  Kovalenko \\& Shchekinov (1999) and Hennebelle \\& P\\'erault (1999) have simultaneously investigated the possibility that a converging flow may dynamically trigger the thermal  transition  from the  WNM  (thermally stable) phase into the CNM phase. They perform 1D simulations and show that if the perturbation lasts long enough (more  than a cooling time) and is strong enough (velocity must be comparable to the sound speed), the  thermal  transition is   possible,  i.e part   of the  WNM  phase condenses  in CNM  leaving  a cold   structures   embedded in  a  warm surrounding medium.  Their underlying idea is that an external forcing like   bubble expansion  or any   phenomena   generating systematic or turbulent  motions  may promote the  formation  of cold structures.  A similar picture has been investigated by Koyama \\& Inutsuka (2000) who simulate  a shock propagating   in   HI and include H$_2$   formation. Hennebelle  \\& P\\'erault (2000) have  also  considered the  case of  a magnetic   field,  important in   the context  of   the  ISM, which is initially oblique with  respect to the  velocity  field. They  show that whatever the  value of the magnetic  field the thermal condensation is still possible provided the initial  angle between ${\\bf B}$ and ${\\bf V}$  is  small   enough.   The smallest angle  below     which thermal condensation is always possible, is about  20$^\\circ$. Due to magnetic tension the  condensation is more  difficult for intermediate magnetic intensity  than for a stronger field.  This paradigm  has been  further investigated  by Koyama  \\& Inutsuka (2002) who  simulate  in 2D a  shock propagating  in HI. They show that several structures of CNM form close  to the shock interface and find that the velocity dispersion of  the cold structures is about 5-6 km/s, i.e a fraction of the sound speed of the WNM.  Gazol et al.  (2001) have performed 2D  simulations of HI at a scale  of about 1 kpc and a numerical resolution of  about 5 pc with a turbulent forcing that mimics star  formation.  They found an important fraction ($\\simeq 50 \\%$) of thermally unstable gas in their simulation.  Kritsuk \\& Norman (2002a, 2002b) have performed 3D simulations of  HI with a  thermal forcing that mimics  the random fluctuations of the heating rate derived  by Parravano et al.  (2002). Their aim is to show  that  interstellar  turbulence   may be   generated by   thermal instability.  Finally, Piontek \\& Ostriker (2004) have recently performed 2D simulations with the aim of studying the development of the magneto-rotational instability as well as the thermal instability in the magnetised warm atomic interstellar gas.  \\subsection{Outline of the paper} In this  paper we further  investigate the dynamical properties  of HI gas by the mean of 2D numerical simulations.  We explore the behaviour of  HI  focusing  on  the  formation of  cold  structures  during  the collision  between  two streams  (converging  flow)  of  WNM.  We  put special attention to the structure morphology, their internal velocity dispersion and to  the fraction of gas in the cold  and warm phase and in an intermediate thermally unstable state.  Our aim is to understand how those features depend on the turbulence of the flow.   In the second section of  the paper we present the equations, describe the  thermal  processes, the  numerical  scheme  and  the initial  and boundary conditions.  We also discuss the drastic numerical resolution which is  needed in order  to properly describe  the HI flow.   In the third section we present our  numerical results for a large scale flow which is weakly or very  turbulent and present statistical analysis of the simulations.   In the fourth section we  emphasize the correlation between  turbulence and the  fraction of  thermally unstable  gas.  We then  develop  a  semi-analytical  model to  understand  the  physical mechanisms  which  are  responsible  for this  phenomena.   The  fifth section summarises our results and concludes the paper. ",
        "conclusions": "We perform 2-d numerical simulations with the aim of understanding the dynamics  of the thermally bistable  interstellar atomic hydrogen. The size of our  1000$^2$ simulation box is 20  pc, leading to a numerical resolution of 0.02 pc. Such resolution is  needed in order to properly describe the cold HI   structures.  In order  to  mimic a large  scale converging  flow that triggers the transition  of  the warm phase into the cold one, warm gas is injected  continuously from 2 opposite sides of  the  box. The gas  can  freely escape at  the  2 other boundaries. Various amplitudes of turbulent fluctuations have been applied.  When the flow is weakly turbulent, a layer of compressed WNM forms and quickly fragments into structures of cold gas.   When the flow is very turbulent,  it is less organised.  Nevertheless the thermally bistable behaviour  is not  erased  and  indeed  locally  very  similar to  the classical picture of the 2-phase medium.   In particular, the 2 phases are connected  through sharp thermal fronts and   are locally in rough pressure equilibrium.  In order to  characterise the CNM  structures  found in  the numerical simulation, we apply a simple threshold algorithm to identify them and give  the pdf of  some of their  intrinsic  properties as mean density, temperature, velocity dispersion,  size   and aspect ratio, $r$.    Typical values for  these parameters  are   respectively: $<n> \\simeq   50$ cm$^{-3}$, T$\\simeq$80 K, $\\delta v \\simeq$ 0.5 km/s, l$\\simeq$ 0.1 pc and r$\\simeq$0.5.  A large fraction  of thermally unstable gas  which increases with  the amplitude of the turbulent forcing,  is found. This thermally unstable gas tends to be organised in filamentary structures.  A semi-analytical model for  the thermal and  dynamical evolution of a fluid particle is presented. It is based on the Lagrangian description of a  fluid element.  This  model predicts  that substantial straining motions may  efficiently stabilise a piece   of thermally unstable gas allowing it to survive a  longer period of time.   We then verify than both locally and globally the  thermally unstable gas in the numerical simulations  is  indeed  very  strongly  correlated  with  substantial straining motions."
    },
    "0410/astro-ph0410548_arXiv.txt": {
        "abstract": "We study the propagation of polarization light through the magnetosphere of neutron stars.  At intermediate frequencies (the optical through the infrared), both the birefringence induced by the plasma and by quantumelectrodynamics influence the observation polarization of radiation from the surface of the neutron star. Because these two processes compete in this regime, we find that polarization observations can constrain the properties of the neutron-star magnetosphere, specifically the total charge density.  We calculate both the phase-resolved and the phase-averaged polarization signatures induced by magnetospheric birefringence. ",
        "introduction": "The propagation of polarized radio waves away from neutron stars (NSs) has been well understood for many years \\citep{Barn86}.  At this energy scale, it is reasonable to neglect the coupling of the light to vacuum-based QED effects.  Lai and Ho (2003ab)\\nocite{Lai:2003,Lai:2003b} have discussed the mechanisms by which polarized X-rays are affected by the atmospheres of NSs. X-ray photons couple the greatest in the atmosphere as the plasma densities and field strengths are sufficiently high.  Lai and Ho included vacuum-based effects in their analysis.  Similar effects exist at IR/optical wavelengths when propagation of the waves through the magnetosphere (possessing a lower-density plasma and lower strength magnetic field) is considered.  This paper presents a model and results that show that there is a wavelength-dependent polarization signature in the IR/optical wavelength as a result of interactions between radiation and the magnetosphere.  For most radio pulsars emission from the magnetosphere itself rather than the surface is thought to dominate the optical and near infrared emission.  However, several radio pulsars exhibit optical radiation from the surface ({\\em e.g.} PSR~B0656+14, PSR~B0950+08 and Geminga) \\citep[\\protect{\\em e.g.}][]{2004A&A...417.1017Z}.  Radio pulsars are not the only members of the neutron-star menagerie.  Tbe optical and near-infrared emission from the radio quiet neutron stars such as RX~J1856.5-3754 and RX~J0720.4-3125\\citep[\\protect{\\em e.g.}][]{2002ApJ...564..981P,2003A&A...408..323M} most likely comes from the surface, and possibly the emission from anomalous x-ray pulsars emerges from a surface layer heated by magnetospheric currents \\citep{2004astro.ph..8538T}.  With these objects in mind let us proceed.  To make progress we will make several simplifying assumptions. First, we assume that the charge carriers both positive and negative have the mass of an electron and that the plasma is cold.  A pair plasma is necessarily hot. Second, we assume that the Goldreich-Julian model is correct.  The Goldreich-Julian picture describes the equilibrium charge density for a rotating neutron star with its magnetic pole aligned (or anti-aligned) with the spin axis.  Here we will assume that the two axes are not necessarily aligned (especially in \\S~\\ref{sec:phase-integrations}).  Finally, the Goldreich-Julian equilibrium even in its realm of applicability is unstable \\citep[see][]{2004AdSpR..33..542M}.  In spite of these deficiencies, the Goldreich-Julian model provides a well-understood starting point to demonstrate that observations of polarization of the surface radiation from neutron stars may provide a window to diagnose the properties of the plasma and to encourage further research to probe more realistic models. ",
        "conclusions": "Of particular attention is the wavelength at which the resonant feature occurs - the near infrared.  As this energy, it would be possible to measure the polarization of the thermal radiation from a NS using a large ground-based telescope or Hubble's NICMOS instrument. This instrument is only sensitive to wavelengths less than 2.1~$\\mu$m, so it is possible that a large portion of the behaviour could not be probed.  If the total plasma density is significantly larger than the Goldreich-Julian plasma density, the feature moves into near infrared or optical energies, and more observing options become available.  This letter presents a model for the evolution of polarized light through the magnetosphere of highly magnetized neutron stars. We found that the polarization has a frequency-dependent signature, where the total polarization decreases significantly in the optical or near IR, dependent on the plasma density near the neutron star and the structure of the magnetic field. Observations of neutron stars in this frequency region may provide useful probes of the properties of neutron-star magnetospheres and elucidate the processes responsible for pulsar emission."
    },
    "0410/gr-qc0410025_arXiv.txt": {
        "abstract": "We have studied the evolution of the massless scalar field propagating in time-dependent charged Vaidya black hole background. A generalized tortoise coordinate transformation were used to study the evolution of the massless scalar field. It is shown that, for the slowest damped quasinormal modes, the approximate formulae in stationary Reissner-Nordstr\\\"{o}m black hole turn out to be a reasonable prescription, showing that results from quasinormal mode analysis are rather robust. ",
        "introduction": "It is well known that the surrounding geometry of a black hole will experience damped oscillations under perturbations. These oscillations are called \"quasinormal modes\" (QNM), which is believed to be a characteristic \"sound\" of black holes, and would lead to the direct identification of the black hole existence through gravitational wave observation, to be realized in the near future. In the past few decades, a great deal of effort has been devoted to the study of QNMs of black holes in asymptotically flat spacetimes (for comprehensive reviews see [1,2] and references therein). Considering the case when the black hole is immersed in an expanding universe, the QNMs of black holes in de Sitter space have also been investigated [3-5]. Motivated by the recent discovery of the Anti-de Sitter/Conformal Field Theory (AdS/CFT) correspondence, many authors have performed the study of QNMs in AdS spaces [6-18]. The study of the QNMs plays an important role in astrophysics, black hole physics and also string theory.  All these previous works on QNMs have so far been restricted to time-independent black hole backgrounds. It should be realized that, for a realistic model, the black hole parameters change with time. A black hole gaining or losing mass via absorption (merging) or evaporation is a good example. The more intriguing investigation of the black hole QNM calls for a systematic analysis of time-dependent spacetimes. Recently the late time tails under the influence of a time-dependent scattering potential has been explored in [19], where the tail structure was found to be modified due to the temporal dependence of the potential. The motivation of our work is to explore the modification to the QNM in time-dependent spacetimes. Instead of plotting an effective time-dependent scattering potential by hand as done in [19], we will introduce the time-dependent potential in a natural way by considering dynamical black holes, with black hole parameters changing with time due to absorption and evaporation processes. We will study the temporal evolution of massless scalar field perturbation. Instead of employing Kruskal-like coordinate in our first attempt before [20], here we start our discussion directly from the Vaidya metric.  The outline of this paper is as follows. In Sec. II we first go over the conventional treatment for the study of the wave propagation in stationary Reissner-Nordstr\\\"{o}m (RN) black hole. In Sec. III, we will present the numerical way for the study of the QNMs of charged Vaidya black hole. We will present our numerical results in Sec. IV. Sec. V contains a brief summary and a discussion. ",
        "conclusions": "We have studied the evolution of the massless scalar field propagating in a time-dependent charged Vaidya black hole background. The two generalized tortoise coordinate transformations were used to study the evolution of the massless scalar field. In the first case we used both the $r_ + $ and  $r_ - $, which is appropriated for the case of large charge in numerical calculation. In the other case we only used the $r_ + $, which is appropriated for the case of small charge. The wave equation with a time-dependent scattering potential is derived. In our numerical study, we have found the modification of the QNMs due to the temporal dependence of the black hole spacetimes. In the absorption process, when the black hole becomes bigger, both the imaginary frequency $\\left| {\\omega _I } \\right|$ and the real frequency $\\omega _R $ decrease with time. However, in the evaporating process, when the black hole loses mass, both $\\left| {\\omega _I } \\right|$ and $\\omega _R $ increase with the increase of time. The study shows that, for the slowest damped QNMs, the approximate formulas in stationary RN black hole is a good description for the time-dependent charged Vaidya black hole, which can be expressed as (\\ref{eq29}).  In general, we conclude that the time dependent processes are well described by an adiabatic approximation with a delay in the definition of the mass function, borrowing the results at each moment from the problem of a black hole of a time-independent masses, instantly equal to the map of the equivalent Vaidya black hole at a retarded time. This implies a huge simplification for the problem of determination of signals from time-dependent processes near black hole. Finally we notice that the period corresponding to QNM is much smaller than the time of flight of a relativistic object in the black hole, actually by several orders of magnitude. We thus conclude that the value we used for the $\\lambda $ parameter is large enough to include realistic cases."
    },
    "0410/astro-ph0410254_arXiv.txt": {
        "abstract": "We have performed three-dimensional, non-LTE (non-Local Thermodynamic Equilibrium) radiative transfer calculations for $^{12}$CO and $^{13}$CO lines, applying them to our high-resolution hydrodynamic models of the `torus' around a supermassive black hole in an active galactic nucleus.  The hydrodynamic simulations reveal inhomogeneous and turbulent gas structure on a sub-pc scale in a circum-nuclear starburst region.  Thick disks interlaced with filaments, clumps and holes are naturally formed due to the interplay among energy feedback from supernovae, self-gravity of the gas, galactic rotation, and radiative cooling.  The intensity maps of the molecular lines for the circum-nuclear disks show a clumpy structure reflecting the intrinsic inhomogeneity and turbulent motion of the gas disk. The fine structure of the `torus' could be resolved in the nearby active galaxies using the Atacama Large Millimeter/submillimeter Array. We also found that the CO-to-H$_2$ conversion factor (X-factor) is not uniformly distributed in the central 100 pc region. The X-factor derived for $^{12}$CO ($J=$1-0) intensity depends strongly on the intensity, whereas the X for $^{12}$CO ($J=$3-2) is nearly constant over two orders of magnitude of the intensity. The suggested conversion factor for the molecular gas mass is $X_{{\\rm CO} (J=3-2)} \\sim 0.27 \\times 10^{20}$ cm$^{-2}$ (K km s$^{-1}$)$^{-1}$.  The line-ratio of high-J transitions of CO is not uniformly distributed in the nuclear disk, and the apparent ratio depends on the beam sizes. `Observed' \\twco (J=4-3)/\\twco (J=2-1) can differ from the intrinsic ratio by maximally 20\\%. ",
        "introduction": "The galactic centers of spiral galaxies are important regions for the evolution of galaxies. It has been widely believed that most `large' galaxies harbor supermassive black holes (SMBHs) at their centers, and also that the observed scaling relation between masses of black holes and the galactic bulges suggests that they have evolved together during the hierarchical formation of galaxies \\citep[e.g.][]{kauf00}. The galactic centers are often very active. Nuclear starbursts and various kinds of active galactic nuclei (AGNs) are observed in nearby galaxies and also in galaxies at high redshifts.  These active phenomena must be related with the evolution of the SMBHs and galaxies, as well as the intergalactic medium through galactic winds and radiative feedback, especially at high redshifts.  All these phenomena are related to the structure and dynamics of the interstellar medium (ISM) in the central 100 pc regions of galaxies. For example, AGN activities are supported by gas fueling, whose mechanisms remain unclear, especially in the inner 100 pc from the central engine (see Wada 2004, and references therein). We have not yet determined what triggers the nuclear starbursts that are connected to the formation of stellar cores and the SMBH \\citep[e.g.][]{norm88}. However, in the next decade, we may expect drastic change in our observational knowledge of the ISM in the circum-nuclear region of galaxies. The $\\sim$ 0.01'' spatial resolution with the next generation millimeter and submillimeter interferometer, ALMA (the Atacama Large Millimeter Array) may reveal a sub-pc structure of the ISM. We now need to improve our theoretical modeling of the ISM in the galactic central regions in order to make direct comparison with future observations.  The ISM in galaxies is inhomogeneous and multi-phase, and this is especially important in studying the gas dynamics in circum-nuclear regions, because the typical scale of the inhomogeneity of the ISM ($\\sim 10$ pc) is not negligible compared to the size of the system ($\\sim 100$ pc). The isothermal approximation (e.g. the sound velocity $\\sim $ 10 km s$^{-1}$), which has often been used for modeling gas dynamics on a galactic scale, is not adequate for the nuclear gas dynamics, particularly when energy and radiative feedback from stars is taken into account. The multi-phase nature is also important for comparing numerical models with observations. Radiative transfer calculations, for continuum and various lines emitted from the AGN regions, require realistic density, temperature and velocity distributions of the ISM as input. In order to achieve realistic numerical models of the ISM, sub-pc spatial resolution for diffuse gases as well as for dense gases and resolving strong shocks is also essential. We should also take into account the self-gravity and three-dimensional dynamics of the ISM.  Instead of adopting phenomenological approaches for the multi-phase ISM, in which many free parameters and assumptions are unavoidable \\citep[e.g.][]{comb85},\\citet{wad01b, wad02} obtain global, inhomogeneous models of the ISM with a sub-pc scale resolution, by solving the governing basic equations of the ISM with a high-accuracy Euler mesh code. In the present paper,  using these results (three-dimensional density, temperature and velocity field data), we have made the first attempt to derive the molecular line intensities emitted from the nuclear starburst region with a supermassive black hole applying our non-LTE (non Local Thermodynamic Equilibrium), fully three-dimensional radiative transfer calculations.   These kinds of `synthesized observations' based on radiative transfer calculations has recently been applied to turbulent molecular clouds \\citep{pad98,oss02, oss02b}. \\citet{pad98} performed non-LTE radiative transfer calculations of CO and CS lines using a Monte-Carlo integration scheme \\citep{juv97} for numerical results of the magneto-hydrodynamic simulations of turbulent clouds. They found that the synthesized maps and spectra bear resemblance to real molecular clouds in many respects. \\citet{oss02} and \\citet{oss02b} made  similar analysis of turbulent clouds, but they assumed a local LVG (Large Velocity Gradient) approximation and isotropic radiation field. Ossenkopft et al. claimed that their simplified analysis is accurate enough for comparison between the observations of molecular clouds and the numerical models of turbulent clouds. These studies focused on the molecular clouds in disks of spiral galaxies, like our Galaxy, therefore we cannot simply apply their results to the interstellar medium in the nuclear starburst with a supermassive black hole. Apparently we should not assume isothermal, isotropic turbulence in the galactic nuclear region. Since the local turbulent motion is of the same order as the global rotation \\citep{wad02}, a fully consistent treatment of the radiative transfer is necessary.  In \\S 2, we describe the numerical method and models. In \\S 3, the typical structure of the ISM in the circum-nuclear region obtained by hydrodynamic simulations is briefly described. Then, $^{12}$CO and $^{13}$CO line intensity calculations are shown. The CO-to-H$_2$ conversion factor, the so-called ``X'' factor and the line-ratios are also discussed. A brief summary and discussion are presented in \\S 4. ",
        "conclusions": "In this paper, we made the first `observational visualization' of a molecular torus (or disk) around a supermassive black hole in an active galactic nucleus under the dynamical influence of starbursts through supernova explosions.  Three-dimensional hydrodynamic simulations of the multi-phase gas disk/torus are used as an input for the three-dimensional non-LTE radiative transfer calculations for $^{12}$CO and $^{13}$CO lines. We found that clumpy and turbulent structure of the nuclear disk whose radius is about 30 pc can be resolved with a few pc spatial resolution, which will be achieved with the ALMA for nearby galaxies.  Our results suggest that the CO-to-H$_2$ conversion factor (X-factor) is not uniformly distributed in the cental 100 pc, and the X-factor for $^{12}$CO ($J=$1-0) is not constant against the density. On the other hand, $^{12}$CO ($J=$3-2) is nearly constant over the wide range, which is $\\sim 0.3 \\times 10^{20} {\\rm cm}^{-2} ({\\rm K} \\;  {\\rm km \\; s}^{-1})^{-1}$, and it can be used to estimate molecular gas mass in gas-rich galactic central regions associated with active star formation. We also found that the observed line-ratios depend on the beam sizes, and the average line-ratios (e.g. \\twco(4-3)/\\twco(2-1)) do not necessarily represent the `real' ratios of the ISM on a local (parsec) scale. We should be careful to use the ratios to infer the physical conditions of the ISM if the spatial resolution is not fine enough.  Although this is the first attempt to investigate the three-dimensional structure and dynamics of the ISM in the galactic central region, which could be compared directly with observations by next-generation radio interferometers, our methodology still needs further improvement in various respects. Firstly, we have not taken into account the radiation from the AGN itself. The strong UV radiation from the accretion disk could affect both the dynamics and excitation of the molecular gases (see also Ohsuga \\& Umemura 2001). We also assumed a uniform UV radiation field. However this uniformity would not be the case in reality, because the radiation from the massive stars distributed in the clumpy nuclear disk should affect the dissociation and excitation of molecules as well as the UV heating of the ISM non-uniformly. We cannot answer how serious these effects are on the `synthesized observations' presented here, before we more consistently solve the radiation field.  The CO abundance, which is assumed constant in our calculations, could affect the intensity and the CO-to-H$_2$ conversion factor \\citep{saka96}. \\citet{arim96} suggested that the observed CO intensities depend on the metallicity in galaxies, and they found smaller X-factors in the central regions than in outer regions of galaxies. In this paper, we assumed the solar metallicity, and a smaller X than the local value is suggested. Thus the X-factor could be even smaller in the nuclear starburst region with higher metallicity than the Galactic local value by a factor of two or maybe as much as ten. In fact, low values of X have been suggested for the central regions of galaxies \\citep[e.g.][]{wall93, reg00}. It is certain at least that $^{12}$CO ($J=$1-0) with the standard X-factor inferred from observations for our Galaxy should not be used without caution, especially for the dense ISM with strong star forming activities, such as starburst galaxies or ultra-luminous infrared galaxies.  The dense molecular gas can be traced not only by CO lines, but also by ionized and atomic carbon, $[{\\rm C_{II}}] 158 \\mu m$ and by $[{\\rm C_I}]$ lines. Under some conditions, CO molecules are photo-dissociated by UV radiation, but H$_2$ is still present. The neutral carbon may be ionized in the photon-dominated region (PDR). Observations with the Kuiper Airborne Observatory revealed that $[{\\rm C_{II}}]$ emission is correlated with the $^{12}$CO(1-0) emission in gas-rich star forming galaxies \\citep{sta91}. The Balloon-borne Infrared Carbon Explorer (BICE) also found $[{\\rm C_{II}}]$ emission in the galactic plane, which shows similar distribution to CO \\citep{nak98}. The Antarctic Submillimeter Telescope and Remote Observatory (AST/RO) has recently revealed that $[{\\rm C_{I}}]$ emission has a spatial extent similar to low-J CO lines in the Galactic central regions \\citep{mar04}. This is also the case in external galaxies, based on observations with the James Clerk Maxwell Telescope (JCMT)\\citep{isr02}. These observations support the PDR picture \\citep[e.g.][]{tie85} in terms of the molecular clouds illuminated by the nearby OB stars. In the present paper, we do not solve the UV radiation field in the torus, and photo-chemistry in terms of CO, $[{\\rm C_{I}}]$, and $[{\\rm C_{II}}]$ is not taken into account. One should note, therefore, that the CO intensities and the conversion factors shown in \\S 4 might be modified if the photo-chemistry and the UV radiative transfer are fully taken into account. Although such calculations are beyond the scope of the present paper, we can estimate how important the photo-chemical effects are using the published PDR models.  \\citet{moc00} investigated PDR models of molecular clouds, solving chemical and thermal equilibrium with radiative transfer calculations for simple, static molecular clouds. Their method is basically based on the PDR model by \\citet{tie85}, but they assume a spherical cloud instead of the plane-parallel geometry. The molecular clouds considered in \\citet{moc00} have typically a mass of $10-10^5 M_\\odot$, a mean hydrogen density $10^2 < n < 10^4$, and size 2-20 pc, and these are similar to the clumps in our torus models \\citep{wad02}. Mochizuki \\& Nakagawa (2000) found that $[{\\rm C_{II}}]$ emission originates in the outer envelope of the molecular cloud. On the other hand, CO emission is concentrated in the inner region of the cloud. They also found that luminosity of CO and the CO-to-H$_2$ conversion factor are insensitive for the far-UV flux over about two orders of magnitude. The conversion factor is close to the standard X-factor for a relatively massive cloud ($ > 10 M_\\odot$). These results imply that the UV radiation affects mainly the outer envelopes of typical molecular clouds/cores; therefore, the CO intensity distribution and conversion factor presented in the previous section would not be significantly changed even if we took into account the photo-chemistry in our model. Of course, this must ultimately be verified by fully self-consistent radiative transfer calculations with chemistry for the clumpy, turbulent medium.     We showed in this paper that spatial resolution ten-times finer than present-day observations in the radio frequency could drastically change our pictures of the ISM in the central regions of the external galaxies. Combining numerical techniques like the one presented here with the near-future observations by ALMA will be a very powerful tool for investigating the unsolved problems associated with galactic nuclei, such as the starburst-AGN connection and fueling mechanism for the AGNs."
    },
    "0410/astro-ph0410018_arXiv.txt": {
        "abstract": "We present high-quality observations of PSR~B0826$-$34 at 1374 MHz. The emission from this pulsar exhibits strong bursts of pulses followed by long periods of `null' pulses. When it is strong, the radiation extends through the whole pulse period. We show for the first time that there is weak emission during the `null' phases, which should therefore be considered to be a different mode rather than a null.  During this weak mode the profile is similar to that observed in the strong mode at low radio frequency.  Using a phase-tracking method, the pattern of drifting subpulses during the strong mode is seen to be coherent across the whole profile.  The drift rate is variable and includes positive and negative values. Thirteen subpulse bands have been directly observed, covering the whole longitude range. The subpulses and their spacings ($P_2$) are wider in one half of the profile than those in the other half.  This difference, and the variation of observed $P_2$ within the two regions, can be accounted for if the magnetic pole is inclined to the rotation axis by about $0.5^\\circ$.  These two regions appear to represent radiation from outer and inner cones.  The intensity modulation of subpulses in all longitude ranges is related to the magnitude of the drift rate. ",
        "introduction": "PSR~B0826$-$34 (PSR~J0828$-$3417) is known as a pulsar whose emission extends through the whole pulse period ($P = 1.848$~sec). It is a relatively old pulsar with a characteristic age of $3\\times 10^7$yr. PSR~B0826$-$34 exhibits some extreme forms of subpulse drifting behaviour, and it switches between strong and weak emission modes. The weak mode is so weak that it has long been regarded as a complete null (Biggs et al. 1985, B85 hereafter)\\nocite{bmh+85}.  PSR~B0826$-$34 was discovered during the second Molonglo pulsar survey (Manchester et al. 1978)\\nocite{mlt+78}. An early study of this pulsar by Durdin et al. (1979)\\nocite{dll+79} at 408 MHz showed that the pulsar appears to exhibit pulse nulling for at least 70 per cent of the time with observed null intervals from a few to 15 000 periods. The drifting subpulses in this pulsar were first observed by B85.  The integrated pulse profile and most individual pulses extend over all longitudes of the pulse period, implying that this pulsar is an almost aligned rotator (i.e. the inclination angle between the rotation and magnetic axes is very small), and that it is viewed from a direction close to the rotation axis. At 645 MHz, B85 found 5 to 6 bands of drifting subpulses in part of the profile, and noted that the drift rate of these bands shows wide variation, including sign reversals but with no significant variations in the drift-band spacing ($P_2$).  We have recorded two long sequences of individual pulses from PSR B0826$-$34 at 1374~MHz.  At this frequency, there is strong emission through most of the pulse period which allows us to investigate the pattern of drifting subpulses in the strong mode and to trace this pattern round a complete track in the emitting region. ",
        "conclusions": "The main results of this work are as follows:  1. The `nulls' in the pulses from PSR~B0826$-$34 are not true `nulls' and the pulsar presents strong and weak emission modes at 1374 MHz. In its strong emission mode, the radiation extends through the whole pulse period, and region I is stronger than region III at this frequency. In the weak emission mode, the profile is similar to that observed in the strong mode at low radio frequency, at which region I is the weaker.  2. Using a phase-tracking method, the drifting subpulses of the pulsar can be traced through the whole pulse period. Thirteen drift bands have been directly observed, covering the whole longitude range, and the drifting occurs in both directions.  3. The subpulses in region I and their spacing ($P_2$) are wider than those in region III.  These differences can be accounted for if the magnetic pole is inclined to the rotation axis by about $0.5^\\circ$.  4. Regions I and III are consistent with radiation from inner and outer cones; in regions II and IV of longitude, the line-of-sight trajectory cuts through a weak emission region between the two cones.  5 .The observed direction of drift is a true reversal within the polar cap and is not a product of aliasing with the rotation period."
    },
    "0410/astro-ph0410532_arXiv.txt": {
        "abstract": "{ The method developed by Stock \\& Stock (1999) to derive absolute magnitudes and intrinsic colors is applied to simulated low-resolution spectra. The simulation is made by convolving real spectra with a Gaussian function, $\\sigma$ (the full width at half maximum), related to the final spectral resolution. The accuracy with which the stellar parameters are determined indicates that the method may be applied to typical objective-prism spectra. We show that changes in the spectral resolution do not significantly affect the stellar parameters obtained with this method for early-type stars, whereas for later-type stars an improved approach is necessary. } ",
        "introduction": "In a previous work, Stock \\& Stock (1999, hereafter Paper I) developed a quantitative method to obtain stellar physical parameters such as absolute magnitude, intrinsic color, and a metallicity index. These parameters were calculated by the use of pseudo-equivalent widths of absorption features in stellar spectra by means of polynomials. For this purpose, the spectral library of Jones (1999, available at ftp://ftp.noao.edu/catalogs/coudelib) was used which contains about 700 stellar spectra in the bands 3820-4500 \\AA\\ and 4780-5450 \\AA\\, with a resolution FWHM of 1.8 \\AA. The calibration (i.e. the determination of the coefficients of the polynomials) was made using 487 stellar spectra of A-K types. In such spectra, the definition of the true continuum is imprecise. For this reason, in Paper I, the authors worked with pseudo-equivalent widths for the identified features. This was done by selecting one inner and two outer regions for each absorption line. The two outer regions (one on each side of the line) were used to determine the pseudo-continuum, while the integration of the inner region allows us to know the pseudo-equivalent width. This procedure is similar to that used by Worthey et al. (1994).  One of the main goals of the quantitative classification method is its application to a large number of stellar spectra, obtained by means of objective-prism observations, which are a kind of low resolution spectra. Therefore, it is essential to simulate low-resolution spectra. This may be done by applying a mathematical treatment to smooth the spectra of Jones library. Once this is done, it is necessary to re-calibrate the method using those smoothed spectra in order to compare the results with those obtained in Paper I. It is important to know if the method is as accurate as in Paper I when it is used with low resolution spectra.  On the other hand, the equivalent width of a line does not depend, in principle, on the spectral resolution. Considering the way that the pseudo-equivalent widths were defined in Paper I, these could be affected by variations in the spectrum profile. It is also important to analyze the sensitivity of the method to variations in the spectral resolution (due, for instance, to seeing variations during the observations). To do this, first the polynomials are generated applying the method to the smoothed spectra. Then these are applied to spectra of two other resolutions, and the respective results are compared.  A short explanation of how the low-resolution spectra are simulated and how the method is applied to these is given in Sections~2 and 3. Section~4 is dedicated to a discussion of the results, and the main conclusions are summarized in Section~5. ",
        "conclusions": "There are two main conclusions drawn from the present study. First, the application of the quantitative classification method (Paper I) to spectra smoothed with the standard resolution ($\\sigma = 250$), which simulates typical objective-prism spectra, yielded results in good agreement with the previous ones: it was possible to derive absolute magnitudes with an average error of 0.26 magnitudes and $(B-V)_0$ colors with an error of 0.020 magnitudes. This method can then be applied to objective-prism observations with high accuracy.  Second, the results indicate that changes in the resolution of the spectra will not affect the determination of the measurable parameters for the stars belonging to group 2 (earlier-type stars). Concerning to very early stars (spectral types O-B), the continuum can be easily determined and the variations in the spectral resolution are not a problem, neither for the calibration nor for the derivation of the fundamental parameters (see Stock et al. 2002). But, for later-type stars (groups 3 and 4), it is necessary to apply a different approach because the shape of these spectra is very sensitive to resolution variations. We propose that this group should be divided into subgroups in order to reduce the differences between the parameters derived with several spectral resolutions."
    },
    "0410/astro-ph0410704_arXiv.txt": {
        "abstract": "We present the results of a series of time-dependent numerical simulations of cold, magnetocentrifugally launched winds from accretion disks.  The goal of this study is to determine how the mass loading from the disk affects the structure and dynamics of the wind for a given distribution of magnetic field.  Our simulations span four and half decades of mass loading; in the context of a disk with a launching region from $0.1\\AU$ to $1.0\\AU$ around a $1\\solarmass$ star and a field strength of about $20\\gauss$ at the inner disk edge, this amounts to mass loss rates of $1\\times 10^{-9}$ -- $3\\times 10^{-5}\\solarmassyear$ from each side of the disk.  We find that, as expected intuitively, the degree of collimation of the wind increases with mass loading; however even the ``lightest'' wind simulated is significantly collimated compared with the force-free magnetic configuration of the same magnetic flux distribution at the launching surface, which becomes radial at large distances.  The implication is that for flows from young stellar objects a radial field approximation is inappropriate.  Surprisingly, the terminal velocity of the wind and the magnetic lever arm are still well-described by the analytical solutions for a radial field geometry.  We also find that the isodensity contours and Alfv\\'en surface are approximately self-similar in mass loading. The wind becomes unsteady above some critical mass loading rate.  The exact value of the critical rate depends on the (small) velocity with which we inject the material into the wind. For a small enough injection speed, we are able to obtain the first examples of a class of heavily-loaded magnetocentrifugal winds with magnetic fields completely dominated by the toroidal component all the way to the launching surface. The stability of such toroidally dominated winds in 3D will be the subject of a future investigation. ",
        "introduction": "\\subsection{Magnetocentrifugal Wind Solutions}  Astrophysical jets are seen in a variety of systems ranging from young stellar objects to active galactic nuclei.  All these systems with jets are also observed to have accretion disks, and it is likely that the two phenomena are related.  Jets provide accretion disks with an efficient means to remove angular momentum, and disks provide jets with a source of material and magnetic fields.  The magnetocentrifugal model has been proposed \\citep{bp82} as a way to generate a fast, collimated jet from an accretion disk.  In this model, gas flows away from the accretion disk along open magnetic field lines that thread the disk. The magnetic field accelerates the gas to super-fast magnetosonic speeds, and the rotation of the gas winds the field up into a predominantly toroidal configuration which is thought to be able to collimate the gas through ``hoop'' stress (see, however, \\citealt{ok99}).  A full analytical treatment of the steady flow is difficult because the equation describing the cross-field force balance, the Grad-Shafranov equation, changes its character from elliptic to hyperbolic upon crossing the fast magnetosonic surface, the location of which is not known a priori. In order to attack the problem, authors have had to make simplifying assumptions.  \\citet{bp82} presented an analytical solution for this flow under the assumption of self-similarity in the original paper on magnetocentrifugal jets, and self-similarity has also been employed by others, e.g.\\ \\citet{cl94}, \\citet{o97}, \\citet{tts97} and \\citet{fc04}. Another approach has been to make an assumption about the magnetic field configuration and then solve for the hydrodynamic quantities, e.g., \\citet{pn86} and \\citet{ks97}.  An asymptotic analysis of the Grad-Shafranov equation can yield the structure of the flow at large distances from the source, as in \\citet{hn89, hn03} and \\citet{shu95}.  Because of the difficulties in finding analytical solutions to the flow structure, it has become common in recent years to attack the problem with time-dependent numerical simulations. Two different approaches have been taken in developing numerical models of jet launching from accretion disks --- including the accretion disk as part of the simulation, or treating the disk as a fixed boundary condition.  We have chosen the latter approach for our study and will focus on it.  We refer the reader to \\citet{su86}, \\citet{sn94}, \\citet{m96}, and \\citet{vonR04} for examples of models which include the disk.  For cases where the disk is treated as a fixed rotating boundary for the simulation, it is important that the hydromagnetic quantities are specified on the disk in a self-consistent manner.  Several different approaches have been taken for specifying both the initial and boundary conditions for numerical simulations of magnetocentrifugal jet launching.  \\citet{u95} performed simulations demonstrating the acceleration and collimation of a jet from a disk with a hot (plasma $\\beta$ parameter $\\gg1$) corona. Their simulations did not reach steady-state, although it is unclear whether this is a feature of their model or the result of terminating the simulation before it became stationary.  A similar study \\citep{r97} using a stronger magnetic field and a cold ($\\beta<1$) corona resulted in stationary wind solutions; however the wind was at best poorly collimated.  \\citet{op97a,op97b} also use a corona that can be described as cold, and two different initial magnetic configurations:  a potential (current-free) field \\citep{op97a}, and a uniform, vertical (parallel to rotation axis) field \\citep{op97b}.  Both cases exhibit acceleration and collimation, but the former comes to a steady-state, while the latter is characterized by episodic behavior.  However, in a follow-up paper \\citep{op99}, they argue that it is not the initial magnetic configuration but rather the mass loading at the base of the wind that determines whether a given model will reach steady-state or display episodic behavior.  Their conclusions are based on simulations they performed for different values of mass load, for both the potential and uniform field configurations.  \\citet{klb99}, hereafter referred to as KLB99, also performed cold simulations in an initially potential field, with an important modification.  They pinned the foot points of the wind-launching field lines on the disk, but left the field lines free to vary in the radial and azimuthal directions both on the disk and away from the disk in the wind zone. These degrees of freedom are needed to prevent sharp discontinuities from developing between the disk boundary and the magnetocentrigual wind (cf.\\ \\citealt{bt99} for the case of a spherical boundary).  All the works mentioned in the preceding paragraph use a magnetic field threading the entirety of the disk (i.e.\\ the entire $z=0$ plane). However, a launching region that extends to the edge of the simulation box results in a portion of the flow reaching the boundary before having been accelerated to super-fast magnetosonic speeds.  As pointed out in KLB99, this has the effect of making the results of the simulation sensitive to the size and shape of the simulation box, which is undesirable. Therefore we have chosen to present simulations with a limited launching region in this paper, as was done in \\citet{klb03} (KLB03 hereafter). The size of the launching region is left as a free parameter, which has the benefit of allowing us potentially to model another proposed candidate for jet production, the X-wind model \\citep{shu94}. This model also uses the magnetocentrifugal effect to accelerate and collimate a wind, but the required magnetic field, rather than coming from a wide area on the accretion disk (the so-called disk wind, see \\citealt{kp00}), instead comes from a narrow region around $\\varpi=R_{X}$, where $R_{X}$ is the point where the accretion disk is in co-rotation with the stellar magnetosphere.  \\subsection{Effects of Mass Loading}  The focus of our investigation will be on the effects of mass loading on the structure of magnetocentrifugal winds. It is motivated in part by the observations of the outflows from young stellar objects (YSOs). \\citet{bon96} established that deeply embedded Class 0 objects typically drive CO molecular outflows with momentum fluxes more than an order of magnitude higher than those of older Class I sources. If the molecular outflow is driven by a magnetocentrifugal wind from a region close to the central object in a momentum conserving fashion, as currently believed, then the wind should have a higher mass flux in the Class 0 phase, since its wind speed can hardly be more than an order of magnitude faster; if anything, the fact that Class 0 sources are yet to accrete most of their final masses points to a smaller Keplerian speed in the wind-launching region, which would argue for a lower wind speed. The decrease in mass flux as the wind ages is probably correlated with the decrease in disk accretion rate. The same trend appears to continue into the optically revealed, classical T Tauri phase, where the primary wind can be probed directly in a number of ways \\citep{erm93}. From the early embedded to late revealed phase, the mass loss rate in a typical YSO wind may span a large range, from $\\sim 10^{-6}\\solarmassyear$ \\citep{bon96} to $\\sim 10^{-9}\\solarmassyear$ or less \\citep{ed93}. The question is then: how does the mass loading affect the structure of the wind?  Part of the answer can be obtained analytically, by considering the wind properties along a flux tube of prescribed opening. \\citet{s96} examined the special case of a cold wind in a radial field geometry. He introduced a dimensionless parameter which, for a non-radial geometry, is equivalent to \\begin{equation} \\mu\\equiv {4\\pi \\rho v_p \\varpi\\Omega \\over B_p^{2}}, \\label{para1} \\end{equation} where the mass density $\\rho$, poloidal flow velocity $v_p$, cylindrical radius $\\varpi$, angular speed $\\Omega$, and poloidal field strength $B_p$ are evaluated at the base of the wind. The parameter measures the ability of the rotating magnetic field $B_p$ to accelerate the mass load, which is proportional to $\\rho v_p$. We will refer to $\\mu$ as the mass loading parameter. In the ``light'' wind limit $\\mu \\ll 1$, \\citet{s96} found that the flow is accelerated centrifugally along the more or less rigid field line, up to the Alfv\\'en radius $\\varpi_A$, which is located about $(3/2)^{1/2}\\mu^{-1/3}$ times the foot point radius. In the opposite, ``heavy'' wind limit $\\mu\\gg 1$, the Alfv\\'en radius moves to $(3/2)^{1/2}$ times the foot point radius. The field lines become tightly wound, and the flow is accelerated mainly by magnetic pressure gradients. In both parameter regimes, the terminal wind speed is given by \\begin{equation} v_{\\infty}=\\mu^{-1/3} v_0, \\label{terminal} \\end{equation} where $v_0=\\Omega\\varpi$ is the rotational speed at the launching surface. The above relation results from the fast magnetosonic point being at infinity for a radially diverging flux tube, as in \\citeauthor{m69}'s (\\citeyear{m69}) minimum energy solution for cold relativistic MHD winds \\citep{gj70}. \\citet{ks97} showed that the above relation holds approximately, at least for light winds, even when the prescribed field geometry is non-radial, as long as the fast point (which is now at a finite radius) is not too close to the foot point. We therefore expect no major surprises for the wind acceleration, which is mainly determined ``locally'' along a field line by the conserved quantities (see equations [\\ref{eqn:kappa}]--[\\ref{eqn:e}] below).  The effects of mass loading on the wind collimation are more difficult to determine. Collimation is a global wind property controlled by the cross-field force balance. One must solve the Grad-Shafranov equation, which has not been possible in general until recently. In this paper, we will use the MHD code developed in KLB03 to solve for the steady-state wind structure through time-dependent simulations. The goal is to study the effects of mass loading on both wind collimation and acceleration. We find that, as expected, a heavier loading leads to a better collimation and a slower wind, and that the slowdown can be described to a good approximation by equation (\\ref{terminal}). We also show that there exists a maximum mass load for a given wind-launching magnetic field, beyond which the magnetocentrifugal mechanism shuts off, and that the maximum corresponds to the dimensionless parameter $\\mu\\gsim 1$, with the exact value depending on the injection velocity at the launching surface. We show that, despite significant collimation in the wind, some of its most important properties can still be well-described by the analytic results originally derived for a wind in a radial field geometry.  The remainder of the paper is organized as follows.  In \\S\\,\\ref{desc} we describe our formulation of the magnetocentrifugal wind problem and the setup of numerical simulation. Our reference simulation is presented in detail in \\S\\,\\ref{reference}.  We compare models with different mass loads and distributions of mass loading in \\S\\,\\ref{loadmagn} and \\S\\,\\ref{loaddist} respectively. A discussion of these results, along with our conclussions, are given in \\S\\,\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  \\subsection{Mass Load Limit and Heavy Magnetocentrifugal Winds} \\label{heavywind}  A novel finding of our paper is that, for a given wind launching magnetic field, there exists a maximum mass load beyond which cold, steady state solution does not exist. For a relatively large wind injection speed of $10\\%$ Keplerian, the maximum corresponds to a dimensionless parameter $\\mu \\sim 1$, which is roughly where the transition from ``light'' to ``heavy'' wind occurs. With the reduction of injection speed to $1\\%$ Keplerian we can obtain rather steady solutions up to $\\mu \\sim 10$. The tendency for the solutions in the heavy wind regime to remain unsteady may not be too surprising, given that their magnetic fields are dominated by the toroidal component all the way to the launching surface. This is illustrated in Fig.\\,\\ref{fig:f13} for the solution shown in lower panel of Fig.\\,\\ref{fig:f10}, where the mass loading is 1000 times higher than the reference solution. The toroidal field in the heavy wind is everywhere greater than the poloidal field in the magnetocentrifugal part of the wind.  For comparison, the toroidal field in the light, reference solution starts to exceed the poloidal field only well beyond the launching surface.   At the launching surface itself, the ``heavy'' wind has a ratio $|B_\\phi|/B_p$ of 4, compared to a value of $0.2$ for the ``light'' wind. The toroidally dominated magnetic configurations are prone to pinch instability in lab experiments \\citep{ks54}. A possible indication of the onset of this type of instability is the waviness of field lines in solutions with mass loads near but below the threshold for transition to chaotic flows (see Fig.\\,\\ref{fig:f9}). The fact that the threshold increases with decreasing injection speed at the launching surface points to another possibility for the nonsteadiness that involves the gravity of the central object --- loss of mechanical balance near the launching surface for heavy winds.  The reason for the loss of balance can be seen from equation (\\ref{eqn:omega}) for the conservation of angular velocity $\\Omega$ along any field line in steady state. The equation can be cast into \\begin{equation} v_\\phi=\\Omega \\varpi + v_p B_\\phi/B_p. \\label{rotation} \\end{equation} At the launching surface, the first term on the right hand side is set to the Keplerian speed in our simulations, and the second represents the speed that the wind fluid lags behind the Keplerian rotation (note that $B_\\phi$ is negative). For a given injection speed $v_p$, the lag increases with the twisting of the field lines (i.e., the ratio of $\\vert B_\\phi\\vert/B_p$), which in turn increases with mass loading. It is conceivable that beyond some threshold in mass load, the fluid rotation near the launching surface becomes too sub-Keplerian for the centrifugal force to balance the inward pull of the central gravity \\citep{k00}. The force imbalance would lead to a radial infall of wind material, which is indeed observed in the initial phase of the development of chaotic flows, such as the one shown in upper panel of Fig.\\,\\ref{fig:f10}. The development of the chaotic behavior is reminiscent of that of the channel flows observed in simulations of weakly magnetized accretion disks (J. Hawley, priv.\\ comm.). As one decreases the injection speed, the lag speed is decreased for a given magnetic field twist. This is consistent with our finding that decreasing the injection speed tends to make a heavy wind more steady. By choosing a small enough injection speed we were able to obtain more or less steady solutions well into the heavy wind regime, with $\\mu \\gg 1$. The fact that such heavy wind solutions can be obtained from {\\it time-dependent} simulations suggests that they are stable to, or at least not disrupted by, axisymmetric (pinch) instabilities, despite their magnetic fields being completely dominated by the toroidal component. Whether they can resist disruption by non-axisymmetric (kink) instabilities in 3D remain to be determined.  Even for an arbitrarily small injection speed, we expect the wind solution to become chaotic beyond a certain mass load. The reason is still related to equation (\\ref{eqn:omega}), which applies to both the launching surface and above. For a slowly injected wind, acceleration above the launching surface increases the poloidal speed, causing the fluid to rotate significantly below the rate needed for radial force balance when the magnetic field becomes toroidally dominated. Material fed into the wind starts to fall toward the central star, dragging magnetic field lines along with it. The inward bending of field lines further reduces the centrifugal acceleration and thus exacerbates the force imbalance, again reminiscent of the development of channel flows in weakly magnetized disks through magneto-rotational instability \\citep{hb91}. This positive feedback can be cut off if the mass loading (which is fixed in our simulations) is allowed to vary in response to the field bending. The mass loading is expected to drop drastically when the (poloidal) field lines bend within roughly $30\\degr$ of the disk normal, the minimum angle for mass-loading by the magnetocentrifugal mechanism (\\citealt{bp82}; \\citealt{ogil01}). The reduction is a natural way for the wind to self-limit its amount of mass loading. Quantifying this process requires a detailed treatment of the disk-wind coupling, which is beyond the scope of the present work (see, e.g., the early analytic work of \\citealt{wk93} and recent numerical simulations of \\citealt{ck02}).  \\citet{op99} addressed the issue of mass loading through numerical simulations. They found a transition from steady to non-steady wind solution as the mass load is decreased below some critical value; this trend is the opposite of the one we find. One possibility for the difference lies in the simulation setup at the launching surface. In their simulations, the toroidal field strength is prescribed at the disk surface, along with the magnetic flux distribution. This prescription essentially fixed the amount of energy (and angular momentum) flux extracted by the rotating magnetic fields from the disk proper (see the last term in the expression, equation~[10], for the specific energy $E$). In general, these prescribed amounts cannot be carried away from the disk by the loaded wind material, as illustrated in \\citet{op97a}. They found a discontinuity in $B_\\phi$ between the disk and the active zones in a test simulation with $B_\\phi$ set to zero on the disk. The discontinuity indicates to us that the toroidal field near the base of the wind is trying to respond to stresses further out, while being constrained by a (formally) mismatched boundary condition. In our simulations, the toroidal field is allowed to adjust in response to the stresses in the wind both on the launching surface and in the active zones. We are able to obtain steady solutions for very low mass loads, limited only by the size of the simulation box (since the fast surface, which needs to be enclosed within the box to minimize boundary effects, increases in size as mass load decreases), and by computation time, because a lighter wind has a larger Alfv\\'en speed, which requires a smaller timestep to simulate. Another possibility for the difference is that our magnetocentrifugal wind becomes super fast-magnetosonic within the computational box, whereas in their simulations a significant fraction of outflow remains sub-fast up to the outer boundary. One worry is that, in such sub-fast regions, signals can propagate from the outer boundary to the launching surface. In the numerical experiments of KLB99, the coupling between the outer boundary and launching surface appears capable of driving the magnetocentrifugal wind unsteady or even chaotic.  The non-steadiness in our high mass load winds has a different origin: it occurs when the field is in some sense too weak to accelerate the prescribed mass load steadily. In our interpretation, the maximum in mass loading has its root in the inability of a heavily loaded wind to find a stable cross-field force balance. Such a force-balance cannot be accounted for in 1D wind models along a prescribed flux tube. Even 2D (axisymmetric) models obtained by directly solving the steady wind equations, such as those of \\citet{sa87} and \\citet{ns94}, fail to uncover this maximum. This is probably because the construction of steady wind solutions directly is difficult, and it has not been possible to explore a large parameter space as we did in this paper using the time-dependent simulations. One type of steady solutions for which parameter exploration is feasible is the self-similar solutions of \\citet{bp82}. However, the cross-field force balance of such solutions must be interpreted with care, because of singularities near the axis and at the infinity. For example, \\citet{sa87} pointed out that the ``standard'' solution of Blandford \\& Payne is less collimated than the potential magnetic field with the same flux distribution on the disk, as the result of a singular current density on the axis. Also, the self-similar solutions tend to recollimate at large distances, which is not observed in our non-self-similar steady solutions.  \\subsection{Scaling Laws for Magnetocentrifugal Winds}  Self-collimation of streamlines is a basic feature of magnetocentrifugal winds. The general expectation is that, for a given distribution of magnetic flux, the wind becomes better collimated as the amount of mass loading increases (e.g., \\citealt{pp92}). This expectation is borne out by our solutions (see Fig.\\,\\ref{fig:f4}). The increase in the degree of collimation with mass load appears to be slow; fitting a power-law distribution to the cylindrical radius of a representative streamline (at a height of $100\\AU$) listed in Table\\,\\ref{tab:t1} as a function of mass load yields a rough relationship, $\\varpi_j\\propto {\\dot M}_w^{0.1}$. This relationship does not apply to all streamlines. In particular, the first (on axis) and last (on equator) streamlines are the same for all of our winds (by design). The space-filling constraint tends to discourage rapid streamline collimation, particularly for winds launched from a region that is small compared to the computational domain. Nevertheless, it is difficult to predict the location of each individual streamline a priori. This needs to be determined through numerical computation.  A somewhat surprising result is that the wind speed along any given streamline scales with the mass load as ${\\dot M}_w^{-1/3}$, originally derived for radial geometry, despite the fact that the streamline collimates to different degrees for different mass loads. A related result is that the Alfv\\'en radius follows closely the analytical relation $\\varpi_A (\\mu)$ for a radial wind. Both results indicate that the energy and angular momentum extraction from the disk and the wind acceleration are insensitive to the relatively small differences in the degree of flow collimation for various mass loads. The fact that the appropriately scaled density contours and Alfv\\'en surfaces closely align implies that magnetocentrifugal winds of different mass loads are approximately self-similar.  The distribution of mass loading also affects flow collimation. We have demonstrated that loading more material in the outer part of the wind tends to produce a better collimated streamline from a given location on the disk. This is probably due to a preferential increase in the toroidal field strength of the outer part, which enhances the cross-field gradient of the product $B_\\phi\\varpi$, which is largely responsible for flow collimation. On the other hand, concentrating more material near the inner edge of the Keplerian disk tends to produce better collimated density contours. The fact that jet-like density structure can be produced naturally, together with the detection of rotation signatures in T Tauri jets (\\citealt{bac02}; \\citealt{coff04}), lends strong support to the magnetocentrifugal model of jet formation (\\citealt{shu95}; \\citealt{and03}).  To summarize, we have numerically obtained axisymmetric magnetocentrifugal wind solutions for a wide range of mass loading for a given wind-launching field configuration. The range is limited from below by computation time, and from above by a form of instability that causes the wind to become chaotic. It testifies to the robustness of the magnetocentrifugal mechanism for outflow production. Whether the mass loading in actual astrophysical systems, such as YSO winds, covers as wide a range remains to be determined. It will probably be limited by the detailed physics of mass supply from the disk, and by the ability of the wind to withstand disruption by instabilities in 3D."
    },
    "0410/astro-ph0410197_arXiv.txt": {
        "abstract": "The eigenfrequencies of nonradial oscillations are a powerful probe of a star's interior structure. This is especially true when there exist discontinuities such as at the neutron star (NS) ocean/crust boundary, as first noted by McDermott, Van Horn \\& Hansen. The interface mode associated with this boundary has subsequently been neglected in studies of stellar nonradial oscillations. We revisit this mode, investigating its properties both analytically and numerically for a simple NS envelope model. We find that it acts like a shallow surface ocean wave, but with a large radial displacement at the ocean/crust boundary due to flexing of the crust with shear modulus $\\mu\\ll P$, the pressure. This displacement lowers the mode's frequency by a factor of $\\sim(\\mu/P)^{1/2}\\sim0.1$ in comparison to a shallow surface wave frequency on a hard surface. The interface mode may be excited on accreting or bursting NSs and future work on nonradial oscillations should consider this mode. Our work also implies an additional mode on massive and/or cold white dwarfs with crystalline cores, which may have a frequency between the f-mode and g-modes, an otherwise empty part of the frequency domain. ",
        "introduction": "A wide variety of nonradial oscillations have been previously investigated for neutron stars (NSs) with the majority of these studies focusing on surface g-modes (oscillations whose restoring force is gravity; McDermott, Van Horn \\& Scholl 1983; Finn 1987; McDermott \\& Taam 1987; McDermott, Van Horn \\& Hansen 1988, hereafter MVH88; McDermott 1990; Strohmayer 1993; Bildsten \\& Cutler 1995, hereafter BC95; Bildsten, Ushomirsky \\& Cutler 1996; Strohmayer \\& Lee 1996; Bildsten \\& Cumming 1998; Piro \\& Bildsten 2004) because their frequencies are similar to many of the oscillations observed from NSs (for example burst oscillations; Muno et al. 2001). These modes are confined to outer regions of the NS surface due to the shear modulus of the NS crust (as illustrated by BC95 and further discussed here), so that their frequencies are independent of the deep NS crust composition and small amounts of energy will lead to large amplitudes. However one mode has consistently been neglected, which is the wave analogous to a shallow water wave in the outermost parts of the star. Unlike a global surface wave (i.e. f-mode), we expect this wave to be confined to the NS ocean just like the g-modes, but a correct prediction of its frequency requires proper treatment of its eigenfunction at the ocean/crust boundary.  A study of nonradial oscillations on non-rotating, non-magnetic NSs was presented by MVH88 for three-component NS models (consisting of an ocean, crust, and core), examining g-modes, acoustic oscillations (p-modes), and f-modes. In addition, MVH88 reported a new set of oscillations, which they called ``interface modes.'' These modes are unique in that they ``live'' (have amplitudes concentrated) around discontinuities in the NS structure. For their three-component model, they found these modes at both the ocean/crust interface and the crust/core interface. We now undertake a detailed investigation of the ocean/crust interface mode so that it may be used for accreting NSs and in the other astrophysical context where a shear modulus is important, crystalline white dwarfs (WDs) (e.g. Hansen \\& Van Horn 1979, hereafter HV79; Montgomery \\& Winget 1999).  To focus our analysis we consider the interface mode for a two-component NS model consisting of an ocean and crust. We use this model because of the simplification provided by the plane parallel geometry of the outer NS layers. The most important property of the interface mode is how it couples with the top of the NS crust (Figure 1). MVH88 showed that the interface mode has a large and approximately constant transverse velocity in the ocean, similar to a shallow water wave. Indeed this mode can be thought of as a surface wave riding on a compressible crust. It is therefore not a toroidal mode as one might expect if the interface wave were like a shear wave (see Appendix A). As its pressure perturbation rolls over the crust's surface, the crust is compressed and sinks. This flexing causes the interface mode to have a large, negative radial displacement at the discontinuity (see the right diagram of Figure 1) in contrast with the typical boundary condition for surface water waves, which treats the crustal boundary as solid with no radial displacement (see the left diagram of Figure 1). Since the pressure perturbation also corresponds to a small, positive swell at the top of the ocean, the radial displacement eigenfunction must necessarily have a single node within the ocean. \\begin{figure*} \\epsscale{0.96} \\plotone{f1.eps} \\figcaption{A schematic diagram of an ocean surface wave showing how it is altered by the flexing of the crust. The vectors show the wave velocity. In both cases the transverse velocity is approximately constant within the ocean and dies off quickly into the crust. If the crust is treated as if it is ``solid'' the radial displacement goes to zero at the ocean floor (left diagram). As shown on the right, the interface mode has a large, negative radial displacement at the crust which is actually larger than the corresponding swell near the ocean surface. This implies that its radial eigenfunction must necessarily have a single node within the ocean.} \\end{figure*}  This simple picture describes the interface mode quite well. The frequency depends on the properties at the base of the ocean, like a surface ocean wave, but with a frequency that is altered by the non-zero radial displacement at the ocean/crust boundary. This boundary condition reduces the frequency by $\\sim(\\mu/P)^{1/2}\\sim0.1$ from what is expected if the crust were ``solid,'' where $\\mu$ is the shear modulus and $P$ is the pressure at the top of the crust. Most of the energy of the mode is in the ocean, just as in a shallow surface wave.  We derive the adiabatic perturbation equations for a plane parallel layer in hydrostatic balance with a shear stress in \\S 2. These are used in \\S 3 to analytically estimate the interface mode's radial eigenfunction and frequency.  We follow this with a numerical study in \\S 4 that confirms our analytic results and also provides more detail about the eigenfunctions. We conclude in \\S 5 with a discussion of possible uses for the interface mode in astrophysical contexts. Maybe most exciting among these is an additional mode of oscillation in WDs with crystalline cores. ",
        "conclusions": "We have discussed the ocean/crust interface wave, investigating its properties both analytically and numerically for a two-component NS envelope. We show that it can be understood as a shallow surface wave, with a different radial displacement boundary condition at the bottom of the ocean. This reduces the frequency by roughly $\\sim\\mu_0^{1/2}\\sim0.1$ from what would be expected if the crust were instead ``solid.'' In general the interface wave is a global mode of oscillation, so that the entire star must be taken into account to calculate the mode to high precision (see discussion in Appendix C). Fortunately, by using our analysis, a fairly good estimate of the mode frequency can be made just from considering the NS properties right near the ocean/crust interface because the majority of mode energy ($\\gtrsim99\\%$) resides in the NS ocean.  The interface wave should be included in future studies of nonradial oscillations on both accreting and bursting NSs. For a non-rotating NS this would result in frequencies of $\\sim10-20\\textrm{ Hz}$ which would then be altered and split due to the NS spin (Bildsten, Ushomirsky \\& Cutler 1996). The observation of multiple mode periods would greatly assist our understanding of the NS crust. The g-modes are sensitive to the depth of the crust (since it provides a bottom boundary), while the interface wave is sensitive to both the depth and the value of the shear modulus. This provides information about the temperature and composition of the crust, which is determined by unstable nuclear burning processes (X-ray bursts and superbursts) that happen in the ocean.  Since the example considered in this work is a relatively ``cool'' envelope, the g-mode frequencies ($\\lesssim6\\textrm{ Hz}$) are well separated from the interface wave frequency ($\\sim10-20\\textrm{ Hz}$). On a bursting NS, which has temperatures in the upper atmosphere as high as $10^9\\textrm{ K}$, at least one g-mode has a frequency higher than the interface wave. The interface wave would then have an additional node in its eigenfunction, due to the relation between g-mode frequency ordering and number of nodes, but its frequency would be unchanged as long as the ocean stays at a temperature of $\\sim4\\times10^8\\textrm{ K}$ (since this is where the mode's energy is concentrated). Woosley et al. (2004) show that the thermal wave during an X-ray burst does not penetrate deeply into the ocean, so that our estimates for the interface mode frequency still hold.  The other important astrophysical context where our work can be applied is for old and/or massive WDs that have crystalline cores. This may be seen in the pulsations of BPM 37093 (Kanaan et al. 1992), a high-mass ZZ Ceti star that should be considerably crystallized (Winget et al. 1997). There is also the exciting possibility that some of the pulsators discovered through the Sloan Digital Sky Survey (Mukadam et al. 2004) may be massive enough to be crystallized. Scaling equation (37) to the fiducial interior properties of a $\\approx1.1M_\\odot$ WD, we find a crustal mode period of \\begin{eqnarray} P&\\approx&95\\textrm{ s} \\left(\\frac{173}{\\Gamma}\\right)^{1/2} \\left(\\frac{10^7\\textrm{ K}}{T}\\right)^{1/2} \\nonumber \\\\ &&\\times\\left(\\frac{A}{16}\\right)^{1/2} \\left(\\frac{R}{5\\times10^8\\textrm{ cm}}\\right) \\left[\\frac{2}{l(l+1)}\\right]^{1/2}, \\end{eqnarray} which for the $l=1$ mode is between the periods of the f-mode ($\\sim10-20\\textrm{ s}$) and g-modes ($\\gtrsim300\\textrm{ s}$). HV79 (and more recently Montgomery \\& Winget 1999) previously calculated the spectrum of g-modes and torsional modes expected for these stars, concentrating on the effects of crystallization. HV79 do not single out any one mode as a ``crustal mode,'' but in their crystallized models their lowest order g-mode (which they denote $_{-1}S_{l}$) matches many of the properties we identify as being characteristic of a crustal mode. This includes a large radial displacement at the crust and a nearly constant transverse displacement in the WD ``ocean'' (see Figures 1 and 2 of HV79). Furthermore, HV79 find the period of their $_{-1}S_{l}$ mode to be between the f-mode and the other g-mode periods, in confirmation of our above discussion.  We thank both Phil Arras and Omer Blaes for helpful discussions. We also thank the referee, Michael Montgomery, for pointing out the HV79 paper and its relevance to our work. This work was supported by the National Science Foundation under grants PHY99-07949 and AST02-05956, and by the Joint Institute for Nuclear Astrophysics through NSF grant PHY02-16783.  \\begin{appendix}"
    },
    "0410/astro-ph0410474_arXiv.txt": {
        "abstract": "The G11.11-0.12 infrared-dark cloud has a filamentary appearance, both in extinction against the diffuse infrared emission of the Galactic plane and in emission at 850$\\micron$.  We use a novel computational technique based on an advanced genetic algorithm to explore thoroughly 3 different models of self-gravitating, pressure truncated filaments and to constrain their parameters.  Specifically, the models tested are the non-magnetic \\citet{ostriker64} model, a generalized version of the magnetic \\citet{stod63} model, and the magnetic \\citet{FP1} model.  Previous results showed that G11.11-0.12 has a much steeper $\\sim r^{-4}$ radial density profile than other filaments, where the density varies approximately as $r^{-2}$, and that this steep density profile is consistent with the \\citet{ostriker64} model.  We present a more complete analysis that shows that the radial structure of G11.11-0.12 is consistent with regimes of each of these models.  All of the magnetic models that agree with the data are threaded by a dominant poloidal magnetic field, and most have dynamically significant fields.  Thus, G11.11-0.12 is an excellent candidate for radial support by a magnetic field that is predominantly poloidal.  We predict the polarization patterns expected for both magnetic models and show that the two magnetic models produce different polarization patterns that should be distingished by observations. ",
        "introduction": "\\label{sec:intro} Infrared-dark clouds (IRDCs) are a class of molecular clouds seen in extinction against the diffuse infrared emission of the Galactic plane.  About 2000 of these objects were discovered by the Midcourse Space Experiment (MSX) satellite in spectral bands from $7-25\\micron$ \\citep{egan98}.  They reside at distances of $2-8$ kpc typically \\citep{egan98,carey98}, and most are located roughly toward the Galactic centre.  The IRDCs have unusual physical characteristics compared to nearby examples of molecular clouds.  \\citet{egan98} found that IRDCs appear to have properties intermediate between massive molecular cloud cores ($n\\approx 10^4-10^6$ {\\rm cm}$^{-3}$; $T\\approx 30-100$ K) and isolated dark clouds ($n\\approx 10^2-10^4 {\\rm cm}^{-3}$; $T\\approx 8-15$ K), having densities typical of molecular cloud cores, but temperatures typical of isolated dark clouds.  Many IRDC filaments appear to be at a relatively early phase of evolution, with fewer signposts of star formation than nearby molecular clouds.  The IRDCs and nearby molecular clouds share the common feature that a large fraction of both populations are filamentary, but the IRDCs are often less fragmented.  Relatively ``pristine'' IRDC filaments have simpler density structures than fragmented, star-forming filaments, which makes them good candidates for detailed comparison with theoretical models.  One drawback is that IRDCs are more distant than many other filaments and are therefore less well resolved.  The G11.11-0.12 cloud studied in this paper is a good example of a filamentary IRDC, which was observed by SCUBA in 850$\\micron$ emission \\citet{carey00}, following its discovery by the MSX survey.  \\citet{johnstone03} (hereafter Paper I) fit a non-magnetic equilibrium filament model to the radial structure of two regions of G11.11-0.12.  This paper presents a more thorough analysis, which compares the SCUBA 850$\\micron$ emission with 3 existing models of filamentary clouds in detail.  These models are the pressure-truncated non-magnetic Ostriker \\citep{ostriker64} model, a generalized \\citet{stod63} (hereafter GS) helical field model, and the Fiege \\& Pudritz (hereafter FP) helical field model.  We discuss these models in the next section.  Note that the Ostriker model is the zero magnetic field limit of both the GS and FP models, but this limit is sufficiently important to warrant separate treatment.  This paper uses a novel genetic algorithm-based technique to explore the parameter space of each model completely and to fit solutions to the data.  This is the first astronomical research application for a new multi-object genetic algorithm (``Ferret'') developed by one of the authors (J.F.).  The technique is discussed in Section \\ref{sec:GAs}. ",
        "conclusions": "\\label{sec:discussion} This paper is the first astronomical research problem attempted using a new multi-objective genetic algorithm ``Ferret'' designed by one of the authors (J.F.).  We have tested this code thoroughly, and it performs extremely well on data-fitting problems where the goal is to find the {\\em class} of allowed solutions, rather than a single representative solution.  This allows the user to understand the detailed structure of the parameter space, and gives an honest representation of the degeneracies of the models being tested.  Our previous analysis of the G11.11-0.12 infrared-dark cloud (Paper I) demonstrated that its density profile is steeper than the $~r^{-2}$ profile inferred from observations of other filamentary molecular clouds (Alves et al. 1998, Lada et al 1999, Johnstone \\& Bally 1999), and that the profile can be fit reasonably well by the non-magnetic Ostriker model.  This paper provides a more detailed analysis that compares the radial structure of the G11.11-0.12 infrared-dark cloud with three models of filamentary clouds: the non-magnetic Ostriker model, the magnetic GS model, and the magnetic FP model.  We found that the density profile of G11.11-0.12 is consistent with the Ostriker model, but also with regimes of both magnetic models.  Our technique allowed us to map out the allowed regions of parameter space thoroughly for each model.  Our most significant result is that nearly all magnetic solutions that are in agreement with the data are dominated by the poloidal component of the magnetic field, or are magnetically neutral.  A few toroidal-field dominated solutions are allowed, but only for Region 2, which has a much ``noisier'' radial structure than Region 1, and even these solutions are close to magnetic neutrality.  G11.11-0.12 appears to be very different from the filaments studied by FP1, which were inferred to have a shallow $~r^{-2}$ density profile and a dominant toroidal field component.  Note that the FP model is the only model of the three models considered that can account for both toroidal-field dominated filaments and the poloidal-field dominated G11.11-0.12 filament.  We predicted the polarization patterns allowed by the GS and FP models of G11.11-0.12 that are allowed by the data.  For the GS model, all of the polarization patterns are Type 2 (poloidal-field dominated) or Type 3 (mixed) patterns \\citet{FP4}.  Thus, the GS model could be ruled out for the G11.11-0.12 filament if future polarimetric observations reveal that this region is more consistent with a Type 1 (toroidal-field dominated) pattern.  For the FP model, Type 3 patterns \\citet{FP4}, which are consistent with a dominant poloidal magnetic field, occupy the greatest fraction of the parameter space.  This is due primarily to the large range of inclination angles (-45$^o$ - 45$^o$) present in the solution set, since \\citet{FP4} showed that type 3 patterns are the most common pattern when the filament is inclined relative to the plane of the sky.  Type 2 patterns occupy a fraction of the parameter space, and occur at moderate values of $\\Gz$ and low values of $\\Gphi$.  We found a few Type 1 patterns that would {\\em appear} to be dominated by a toroidal field.  Most of these solutions have very weak magnetic fields that are close to magnetic neutrality or slightly dominated by the poloidal field, according to the virial equation \\ref{eq:virial}.   We demonstrated that the polarization class, combined with the maximum polarization percentage, provides an estimate of the ratio of magnetic flux to mass values $\\Gz/\\Gphi$.  This is significant because \\citet{FP2} demonstrated that this quantity is the dominant factor that controls the fragmentation timescale of a filament, the fragmentation wavelength, and ultimately the masses of clumps that form along the filament.  Thus, a polarization map of G11.11-0.12 might constrain the magnetic field sufficiently well to provide an indication of the magnetic and dynamical conditions within the filament.  Improved constraints could be obtained by detailed multi-objective modeling of the filament, which would fit the intensity structure, polarization vector orientations and polarization percentages simultaneously.  Our polarization modeling also determined that the maximum polarization percentage may vary from 2-15\\% for the FP model, but is restricted to 4-15\\% for the GS model.  Thus, observations of a Type 2 pattern with a low peak polarization percentage between 2-4\\% would not be consistent with the GS model.  We have not considered non-axisymmetric magnetic field geometries in our analysis.  It is interesting to note that a filament threaded by a straight magnetic field perpendicular to the filament axis would generate a Type 1 polarization pattern that may be similar to that of a toroidal field-dominated filament, if {\\em either} the axis of the filament or the magnetic field is parallel to the plane of the sky.  However, this is not true for the general case where neither the filament axis nor the field is parallel to the sky plane, because of projection effects.  In this case, the projection of the field in the sky plane, and hence the polarization vectors, would have a contstant angle relative to the projected axis of the filament.  The symmetry of a helical field results in polarization vectors that are always parallel or perpendicular to the projected filament axis, unless the filament is bent, in which case the polarization angle changes along the filament \\citep{MWF, MFM}.  A helical field cannot produce vectors that maintain a constant angle relative to the filament axis.  Thus, it should be possible to distinguish between toroidal field-dominated helical field models and models with a non-axisymmetric orthogonal field for an ensemble of objects, although it may be difficult to distinguish between these possibilities for a single filament.  It would be interesting to include non-axisymmetric equilibrium models in future analyses.  However, no such models exist at present.  Difficult data-modeling problems are encountered frequently in astronomy and physics, which may feature large parameter spaces, noisy data, multiple constraints, and multiple modeling objectives applied to heterogenous data sets.  Our multi-objective genetic algorithm solves such problems reliably and efficiently.  We hope that our code will prove useful for a variety of problems in the physical sciences.  Our immediate goal is to build an automated system for multi-objective 2-dimensional modeling (i.e., 2 spatial map coordinates, rather than a single radial distance) of continuum and polarization maps, which would allow us to model simultaneously the intensity, polarization position angle and polarization percentage for filaments and cores.  In fact, we have already designed and tested a preliminary system that models the intensity and polarization data for one of the cores in the Orion Molecular Cloud $\\int$-shaped filament, by varying the magnetic geometry, the density structure, the temperature and dust emissivity, the position angle and inclination of the core relative to the plane of the sky.  This analysis will be the subject of future work."
    },
    "0410/astro-ph0410368_arXiv.txt": {
        "abstract": "We discuss a clumpy model of obscuring dusty tori around AGN. Cloud-cloud collisions lead to an effective viscosity and a geometrically thick accretion disk, which has the required properties of a torus. Accretion in the combined gravitational potential of central black hole and stellar cluster generates free energy, which is dissipated in collisions, and maintains the thickness of the torus.   A quantitative treatment for the torus in the prototypical Seyfert~2 nucleus of NGC~1068 together with a radiative transfer calculation for NIR re-emission from the torus is presented. ",
        "introduction": "Starting with the interpretation of spectropolarimetry of NGC~1068 by Antonucci \\& Miller \\cite{Miller83,Anton85}, the difference between type~1 and type~2 AGN is attributed to aspect-angle-dependent obscuration. The central continuum source is thought to be surrounded by a dusty molecular gas distribution, which is flattened by rotation into a torus or thick disk. The simplest unification scheme assumes that all Seyfert~2 nuclei harbor a Seyfert~1 core, so that the ratio of type 1s to 2s measures the thickness of the torus. Present estimates range form 1:4 \\cite{Maiolino95} to 1:1 \\cite{Lacy04} when selected in the mid--infrared. The torus should therefore have a thickness $H/R \\sim 1$.   For more than a decade, clumpy models for the circumnuclear torus have been discussed \\cite{Krolik88}. It was recognized early on that dust will only survive in distinct clouds, which are embedded in the inter-cloud medium of the geometrically thick torus. This idea was nonetheless not included in the following radiative transfer calculations \\cite{Pier92,Pier93,Granato94,Efstathiou95} of dust emission from these tori. Only recently, Nenkova et al. \\cite{Nenkova02} developed a scheme which uses the clumpiness of the torus for an approximative and statistical calculation of the thermal IR-emission from clouds, which are individually optically thick $\\tau_V > 40$. This approach resolved some of the problems with earlier radiative transfer calculations. The toroidal distribution of atomic and molecular material is also inferred from observed H{\\sc i} absorption and H$_2$ emission, for example in the nucleus of NGC~4151 \\cite{Mundell03}.  In two papers \\cite{Vollmer04,Beckert04} we studied the consequences of the proposed clumpiness. It turns out that a scenario with distinct, quasi-stable clouds, which experience frequent cloud-cloud collisions, is unavoidable for geometrically thick torus. In this paper we emphasis the importance of cloud collisions and present the results of radiative transfer calculations for the torus in NGC~1068. ",
        "conclusions": "We described a dynamical approach to the radial and vertical structure of a clumpy torus around an AGN following the ideas of Krolik \\& Begelman \\cite{Krolik88} and Vollmer et al. \\cite{Vollmer04}. The vertical thickness is derived from an accretion scenario for a geometrically thick torus, which is based on cloud-cloud collisions. The inferred mass accretion rates for scale heights $H > 0.5 R$ in the torus are larger than the Eddington rate for the central black hole and the model requires massive outflows from within the sublimation radius for dust. It is found that along typical lines of sight to the center (for angles $< 45^\\circ$ w.r.t. the torus midplane) between one and ten clouds on average obscure the AGN. Only a small fraction of the mass accreted through the torus eventually reaches the black hole but is sufficient to generate the ionizing luminosity. An in depth analysis of cloud collisions in the environment of an AGN is required to test our model assumptions for $\\epsilon$ and the structure of individual clouds."
    },
    "0410/astro-ph0410642_arXiv.txt": {
        "abstract": "{ We have shown that the cluster-mass reconstruction method which combines strong and weak gravitational lensing data, developed in the first paper in the series, successfully reconstructs the mass distribution of a simulated cluster. In this paper we apply the method to the ground-based high-quality multi-colour data of \\RXJ, the most X-ray luminous cluster to date. A new analysis of the cluster core on very deep, multi-colour data analysis of  VLT/FORS data reveals many more arc candidates than previously known for this cluster. The combined strong and weak lensing reconstruction confirms that the cluster is indeed very massive. If the redshift and identification of the multiple-image system as well as the redshift estimates of the source galaxies used for weak lensing are correct, we determine the enclosed cluster mass in a cylinder to $M(<360\\: {\\rm h}^{-1}\\mbox{kpc})= (1.2 \\pm 0.3) \\times 10^{15} M_{\\odot}$. In addition the reconstructed mass distribution follows the distribution found with independent methods (X-ray measurements, SZ).  With higher resolution (e.g. HST imaging data) more reliable multiple imaging information can be obtained and the reconstruction can be improved to accuracies greater than what is currently possible with weak and strong lensing techniques. ",
        "introduction": "\\label{sec:introduction}   Clusters of galaxies have been a focus of a very intense ongoing research. Especially important for many cosmological applications is a good determination of their mass. One way to obtain their masses is to use the gravitational lensing information, both from multiple image systems (strong lensing) as well as from distortions of background sources (weak lensing). Many weak and strong lensing cluster mass reconstructions have been successfully performed in the past (see e.g. \\citealt{clowe01, clowe02, gavazzi04} for examples of weak lensing and e.g. \\citealt{kneib03,smith04} for a combination of weak and strong lensing). While weak lensing mass reconstructions have an advantage in constraining the mass at much larger radii than strong lensing, one of main limitations for both strong and weak lensing is the problem of the mass-sheet degeneracy (i.e. the mass profile of the cluster can only be determined up to a constant).  In the absence of redshift information from individual sources and the lens, one can break this degeneracy only by making assumptions about the underlying potential. Different assumptions, however, can lead to discrepant results on the cluster mass.  In this work we therefore use individual redshifts of background sources to overcome this problem.  As shown in \\citet{bradac03b}, by using these and by extending the reconstruction to the inner parts of the cluster we are effectively able to break this degeneracy.   \\defcitealias{bradac04a}{Paper I} This is the second of the series of papers in which we develop and test a cluster mass reconstruction technique that combines strong and weak lensing information. In \\citet{bradac04a} (hereafter \\citetalias{bradac04a}) we describe the method in which we extend the weak lensing formalism to the inner parts of the cluster, use redshift information of the background sources and combine these with the constraints from multiply imaged systems. Using simulated data we have shown that the method is successful in reconstructing the mass distribution of a cluster, and yields an excellent agreement between the input and reconstructed mass also on scales within and beyond the Einstein radius.   Encouraged by the success of our method, we apply it to the weak and strong lensing data for the redshift $0.451$ cluster {\\RXJ} \\citep{schindler95}, the most X-ray luminous cluster known to date. Due to its record holding, this cluster has been a subject of many studies in X-ray \\citep{schindler95, schindler97,allen02,ettori04,gitti04} and optical \\citep{fischer97,sahu98,cohen02,ravindranath02}. It has also been detected through the Sunyaev-Zel'dovich effect \\citep{pointecouteau01,komatsu01,kitayama04}. Yet the mass determinations based on X-ray properties, SZ effect, velocity dispersion measurement, strong and weak lensing have all yielded discrepant results (see \\citealp{cohen02} for a summary).    For the purpose of mass reconstruction we use VLT/FORS data on a field of $3.8\\times3.8\\mbox{ arcmin}^2$ in U, B, V, R, and I bands.  We also use Ks-band data from VLT/ISAAC to obtain more reliable photometric redshift estimates.  The shape measurements for the weak lensing reconstruction is performed on two FORS bands, R and I. The strong lensing properties of this cluster are analysed. From previous data sets five arc candidates were reported \\citep{schindler95,sahu98}; using the new multi-colour data we conclude that only two possibly belong to the same multiple image system. Furthermore, we searched for additional images belonging to this system and identified a third possible member. Several new arc candidates were found as well and are presented in this work.  Particularly interesting is a very red arc candidate with two components, located at a distance of $1\\mbox{ arcmin}$ from the brightest cluster galaxy (BCG). In addition, we detect further elongated structures, some of them have been previously indicated by \\citet{lenzen04} who developed and use an automated arc searching routine.   This paper is organised as follows. In Sect.~\\ref{sc:obs} we describe the observations and give a brief outline of the data reduction process. In Sect.~\\ref{sc:RXJ_strong} we describe how we search for multiply imaged systems. In Sect.~\\ref{sc:RXJ_rec} we give the results of a combined strong and weak lensing reconstruction and the cluster luminosity measurements. We conclude in Sect.~\\ref{sc:RXJ_concl}. ",
        "conclusions": "\\label{sc:RXJ_concl}   The case of {\\RXJ} has been a cause of many puzzles in the past. Very discrepant mass estimates are given in the literature, and unfortunately this cluster is not the only case where the mass measurements have proven to be difficult. We have applied a new mass reconstruction method to deep optical data using a multiple-image system with three images selected based on their colours and redshifts. Our main conclusions are the following. \\begin{enumerate} \\item The combined strong and weak lensing mass reconstruction confirms that the most X-ray luminous cluster is indeed very massive. If the redshift and identification of the multiple-image system, as well as redshifts used in weak lensing data, are correct we estimate the enclosed cluster mass within $360\\: { h}^{-1}\\mbox{kpc}$ to $M(<360\\: {h}^{-1}\\mbox{kpc})= (1.2 \\pm 0.3) \\times 10^{15} M_{\\odot}$. \\item  The reconstruction shows a south-east mass extension, compatible with the X-ray measurements (see e.g. \\citealp{gitti04}, \\citealp{allen02}). \\item A single SIS fit to the average tangential ellipticities does not give a reliable estimate for the enclosed mass within $360\\: { h}^{-1}\\mbox{kpc}$; detailed modelling needs to be performed. \\item We have demonstrated the feasibility of breaking the mass-sheet degeneracy in practice by using shape measurements and adding the information on individual redshifts, without any assumptions regarding the cluster potential. \\end{enumerate} In addition we measured the corresponding mass-to-light ratio of the cluster within $360\\:{h}^{-1}\\mbox{kpc}$. We find that the cluster is more luminous in the rest-frame I-band than R-band, which is expected due to the presence of many old (red) elliptical galaxies in clusters. The resulting mass-to-light ratios are high, both in rest-frame I- and R-band, giving $M/L_{\\rm I} = 400 \\pm 150 M_{\\odot} / L_{\\rm I,\\odot}$ and $M/L_{\\rm R} = 550 \\pm 150 M_{\\odot} / L_{\\rm R,\\odot}$. These values are higher than typical values for clusters claimed in the literature ($\\sim 200$ for R-band). However it is difficult to compare our results with existing measurements of the mass-to-light ratios, since they are usually performed at larger radii not accessible with our data.    In the course of this research we discovered one new extremely red arc candidate (system B) at $\\sim 1\\mbox{ arcmin}$ distance from the BCG. Unfortunately its redshift can not be measured, as it is significantly detected only in the Ks band. Further arc candidates are discovered from the combined colour image, suggesting that the cluster is indeed very centrally concentrated.  In addition, the enclosed mass we obtain using the combined reconstruction also fits reasonably well the standard mass vs. X-ray luminosity relation (see \\citealp{reiprich02}), provided we assume the model to be isothermal (which for the same enclosed mass as our reconstructed model means $\\sigma \\simeq 1800\\mbox{ kms}^{-1}$) to a radius of $r_{200}$ frequently used to determine the relation.   The mass-reconstruction of {\\RXJ} can be significantly improved. Deep HST imaging would greatly help in identifying and confirming new multiple-image systems, thus allowing more detailed modelling.  In addition, not only the centre of the light for each of the arcs can be used as constraints, but also their morphology.  As mentioned in \\citetalias{bradac04a}, the reconstruction technique with adaptive grid at image positions can be used for these purposes. Further, spectroscopic redshifts need to be obtained for the multiple-image system candidates as well as for the cluster members (to obtain velocity dispersion measurements from a large sample). Deep, wide-field imaging data of this cluster will help us to improve the weak lensing constraints also at larger radii than presented here. A large number density of sources that can be used for weak lensing accessible by ACS ($\\gtrsim$ $120\\mbox{ arcmin}^{-2}$) would greatly improve the accuracy of the mass estimate and enable us to resolve substructures in the cluster. The details of the reconstruction can be used to reliably determine the cluster profile.   In conclusion, even without the best data quality that can be reached at present, we were able to perform a detailed cluster-mass reconstruction of the most X-ray luminous cluster {\\RXJ}. The method has also shown a high potential for the future. If the highest quality data is used, a combination of strong and weak lensing has proven to offer a unique tool to pin down the masses of galaxy-clusters as well as their profiles and accurately test predictions within the CDM framework."
    },
    "0410/astro-ph0410656_arXiv.txt": {
        "abstract": "We analyse oscillations observed in the X-ray light curve of the late-type star \\object{AT Mic}. The oscillations occurred during flare maximum. We interpret these oscillations as density perturbations in the flare loop. Applying various models derived for the Sun, the loop length and the magnetic field of the flare can be estimated. We find a period of 740~s, and that the models give similar results (within a factor of 2) for the loop length ($\\sim 5.4\\ 10^{10}~{\\rm cm}$) and the magnetic field ($\\sim 100~{\\rm G}$). For the first time, an oscillation of a stellar X-ray flare has been observed and results thus obtained for otherwise unobservable physical parameters. ",
        "introduction": "\\label{intro}  Oscillations in the solar corona have been observed for many years. Wavelength regimes ranging from hard X-ray right down to radio have been investigated to search for evidence of waves. Periods have been found ranging from 0.02 to 1000 s. Table 1 in \\cite*{aschwanden1999b} provides an excellent summary of the different periods that have been found, and an explanation for their existence. Most of these waves have been explained by MHD oscillations in coronal loops. \\cite*{roberts2000} has provided an excellent review of waves and oscillations in the corona.  Many of the observations of waves have been determined from variations in intensity brightness. Recently, however, a huge step forward has been achieved in solar coronal physics due to the high spatial resolution available with the Transition Region and Coronal Explorer (TRACE). The first spatial displacement oscillations have been observed in coronal loops (\\cite{aschwanden1999b}). It was suggested that these oscillations were triggered by a disturbance from the core flare site. Various MHD waves were investigated and it was found that a fast kink mode wave provides the best agreement with the observed period of 280$\\pm$30s.  One of the most exciting aspects of observing waves in this fashion is that it potentially provides us with the capability of determining the magnetic field in the corona. It is notoriously difficult to measure the magnetic field in the corona. Techniques using the near-infrared emission lines have been successful, but have poor spatial resolution. Most frequently, indirect methods are used such as the extrapolations of the coronal magnetic field from the photospheric magnetic field that can be measured using Zeeman splitting. \\cite*{nakariakov2001} made use of the flare-related spatial oscillations sometimes observed when a flare occurs, to determine the magnetic field. They assume that the oscillation is due to a global standing kink wave, and hence the magnetic field is related to the period of the oscillation, the density of the loop, and the length of the loop. In the case of \\cite*{nakariakov2001} they found the magnetic field to be 13$\\pm$9G. The errors on this are large due to errors on the determination of density, loop length and period. \\cite*{schrijver2000} suggest that the oscillations are due to a sudden displacement of the magnetic field at the surface which causes an oscillatory relaxation of the field. Recent work by \\cite*{nakariakov2004} suggests that the oscillations are second standing harmonics of acoustic waves tied to the loop length. The model gives a relationship between the oscillation period, the loop length and the average plasma temperature, all of which can be independently observed on the Sun.  Waves are observed across the electromagnetic spectrum on the Sun. However, observations on other stars are rare. One reason for this is that on the Sun, we have the spatial resolution to zoom in on small regions. For example, the transverse TRACE loops that have been observed are away from the main flare site, and hence have lower emission measure. Analysing an X-ray light-curve, for example, would be unlikely to provide evidence of waves if the emission was averaged over the whole Sun. Oscillations have been observed for optical stellar flares (e.g., \\cite{mathioudakis2003} and \\cite{mullan1992}). \\cite*{mathioudakis2003} found a period of 220~s in the decay phase of a white-light flare on the RS CVn binary \\object{II Peg}. Using this they determine a magnetic field of 1200~G. \\cite*{mullan1992} again found oscillations in X-ray active red dwarfs. They compared their observations to the possibility that they were observing p-modes. On the Sun p-modes give rise to 5-minute oscillations on the surface. They concluded that p-modes could not produce the amplitudes they observed and that it must be coronal oscillations, which have, until now, not been observed.  In this work we observe for the first time an oscillation in the corona of \\object{AT Mic}. We measure the periodicity, and hence determine the magnetic field. This is the first observation of an X-ray oscillation during a stellar flare. ",
        "conclusions": "We have used three different approaches to determine the loop length of the \\object{AT Mic} flare and two different approaches for the magnetic field. The loop length derived in Sect.~\\ref{oscill} from the flare oscillation is the largest with $L = 5.4\\cdot 10^{10}~{\\rm cm}$, while the loop length derived from radiative cooling times (Sect.~\\ref{radcool}) is somewhat smaller by roughly a factor of 2, $L=2.3~10^{10}~{\\rm cm}$. The loop length derived from pressure balance (Sect.~\\ref{pressbal}) is similar to the latter one, $L=(2.8\\pm1.7)\\ 10^{10}~{\\rm cm}$. The loop length derived by assuming that the variations seen are due to oscillations is a factor of 2 different to the loop length determined by more usual methods. Considering the assumptions that have to be made, this is good agreement, and gives us confidence that we are, for the first time, observing a stellar coronal loop oscillating. The magnetic field ranges from 50~G (radiative cooling) to 115$\\pm$56~G (oscillation), with the value from pressure balance in between with 70$\\pm$40~G. Again, these three values are all consistent.  This was the first time that an oscillation during flare peak was observed in X-rays in a stellar flare and flare loop length and magnetic field derived from it. The values are consistent with other flare models."
    },
    "0410/astro-ph0410183_arXiv.txt": {
        "abstract": "We study the influence of general lens and source velocities on the gravitational deflection of light by single and two-point-mass microlenses with general axis orientation. We demonstrate that in all cases the lens equation preserves its form exactly. However, its parameters -- the Einstein radius and the binary-lens separation -- are influenced by the lens velocity. In Galactic microlensing settings the velocity mainly affects the inferred separation for wide binary-star or star+planet microlenses oriented close to the line of sight. We briefly discuss the case of lenses moving with highly relativistic velocities. ",
        "introduction": "The time-dependence of Galactic gravitational microlensing events \\citep{pacz96} is given by the relative motion of the source of light, the lens, and the observer near perfect alignment. These events have been generally successfully analyzed using a quasi-static approach, assuming that light is deflected at each instant in a static source-lens-observer configuration. This approach is well justified -- the lens and source velocities relative to the observer are of the order of a couple hundred $km\\,s^{-1}$, i.e., $\\sim 10^{-3}$ of the speed of light $c$.  Nevertheless, with the introduction of image subtraction techniques \\citep{allu98} the accuracy of measured microlensing light curves has increased in ideal cases to sub-percent levels. In addition there are prospects of high-precision observations of the astrometric microlensing effect (see \\citealt{bsv98}; \\citealt{han01} for the single- and binary-lens cases, respectively) by space-based interferometers such as the Space Interferometry Mission (SIM)\\footnote{\\url{http://sim.jpl.nasa.gov}}. Given these developments, one can expect that corrections to the light curve or the angular image geometry down to the order $\\sim10^{-3}$ might be detectable. It is therefore interesting to avoid the quasi-static approach and investigate the effect of general lens and source velocities on light deflection.  Previous theoretical research has mostly concentrated on the effect of a relative lens to observer velocity on the deflection angle by a single lens. \\cite{pybi93} derived the first-order effect for a velocity of arbitrary direction, showing that the deflection angle increased for lenses moving away from the observer and decreased for those moving towards the observer. The same result was later confirmed by \\cite{frit03}. Recently \\cite{wusp04} presented results for an arbitrarily large but purely radial lens velocity. The most general theoretical result can be found in the detailed treatise by \\cite{kosch99}, which includes a derivation of the deflection angle for an ensemble of lenses of general velocities.  In this paper we study light deflection for general lens and source velocities (arbitrarily oriented and arbitrarily large) specifically for single and two-point-mass lens microlensing events. In \\S2 we derive the general lens equation and demonstrate the velocity dependence of its parameters. In \\S3 we explore the magnitude of the velocity effects for low velocities and illustrate the highly relativistic limit. We conclude in \\S4 by discussing observational aspects and the appropriateness of some of the assumptions. ",
        "conclusions": "In \\S3.1 we have demonstrated the low-velocity effects for different binary-lens orientations and intrinsic separations. Nevertheless, from the observational perspective we must take into account the fact that not all binary-lens configurations can lead to microlensing events detectable as two-point-mass lens events. While this depends on the exact source trajectory, from a statistical point of view the Einstein-radius scaled lens separation $d$ plays the main role. If the separation is too small or too large, most events would appear as single-lens events.  For the event to be detectable as a two-point-mass lens event the lens separation has to fulfil $d(0)\\in(d_{min},d_{max})$. This leads to limits on the binary-lens axis orientation \\beq \\cos\\zeta\\in(d_{min},d_{max})\\times\\left(\\frac{R}{4\\,AU}\\right)^{-1} \\left(\\frac{D_L}{4\\,kpc}\\right)\\left(\\frac{\\theta_E(0)}{1\\,mas}\\right)\\;. \\label{eq:cosine} \\eeq The values of the binary separations from the 21 binary events published by the MACHO team \\citep{alc00} range from 0.421 to 2.077 with an outlier value of 7.454 (a possible binary-source event). The 18 events detected by OGLE-II \\citep{jaro02} have values from 0.355 to 2.917, and the values for the 15 events detected by OGLE-III \\citep{jaro04} range from 0.352 to 3.457.  To obtain a rough estimate we set $d_{min}=0.3$ and $d_{max}=4$. Keeping $D_L$ and $\\theta_E$ fixed at the values used in equation~(\\ref{eq:cosine}), for an intrinsic binary separation $R=15 AU$ we get limits on the axis angle $\\zeta\\in(0^\\circ,85^\\circ)$, for $R=60 AU$ we get $\\zeta\\in(75^\\circ,89^\\circ)$, and for $R=240 AU$ we get $\\zeta\\in(86^\\circ,89.7^\\circ)$. In the first case we have a wide range of possible orientations -- however, if we approach $90^\\circ$ where the velocity effects are strongest, it would be difficult to detect such an event as a binary-lens event. The two cases with higher intrinsic separations demonstrate that such lenses can have stronger velocity effects while being detectable as binary events, albeit for a narrower range of axis orientations.  In this work we concentrated on the center-of-mass velocity of the binary lens and neglected its orbital velocity. However, for binaries or star+planet systems with semi-major axis $\\ll 1 AU$ the assumption of small orbital velocity breaks down. The results of this paper indicate that the velocity effects are proportional to the intrinsic lens separation, being caused by the lens motion during the passage of the light ray in its vicinity. While the orbital velocity grows as $R^{-1/2}$ with decreasing $R$, its product with the separation decreases as $R^{1/2}$. We conclude that the effects of orbital lens velocity on light deflection are not significant even in binary-lens systems closer than $1 AU$."
    },
    "0410/astro-ph0410240_arXiv.txt": {
        "abstract": "We have obtained {\\it HST} WFPC2 observations of three very luminous but morphologically different giant \\hii\\ regions (GHRs) in M101, NGC\\,5461, NGC\\,5462 and NGC\\,5471, in order to study cluster formation in GHRs. Images obtained in the F547M and F675W bands are used to identify cluster candidates and for photometric measurements, and images in the F656N band are used to show ionized interstellar gas. The measured colors and magnitudes are compared with the evolutionary tracks generated by the Starburst99 and Bruzual \\& Charlot population synthesis models to determine the ages and masses of the cluster candidates that are more luminous than $M_{\\rm F547M}=-9.0$. The brightest clusters detected in the PC images are measured and found to have effective radii of 0.7--2.9 pc.  NGC\\,5461 is dominated by a very luminous core, and has been suggested to host a super-star cluster (SSC). Our observations show that it contains three R136-class clusters superposed on a bright stellar background in a small region. This tight group of clusters may dynamically evolve into an SSC in the future, and may appear unresolved and be identified as an SSC at large distances, but at present NGC\\,5461 contains no SSCs. NGC\\,5462 consists of loosely distributed \\hii\\ regions and clusters without a prominent core. It has the largest number of cluster candidates among the three GHRs studied, but most of them are faint and older than 10 Myr. NGC\\,5471 has multiple bright \\hii\\ regions, and contains a large number of faint clusters younger than 5 Myr. Two of the clusters in NGC\\,5471 are older than R136, but just as luminous; they may be the most massive clusters in the three GHRs studied.  The fraction of stars formed in massive clusters has been estimated from the clusters' contribution to the total stellar continuum emission and from a comparison between the ionizing power of the clusters and the ionizing requirement of the associated \\hii\\ regions. Both estimates show that $\\lesssim$ 50\\% of massive stars are formed in massive clusters; consequently, the \\ha\\ luminosity of an \\hii\\ region does not provide a sufficient condition for the existence of SSCs. The cluster luminosity functions (LFs) of the three GHRs show different slopes. NGC\\,5462 has the steepest cluster LF and the most loosely distributed interstellar gas, qualitatively consistent with the hypothesis that massive clusters are formed in high-pressure interstellar environments. The combined cluster LF of the three GHRs has a slope similar to the universal cluster LFs seen in starburst galaxies and non-starburst spiral galaxies. ",
        "introduction": "Giant \\hii\\ regions (GHRs) are sites of intense massive star formation. Their \\ha\\ luminosities, $10^{39}$--$10^{41}$ ergs~s$^{-1}$ \\citep{Ke84}, require an ionizing power equivalent to that of 24--2400 O5V stars \\citep{SdK97}. With such high concentrations of massive stars, GHRs provide an excellent laboratory to study the modes of massive star formation, and in particular to probe whether they are birthplaces of globular clusters \\citep{KC88}.  In the two nearest GHRs, 30 Dor in the Large Magellanic Cloud (LMC) and NGC\\,604 in M33, where stellar contents are well resolved, two distinct types of stellar groupings have been observed: 30 Dor is dominated by one central massive cluster R136 \\citep{Hu95,WB97}, while NGC\\,604 contains multiple OB associations spreading over a large area \\citep{Hu96}. Evidently, not all GHRs contain massive compact clusters; what physical environments give rise to the various cluster morphologies is currently under investigation.  One obvious way to elucidate this issue is to carry out detailed examinations of relatively nearby clusters and their environments. \\citet{MA01} studied 27 nearby ($<$ 5 Mpc) clusters of varying morphological types, primarily classifying clusters based on their core and halo sizes. He suggests that compactness of clusters is predominantly related to the central density of the progenitor giant molecular cloud, i.e., extremely high pressure environments may be required to form massive compact clusters. However, this scenario has not been observationally tested; we do not know the pressures and densities of the giant molecular cloud in which optically visible clusters were formed. Examining clusters in a range of present-day environments may help us gain insight into their properties and relationship to their natal interstellar medium.  Of all massive compact clusters, the most impressive ones are the super-star clusters (SSCs) with masses of 10$^5$--10$^6$ M$_\\odot$. SSCs are frequently observed in galaxy mergers and starburst galaxies \\citep[][and references therein]{Wh03}, and they are believed to be preferentially formed in high-pressure interstellar conditions \\citep[e.g.,][]{EE97}. However, some GHRs in non-interacting, late-type spiral galaxies may also host SSCs, especially those GHRs that are several times as luminous as 30 Dor and require ionizing powers rivaling those of young SSCs \\citep[e.g.,][]{LP01}. It is thus intriguing to examine the cluster content of such GHRs and investigate whether these relatively quiescent environments can also produce SSCs.  The giant spiral galaxy M101 contains a large number of very luminous GHRs whose stellar content can be resolved and studied with {\\it Hubble Space Telescope} ({\\it HST}) images. We have therefore obtained {\\it HST} observations of three M101 GHRs with different morphologies and galactic locations: NGC\\,5461, NGC\\,5462, and NGC\\,5471. The locations of these GHRs in M101 are marked on a Second Palomar Observatory Sky Survey (POSS-II) red image in Figure~\\ref{fig:m101}. The properties of these GHRs are summarized in Table~\\ref{tbl:GHRs}; for comparison, 30 Dor is also included in this table. We have used the {\\it HST} continuum and \\ha\\ images of these three GHRs to carry out a detailed photometric study of their clusters. This paper reports our observations (\\S2) and methodology (\\S3), describes the cluster content in each GHR (\\S4), discusses cluster properties and their implications in studying massive star formation and cluster formation (\\S5), and summarizes our results (\\S6). ",
        "conclusions": "\\subsection{Nature of Faint Cluster Candidates in GHRs}  A large number of cluster candidates have been identified in the three M101 GHRs, but the nature of the faintest objects is uncertain. It is possible that some of these faint cluster candidates consist of multiple OB associations as observed in the nearby GHR NGC\\,604 \\citep{Hu96} and some are simply luminous supergiants frequently seen near high concentrations of massive stars, such as 30 Dor \\citep{WB97}. We have therefore simulated WFPC2 images of 30 Dor and NGC\\,604 at a distance of 7.2 Mpc and searched for ``clusters'' using the same criteria as we did for cluster candidates in the M101 GHRs. The spurious clusters in 30 Dor and NGC\\,604 can be identified because their resolved stellar contents are known. The real and spurious clusters in these two GHRs can then be compared with the cluster candidates in M101 to better assess the nature of the latter.  To simulate a WFPC2 image of 30 Dor at 7.2 Mpc, we have used a green continuum ($\\lambda_c$ = 5130 \\AA, $\\Delta \\lambda$ = 155 \\AA) image from the Magellanic Clouds Emission-Line Survey \\citep[MCELS,][]{Smith99}, and binned the data to 3.5 pc per pixel. The resultant image is displayed in Figure~\\ref{fig:ghrs}a. Using the same identification criteria for clusters in M101 GHRs, the two brightest objects in 30 Dor will be selected as clusters: R136 and R131. While R136 is a bona fide cluster, R131 ($=$ HD\\,269902) is an A0 supergiant with $V = 10.0$ (corresponding to $M_V + A_V = -8.5$) and hence a spurious cluster. The mis-identification of R131 as a cluster bolsters our choice of a luminosity cutoff of $M_{\\rm F547M} = -9.0$ for cluster mass estimates (\\S3.2), as the fainter cluster candidates may be single supergiants. It is interesting to note that two other concentrations of stars in 30 Dor are not identified as clusters: the Hodge 301 cluster \\citep{H88} and the OB association LH\\,99 \\citep{LH70}. The Hodge 301 cluster, with an age of $\\sim$ 20--25 Myr and a mass of a few $10^3$ M$_{\\odot}$ \\citep{GC00}, is too faint to meet our cluster identification criteria. LH\\,99, on the other hand, is too distributed to mimic a cluster.  A WFPC2 image of NGC\\,604 at 7.2 Mpc is simulated with its archival WFPC2 F547M image. Adopting a distance of 0.84 Mpc to M33 \\citep{FWM91}, the data are binned to 3.5 pc per pixel, and the resultant image is displayed in Figure~\\ref{fig:ghrs}b. The OB associations in NGC\\,604 appear as small concentrations on top of an irregular stellar background. The four brightest concentrations have $M_{\\rm F547M}=-8.0$ to $-9.0$, luminous enough to meet our identification criteria for M101 clusters. However, most of the concentrations have irregular shapes, and the brightest one does not even have an obvious boundary. Therefore, the NGC\\,604-type OB associations may mimic faint clusters at best, with luminosities rivaled by those of supergiant stars.  For a direct comparison between M101 GHRs and the simulated 30 Dor and NGC\\,604 at 7.2 Mpc, Figure~\\ref{fig:ghrs} displays their images in the same spatial and intensity scales\\footnote{The green image of 30 Dor was taken in a similar but not identical wavelength band.  The intensity scale of the 30 Dor image is selected to match that of the F547M images as much as possible so that objects with similar magnitudes appear similar in both F547M and green band images.} over a 350 pc $\\times$ 350 pc field-of-view. It is immediately clear that the brightest cluster candidates in M101 GHRs are more luminous than R136 and are most likely bona fide clusters. The nature of the cluster candidates fainter than $M_{\\rm F547M} = -9.0$, marked in Figure~\\ref{fig:ghrs}, is less obvious. Some faint cluster candidates may be blue supergiants because they have sharp images and appear isolated, and the blue color excludes the possibility of post-core-collapse clusters; examples of these include NGC\\,5461-7, NGC\\,5462-20, and perhaps NGC\\,5471-13 and 14. Some faint cluster candidates may be OB associations because they appear extended without a sharp boundary; examples of these include NGC\\,5461-5, NGC\\,5462-13 and 15, and NGC\\,5471-10, 11, and 17.   \\subsection{The Fraction of Massive Stars Formed in Clusters}  It has been suggested that massive stars form preferentially in associations and clusters \\citep[][and references therein]{SPH00}. Our {\\it HST} WFPC2 images of the M101 GHRs show that several young R136-class clusters are superposed on discrete regions of unresolved stellar emission, e.g., clusters \\#8, \\#9, and \\#10 at the core of NGC\\,5461 and cluster \\#16 at the core of NGC\\,5471A. The unresolved stellar backgrounds are most likely star clouds that contain field stars and loosely assembled associations. The similarities in locations and colors suggest that the clusters and the background star clouds are formed from the same episode of star formation. It is then interesting to determine the fraction of massive stars that are formed in R136-class clusters in these regions. We have used two different methods to determine this fraction: one is based on the contribution of cluster light to the total light, and the other is based on a comparison between the ionization flux expected from the clusters and the ionizing flux required by the surrounding \\hii\\ region.  We have selected four regions for this analysis: two in NGC\\,5461, one in NGC\\,5462, and one in NGC\\,5471. These regions are listed in Table~7, and their close-up F547M and \\ha\\ images are presented in Figure~\\ref{fig:4regions}. The F547M images show that the clusters in these four regions are all superposed on discrete bright diffuse stellar backgrounds, and the \\ha\\ images show that all are at the cores of bright compact \\hii\\ regions. We have measured the total light from these four regions in both F547M and F675W$'$ bands using the apertures marked on the F547M images in Figure~\\ref{fig:4regions} and described in column 3 of Table~7. The background, determined from the median of an annular region outside the \\hii\\ region, has been subtracted, although it contributes to only 1--2\\% of the total light. For the clusters, we have applied aperture corrections of $\\sim -0.2$ mag to our photometric measurements made with a 0\\farcs2-radius APPHOT aperture to account for the missing light. The cluster-to-total light ratios,  $L_{\\rm cluster}/L_{\\rm total}$ given in columns 4 and 5 of Table~7, are in the range of 0.25--0.5; the uncertainties are dominated by the photometry and aperture corrections of the clusters, as they are superposed on bright local stellar background. The light ratios are slightly larger in the F547M band than in the F675W band, because the clusters are 0.1--0.2 mag bluer than their diffuse stellar background. This color difference can be caused by an age difference of a few Myr, assuming that the clusters and the underlying star clouds have the same initial mass function. As the precise ages are unknown, we cannot model the star clouds to determine their masses; therefore, the cluster-to-total light ratio can be considered only as an approximation of the fraction of massive stars formed in clusters. In the four regions we analyzed, about 25--50\\% of the massive stars are formed in R136-class clusters.  The \\ha\\ images in Figure~\\ref{fig:4regions} show that the \\hii\\ regions around the clusters have rather well-defined boundaries where the surface brightness drops off sharply. Such morphology suggests that the \\hii\\ regions are likely ionization-bounded, or optically thick to ionizing radiation. We have measured the \\ha\\ fluxes\\footnote{Owing to the $\\sim300$ \\kms\\ redshift of M101, two corrections need to be considered. First, the filter transmission of the red-shifted \\ha\\ line is $\\sim93$\\% of the peak transmission, thus the extracted \\ha\\ flux should be multiplied by a correction factor of 1.07. Second, the \\nii$\\lambda$6548 line is red-shifted into the \\ha\\ bandpass at $\\sim91$\\% of the peak transmission and needs to be removed.  The \\nii\\ contamination, estimated from the \\nii/\\ha\\ ratios reported by \\citet{KG96}, amounts to 1--3\\% of the \\ha\\ flux in most cases.} of these four regions using the apertures marked on the \\ha\\ images in Figure~\\ref{fig:4regions} and described in column 6 of Table~7. The continuum-subtracted \\ha\\ images are used for the flux measurements, and extinction corrections are made. Assuming a 10$^4$ K optically thick \\hii\\ region, the derived \\ha\\ luminosity, $L_{{\\rm H}\\alpha}$, can be used to determine the required ionizing luminosity, $Q({\\rm H}^0)$, through the relation \\begin{equation} Q({\\rm H}^0) = 7.4\\times10^{11}~L_{{\\rm H}\\alpha}~{\\rm photons~s}^{-1}, \\end{equation} where $L_{{\\rm H}\\alpha}$ is in units of ergs~s$^{-1}$. The resultant ionizing luminosities ($Q_{\\rm HII}$) of the four regions are given in column 7 of Table~7.  The ionizing luminosity expected from the clusters at different ages can be calculated using the Starburst99 models. The ages of the clusters in the four selected regions are $<$ 5 Myr, as indicated by the associated bright \\hii\\ regions. Unfortunately, during the first 5 Myr a cluster's ionizing luminosity decreases rapidly, dropping from the maximum at $\\sim$ 1 Myr to a factor of 5--7 lower at 5 Myr \\citep{Le99}; therefore, the uncertainty in cluster age directly propagates into the uncertainty in ionizing luminosity of a cluster. We have adopted a cluster age of 3 Myr and calculated the expected ionizing luminosities of the clusters ($Q_{\\rm cluster}$) and their ratios to those required by the surrounding \\hii\\ regions. These results are given in columns 8 and 9 of Table 7. These ratios, 0.2--0.6, can be viewed as a very crude approximation of the fractions of massive stars formed in clusters. An interesting corollary of this result is that the ionizing luminosity, or \\ha\\ luminosity, of an \\hii\\ region is not a sufficient diagnostic for the existence of SSCs because the majority of stars may reside outside clusters.  In the four regions of star formation we considered, the fraction of massive stars in clusters estimated from the cluster-to-total light ratio is, within the uncertainties, consistent with that estimated from the cluster-to-\\hii-region ionizing luminosity ratio -- no more than about half of the massive stars are formed in the R136-class clusters. On the other hand, some R136-class clusters, such as \\#4 and \\#5 in NGC\\,5471D, are not superposed on a bright stellar background, and constitute the dominant components in their associated episode of star formation. The fraction of stars formed in clusters must cover a range, which varies according to the physical conditions of star formation.   \\subsection{Interstellar Environments of the GHRs}  We examine the distribution of interstellar molecular clouds and \\ion{H}{1} gas in the three M101 GHRs in order to gain insight on the cluster formation process. Molecular CO observations of these GHRs have been made by \\citet{GC99} using both single-dish telescopes and interferometers. NGC\\,5461 has the strongest CO emission among the three GHRs, and its CO peaks appear concentrated toward the peaks of the \\ha\\ emission \\citep[see Figure~4 in][]{GC99}, where the massive clusters are located. The CO emission toward NGC\\,5462 is detected in single-dish observations but not in interferometric observations, indicating that the molecular gas is distributed over a scale larger than the synthesized beam, $\\sim$3$''$. No CO emission is detected in NGC\\,5471, which may be attributed to its low metallicity. In the two GHRs where CO is detected, the distribution of molecular clouds is similar to that of clusters: concentrated in NGC\\,5461 and distributed in NGC\\,5462.  The distribution of \\hi\\ gas in M101 and its relation with ionized gas have been reported by \\citet[][see their Figures~1--2]{Setal00}. Their \\hi\\ map shows that NGC\\,5461 and NGC\\,5462 are in the same spiral arm. Assuming a trailing arm, the offsets of \\hi\\ ridge downstream from the stars in NGC\\,5461 and NGC\\,5462 are consistent with the expectations of star formation triggered by density waves \\citep{Ro69}. NGC\\,5471 has a concentration of \\hi\\ gas but the large-scale distribution of \\hi\\ show a complex inter-arm structure, which may have resulted from tidal interactions during the last 10$^9$ yr \\citep[][and references therein]{Wetal97}. As \\hi\\ gas can be produced by photodissociation of the natal molecular clouds \\citep{AAT85,AAT86,Setal00}, converted to \\hii\\ by photoionization, and dispersed by fast stellar winds and supernova explosions, the distribution of \\hi\\ does not provide adequately pertinent information about the cluster formation.  To study the physical conditions for cluster formation, it is necessary to examine the interstellar environment of the youngest clusters before the interstellar conditions have been altered by stellar energy feedback. The embedded young clusters that are observable in the infrared but not yet in the optical wavelengths \\citep[e.g.,][]{KJ99,TBH00,JK03} provide promising locations to study the physical conditions of cluster formation.  \\subsection{Cluster Luminosity Function}  The luminosity functions (LFs) of young compact clusters have been studied in various types of galaxies with different star formation rates as a means to gain insight into the cluster formation process. To first order, the measured LFs for young compact cluster systems in merging or starburst galaxies are remarkably universal, and can be approximated by a power-law of the form $dN(L)/dL \\propto L^{\\alpha}$, with the exponent $\\alpha \\approx -2 \\pm 0.2$ \\citep[see][and references therein]{Wh03}. The cluster LFs for a sample of nearby non-starburst spiral galaxies also show similar $\\alpha$, $-$2.0 to $-$2.4 \\citep{La02}. It has been suggested that this roughly universal LF is the result of fractal structure in turbulent gas \\citep{EE97}. Since the three M101 GHRs show different age and spatial distributions of clusters, it is interesting to intercompare their cluster LFs and see if they also follow the universal cluster LFs.  The LFs of clusters in NGC\\,5461, NGC\\,5462, and NGC\\,5471 are presented in Figure~\\ref{fig:lum_func}. These LFs are constructed using raw $M_{\\rm F547M}$ (without extinction correction). No completeness correction to the LFs is needed because the cutoff of our sample at the faint end, $m_{\\rm F547M} \\leq 21.3$, is much brighter than the detection limit, $m_{\\rm F547M} \\sim 25.5$. Note however that the two faintest bins, $M_{\\rm F547M} = -8.0$ to $-$9.0, should be viewed with caution, as some of the ``clusters'' may be spurious as discussed in \\S5.1. The number of clusters in each GHR is modest, so we have also constructed a combined cluster LF of the three GHRs, shown in the bottom panel of Figure~\\ref{fig:lum_func}. We have carried out linear least-squares fits to the logarithmic LFs of clusters for the three GHRs individually and combined. The logarithmic LFs of clusters in the individual GHRs appear to have different slopes; the best-fit slopes for NGC\\,5461, NGC\\,5462, and NGC\\,5471 are $-1.5\\pm0.3$, $-3.0\\pm0.2$, and $-1.9\\pm0.4$, respectively. The logarithmic LF of clusters in all three GHRs has a best-fit slope of $-2.3\\pm0.1$.  Compared with the universal cluster LFs, NGC\\,5461 and NGC\\,5471 are on the flatter side, and NGC\\,5462 is on the steeper side. The small difference between NGC\\,5461 and NGC\\,5471 is not statistically significant, as each has only a small number of clusters and the numbers of clusters in the brightest bins are only 1--2. On the other hand, it may be statistically significant that NGC\\,5462 has a much steeper LF than NGC\\,5461 and NGC\\,5471, or NGC\\,5462 has a larger proportion of low-mass clusters. It is possible that the slope of a cluster LF varies according to the interstellar environment at the time when clusters were formed. As discussed in \\S5.3, the current molecular environments of NGC\\,5461 and NGC\\,5462 are quite different, with molecular CO highly concentrated in NGC\\,5461 and diffuse in NGC\\,5462. If the current environments reflect the conditions when the clusters were formed, the clusters in NGC\\,5462 would have been formed in a lower-pressure, lower-concentration interstellar environment. The association between a steep cluster LF and a low-pressure, low-concentration star formation condition is consistent with the previously suggested hypothesis that massive clusters are formed in high-pressure, high-concentration molecular clouds.  The cluster LF for NGC\\,5461, NGC\\,5462, and NGC\\,5471 combined has an $\\alpha$ within the range of the universal value, $-$2 to $-$2.4. If clusters in spiral or starburst galaxies are formed under conditions similar to NGC\\,5461 and NGC\\,5462 in random proportions, the cluster LFs should show a larger range of $\\alpha$. The scarcity of cluster LFs with $\\alpha$ steeper than $-$2.5 might be an observational effect to some extent. Surveys of clusters in galaxies preferentially detect the most luminous clusters that are likely formed in high-pressure environments and follow an LF similar to those of NGC\\,5461 and NGC\\,5471; therefore, the cluster LF of a galaxy would be biased toward $\\alpha = -2$.    \\subsection{Evolutionary Aspects of the Clusters}  The cluster mass, age, and size distribution of a cluster system may be used to investigate the dynamic evolution of clusters. Recent studies of rich clusters in the LMC have shown that the spread in core radius increases with cluster age, suggesting that all clusters were formed with small core radii but subsequently some experienced core expansion while others did not \\citep{EFL89,MG03}. It would be interesting to examine the clusters in M101 GHRs to see whether they follow the same core radius-age relation.  The mass, age, and core radius of the LMC clusters have been derived by \\citet{MG03}, using the \\citet{FRV97} population synthesis code, the \\citet{KTG93} IMF slope, and a stellar mass range of 0.1--120 M$_\\odot$. To compare the M101 clusters with the LMC clusters, we have followed the \\citet{MG03} method and re-estimated the masses of M101 clusters, using the same code, IMF slope, and stellar mass range. The new cluster mass estimates are $\\sim$ 2--3 times as high as the cluster masses estimated earlier in this paper, owing to the addition of stars in the mass range of 0.1--1 M$_\\odot$. The M101 clusters are not sufficiently resolved for measurements of their core radii ($r_{\\rm c}$); therefore, we have adopted the relation $r_{\\rm c} \\sim 0.35 R_{\\rm eff}$ derived from a King profile with a concentration parameter of $c=1.5$, a median value for LMC clusters \\citep{EFF87}. The core radii of M101 clusters thus estimated are 0.25--1 pc, and may be uncertain by up to a factor of 2, if its concentration parameter spans the same range as that in the LMC clusters.  To compare the M101 clusters with the LMC clusters, we present a 3-D diagram of cluster mass, age, and core radius in Figure~\\ref{fig:3d-1}. The data of the LMC clusters are adopted from \\citet{MG03}. The M101 clusters are plotted in open rhombuses, while the LMC clusters are plotted in filled ellipses, with the R136 cluster in a larger ellipse for easy identification. The M101 cluster masses were estimated using 7-pc-radius apertures, so we repeated Mackey \\& Gilmore's derivation of the mass of R136 using this larger aperture and obtained a mass that is 40\\% higher, shown as a large open ellipse in Figure~\\ref{fig:3d-1}. While the LMC clusters span a large age range and show an increasing spread in core radii with the cluster age, the M101 clusters we analyzed are all young and small, sharing a similar parameter space with the R136 cluster. Among the small number of M101 clusters with size measurements, the younger clusters are generally smaller, but the accuracy of the size measurements is too limited by the linear resolution for definitive conclusions.   Finally, we discuss the disruption time of our M101 clusters and investigate whether dynamical evolution has caused a significant mass change in these clusters. The disruption time can be assessed empirically from the break in slopes of the logarithmic age distribution of clusters \\citep{BL03}, or from comparisons with $N$-body simulations \\citep{BM03}. Recent studies using the empirical method have shown that the disruption time for a 10$^4$ M$_\\odot$ cluster varies greatly among galaxies, from 10$^7$ to 10$^{10}$ yr, with the shortest being $\\sim$30--40 Myr in M82 and at 1--3 kpc from the nucleus of M51 \\citep{BL03,dG03}. Our M101 clusters in GHRs are apparently younger than these disruption timescales. Thus we determine the disruption time, $T_{\\rm dis}$, following the simulations of \\citet{BM03}: \\begin{equation} \\frac{T_{\\rm dis}}{\\rm Myr} = \\beta~\\left[\\frac{N}{{\\rm ln}(\\gamma N)} \\right]^{x} \\frac{R_{\\rm G}}{\\rm kpc} \\left( \\frac{V_{\\rm G}}{220~ {\\rm km~s}^{-1}} \\right) ^{-1} \\end{equation} where $R_{\\rm G}$ is the galactocentric radius, $V_{\\rm G}$ is the circular velocity in a galaxy, $N$ is the number of stars in a cluster, $\\beta \\sim$ 1--2, $\\gamma$ = 0.02, and $x \\sim$ 0.8. For a cluster with a Salpeter mass function and a stellar mass range of 1--100 M$_\\odot$, the average mass of a star is 3.09 M$_\\odot$ and the number of stars is $N$ = cluster mass/(3.09 M$_\\odot$). NGC\\,5461 and NGC\\,5462 have $R_{\\rm G} \\sim 10$ kpc and $V_{\\rm G}$ = 185 km~s$^{-1}$, and NGC\\,5471 has $R_{\\rm G} \\sim 25$ kpc and $V_{\\rm G}$ = 195 km~s$^{-1}$ \\citep{RR73}. A 10$^4$ M$_\\odot$ cluster in NGC\\,5461/NGC\\,5462 or NGC\\,5471 would have disruption times of 2.4--4.8 $\\times10^9$ and 5.8--11.6 $\\times10^9$ yr, respectively. The real disruption time must be shorter because the above estimates do not take into account processes that are important in disrupting clusters in GHRs, i.e., interactions with other clusters and with giant molecular clouds. In cases where massive clusters are concentrated in a small volume, such as the core of NGC\\,5461, cluster merger is a more important dynamic process and operates in a much shorter timescale than tidal disruption \\citep{BBV03}. Future simulations of dynamical evolution of clusters in GHRs using realistic conditions are needed."
    },
    "0410/astro-ph0410289_arXiv.txt": {
        "abstract": "We study analytically a gravitational lens due to a deformed star, which is modeled by using a monopole and a quadrupole moment. Positions of the images are discussed for a source on the principal axis. We present explicit expressions for the lens equation for this gravitational lens as a single real tenth-order algebraic equation. Furthermore, we compute an expression for the caustics as a discriminant for the polynomial. Another simple parametric representation of the caustics is also presented in a more tractable form. A simple expression for the critical curves is obtained to clarify a topological feature of the critical curves; the curves are simply connected if and only if the distortion is sufficiently large. ",
        "introduction": "Some recent highly accurate numerical simulations of rotating compact objects such as a neutron star elucidate many features of compact objects such as an equilibrium configuration and its stability (e.g., Shibata et al. 2002, Stergioulas et al. 2004, Zanotti and Rezzolla 2002, Zdunik et al. 2004). The deformation of their configurations causes a non-spherical distortion of the surrounding spacetime, which has not been observed yet. Probing this warp of the spacetime provides a strong verification of the theory of the general relativity (Will 1993) and a deep insight into the relevant physics especially about the nuclear matter (e.g., Shapiro and Teukolsky 1983), because the deformation depends not only on the gravity but also on the equation of state of the internal matter.  One possibility is to study effects of the deformation on the light propagation near the compact object. Before doing numerical studies in detail, it is important to find some physical properties by using a simple analytic model. One reason is that one needs many light ray tracings in numerical investigations, and consequently it is difficult to clarify quantitatively the parameter dependence.  We perform the multipole expansion of the external gravitational field, which can be assumed to be weak along the line of sight even if the internal gravity is strong. Namely, one can consider only the linearized potential. This does not mean that our system is linear, because the lens equation is non-linearly coupled with respect to the image positions as shown in the next section.  The quadrupole moment depends on the rotation speed and the equation of the state for the internal matter. Many years ago some comprehensive studies had been done for equilibrium configurations of rotating stars in the Newtonian gravity, some of which are summarized in a famous textbook (Chandrasekhar 1987). Numerical investigations on this issue in the general relativity are still going on as mentioned in the beginning of this section (e.g., Shibata et al. 2002, Stergioulas et al. 2004, Zanotti and Rezzolla 2002, Zdunik et al. 2004). In an extreme case called as rapid rotation, the quadrupole deviation of order of $10^{-3}$ may be possible even for a neutron star. These results show that the deformation is dominated by the quadrupole moment and higher moments such as the octapole moment are much smaller. In addition, a contribution of the octapole moment in the external field decays faster than that of the quadrupole moment, because the ratio between the octapole potential and the quadrupole potential is of the order of the ratio of the stellar radius to the distance between the star and the field point, which becomes the impact parameter in the gravitational lensing. This ratio is so small that one can ignore approximately the octapole and higher moment in the lensing study. Furthermore, radio observations such as the VLBI will be improved. If the accuracy is improved sufficiently, the present result will have a direct relevance with such an observation. Hence, we consider in this paper the truncated multipole expansion up to the quadrupole moment.  This consideration in the gravitational lens is not new but rather classical though it was done in a different context. For instance, the light deflection caused by the quadrupole moment of the Sun was discussed in the context of tests of the general relativity (Epstein and Shapiro 1980). Except for a case of the source on the equatorial plane, we need solve the lens equation numerically. A spherically symmetric galaxy and a quadrupole tidal part were extensively studied by Kovner (1987), in which most models are an extended sphere and the quadrupole part, and another model is a star with quadrupole shear induced by a galaxy where the star is located. The last model seems similar to our system of a deformed star. However, it should be noted that this extra shear does not come from a potential of the star but from that of the galaxy. In other words, the dependence of the shear on the radial coordinate is different from that for the deformed star.  It is believed that gravitational lensing study needs some numerical techniques except for a few special cases such as a point mass lens or an axisymmetric mass distribution along the line of sight, because the lens equation is nonlinearly coupled (for instance, Schneider et al. 1992). In recent, however, some development has been achieved: It has been shown that for two cases of a binary lens (Asada 2002) and an ellipsoidal lens (Asada et al. 2003) the lens equation is reduced to a single polynomial. An advantage of this approach is saving time and computer resources, and another one is to enable us to study analytically the lensing, for instance to find out an analytic expression for caustics (Asada et al. 2002, 2003) which are curves on the source plane, and the number of images changes if a source crosses the caustics. This expression for the caustics is nothing but a discriminant for the polynomial, which determines a change in the number of roots for the equation: As known in algebra, the discriminant vanishes in a case of multiple roots, which corresponds to merging images or giant arcs in the gravitational lens. Classification of the critical curves and caustics is of great importance. For instance, it was done intensively for a binary gravitational lens (e.g., Dominik 1999).  Furthermore, it has been shown also that one can derive a single polynomial equation for any set of the lens equations in polynomials by using the Euclidean algorithm (Asada et al. 2004). In practice, however, in a case of the lens equation which is higher than fifth degree, this algebraic procedure needs a computer with a fast CPU and a huge memory, and furthermore the resultant equation may become too complicated to handle any more in the Cartesian coordinates. In order to avoid this problem, the polar coordinates are adopted in this paper so that a simple result can be obtained.  This paper is organised as follows. First, we show that the lens equation is reduced to a single polynomial equation. Next, the analytic expression for the caustics is computed as the discriminant for the polynomial. Finally, a parametric representation for the critical curves and caustics are obtained to clarify a topological feature of the critical curves. ",
        "conclusions": "We have re-examined the lens equation for a gravitational lens due to a deformed star modeled by using quadrupole moment. First, we reduce the lens equation to a single real tenth-order polynomial. Consequently, an analytic expression for the caustics is given by Eq. ($\\ref{caustics}$) as $K=0$, though it is too lengthy to handle. The critical curves are simply expressed as Eq. ($\\ref{critical}$) or Eq. ($\\ref{critical-r}$). The parametric representation of the caustics is given by Eqs. ($\\ref{caustics-a}$) and ($\\ref{caustics-b}$). Hence, the present result must be helpful for understanding the gravitational lensing due to a deformed star. In particular, it gives a nice approximation to a lensing study based on a realistic numerical simulation of a compact object.  For instance, the image position is unstable for a source close to the caustics, so that careful numerical computations are needed. The present approach will make it easy to study such a case. This is a future subject."
    },
    "0410/astro-ph0410595_arXiv.txt": {
        "abstract": "The Supernova Legacy Survey (SNLS)\\footnote{For a complete author list and candidates see http://cfht.hawaii.edu/SNLS} will discover and obtain $g' r' i' z'$ lightcurves for more than 700 spectroscopically confirmed SNe Ia $(0.3<z<0.9)$ to differentiate between competing models for Dark Energy.  We fit the multicolor lightcurves of the candidates to determine which are likely SNe Ia and send them for follow-up spectroscopy to the Keck, VLT, and Gemini telescopes.  Here we show the results from Gemini, where we send most of our highest redshift $(0.6<z<0.9)$ targets. ",
        "introduction": "As a part of the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS), the SNLS uses a rolling search strategy to discover and monitor SNe in four one degree fields over the course of five years. SNe are observed in $g' r' i' z'$, allowing unprecedented measurement of reddening and SN intrinsic colors.  Observations are made every 2-3 rest frame days during dark and gray time.  A typical run includes 10 observing blocks of 5 visits to each of two deep fields. Because we always observe the same fields, every visit brings new SN discoveries and simultaneously adds points to existing lightcurves.  See the contribution of Sullivan et al. to these proceedings for an overview of the SNLS. ",
        "conclusions": ""
    },
    "0410/astro-ph0410306_arXiv.txt": {
        "abstract": "We have obtained deep, multi-band imaging observations around three of the most distant known quasars at redshifts $z>6$. Standard accretion theory predicts that the supermassive black holes present in these quasars were formed at a very early epoch. If a correlation between black hole mass and dark matter halo mass is present at these early times, then these rare supermassive black holes will be located inside the most massive dark matter halos. These are therefore ideal locations to search for the first clusters of galaxies. We use the Lyman-break technique to identify star-forming galaxies at high redshifts. Our observations show no overdensity of star-forming galaxies in the fields of these quasars.  The lack of (dust-free) luminous starburst companions indicates that the quasars may be the only massive galaxies in their vicinity undergoing a period of intense activity. ",
        "introduction": "The existence of inactive supermassive black holes in the nuclei of massive nearby galaxies (Magorrian et al. 1998) is a sure sign that most galaxies underwent substantial active phases. Furthermore, the correlation between the amount of material accreted (black hole mass) and the galaxy mass indicates important links between galaxy formation and AGN activity.  One implication is that the most massive black holes live in the most massive galaxies which reside in the most massive dark matter halos.  The $z>6$ quasars being discovered in the Sloan Digital Sky Survey (see Fan, this volume) have exceptionally high luminosities ($M_{1450}<-27$). The available evidence suggests that these luminosities are not amplified by gravitational lensing (Fan et al. 2003) or beaming (Pentericci et al. 2002; Willott et al. 2003). Assuming that these quasars are radiating at the Eddington limit gives a lower limit on their black hole masses of several billion solar masses. These are comparable to the black holes inside the largest dominant cluster ellipticals at the present time. Applying local calibrations from this black hole mass to galaxy and halo masses suggests that these quasars reside in halos with mass $\\sim 10^{13} \\Msolarcjw$ (e.g. Fan et al. 2001). LCDM simulations show that structure grows `bottom-up' with the largest halos typically collapsing at the latest epochs. Halos with mass $> 10^{13} \\Msolarcjw$ at $z=6$ are therefore extremely rare (the fraction of mass in such halos is of order $10^{-8}$; Sheth \\& Tormen 1999)  There is evidence that at least some of these quasars are located in massive galaxies. Sub-millimetre photometry of several of them show very high far-infrared luminosities, implying extreme starbursts of $> 1000 \\Msolarcjw\\,{\\rm yr}^{-1}$ (Bertoldi et al. 2003; Priddey et al. 2003). The most distant quasar, SDSS\\,J1148+5251 at $z=6.42$, contains $ \\gtsimeqcjw 10^{10} \\Msolarcjw$ of molecular gas, again indicative of a massive, primeval galaxy (Walter et al. 2003). These distant quasars, which are thought to be residing in massive galaxies, are therefore ideal places to search for the first clusters of galaxies. The most straightforward method for identifying galaxies at such high redshifts is the Lyman-break technique -- a large discontinuity in the spectral slope due to absorption by neutral hydrogen (e.g. Steidel et al. 1996). The increase in the optical depth of hydrogen absorption at such high redshifts (e.g. Songaila \\& Cowie 2002) makes the Lyman-break technique even more effective at $z\\sim 6$ than at lower redshifts. To search for galaxies in these putative mass overdensities, we carried out a search for Lyman-break galaxies in fields of three $z>6.2$ SDSS quasars. ",
        "conclusions": ""
    },
    "0410/astro-ph0410441_arXiv.txt": {
        "abstract": "Constraints on an exact quintessence scalar-field model with an exponential potential are derived from gravitational lens statistics. An exponential potential can account for data from both optical quasar surveys and radio selected sources. Based on the Cosmic Lens All-Sky Survey (CLASS) sample, lensing statistics provides, for the pressureless matter density parameter, an estimate of $\\1 = 0.31_{-0.14}^{+0.12}$. ",
        "introduction": "Recently, astronomers have suggested that the universe is roughly geometrically flat and accelerating its expansion. Usual types of matter, i.e. baryons, photons and neutrinos, fail to close the universe and generate attraction, which leads to a decelerated expansion. To reconcile observations with predictions of general relativity, two new components have been added to the energy budget of the universe: pressureless cold dark matter (CDM) and dark energy, also known as quintessence, with negative pressure.  Whereas the simplest explanation for dark energy is a cosmological constant, one of the more interesting physical realizations of the quintessence is a cosmic scalar field minimally coupled with the usual matter action \\cite{pe+ra88,cal+al98}. Such a field induces the repulsive gravitational force dynamically, explaing the accelerated expansion of our universe. $k$-essence, namely a scalar field with non canonical kinetic terms \\cite{arm+al99}, models based on branes and extra dimensions, such as the Cardassian scenario \\cite{zh+fu04}, and the generalized Chaplygin gas \\cite{kam+al01,zhu04} can also drive an accelerated expansion.  A flat model of universe with a subcritical CDM energy density $\\1$, with two-thirds of the critical density existing in the form of dark energy and with negligible amount in standard matter, matches observations from galaxy clustering \\cite{ba+fa98,car&al98}, large-scale structure \\cite{pe&al01,ver+al01}, cosmic microwave background radiation \\cite{deb&al00,bal+al00,ja&al00,pr+al02}, type Ia supernovae \\cite{ri&al98,pe&al99}, age of the universe and Ly$\\alpha$ forest \\cite{ha&ro01,wan+al00}. So, to explain the bulk of the evidence, we require a universe filled, nearly exclusively, with exotic matter.  Gravitational lensing, based on the quite simple and well understood physics of light deflection in a weak gravitational field, provides useful tools for constraining cosmological parameters. Gravitational lensing statistics \\cite{cha+al02,wa+mi98,co&hu99,wan+al00,zhu00,jai+al01}, effects of large-scale structure growth in weak lensing surveys \\cite{be+be01}, Einstein rings in galaxy-quasar systems \\cite{fu&yo00,ya+fu01}, clusters of galaxies acting as lenses on background high redshift galaxies \\cite{io02cl,se+lo04}, offer very promising opportunities to probe quintessence. Gravitational lens time delay measurements can also produce precise estimates of cosmological parameters \\cite{sch04}, and the effect of quintessence in such observations has been investigated \\cite{gi+am01,le+ib02}. Results from techniques based on gravitational lensing  are complementary to other methods and can provide restrictive limits on the dark energy contribution, sometimes in disagreement with the concordance value \\cite{koc96}.  In this paper, gravitational lens statistics is used to study a class of quintessence models with an exponential potential. Such a potential can drive a scaling solution at late time, i.e. the scalar field mimics the scaling of the dominant energy component. Usually, associated with an exponential potential, a scalar field is considered such that its fractional density, $\\Omega_\\varphi$, is practically constant during part of the matter-dominated era. The usual view is that, $\\Omega_\\varphi$ being small at the beginning of the matter-dominated era due to bounds from nucleosynthesis theory, the exponential potential cannot produce a today accelerated expanding universe \\cite{bea+al01}. This conclusion is mainly based on considering the equation of state of the dark energy as nearly constant. Recently, Rubano \\& Scudellaro \\shortcite{ru+sc02} showed how, without such an a priori assumption, the exponential potential can reveal very interesting features. This new position in favour of the exponential potential relies on a very strong argument, i.e. on general exact solutions of the field equations, allowed by a suitable choice of the exponent of the scalar potential. It has been shown that, since the equation of state is not constant, the tracker condition cannot be treated in the usual way. Despite of its simplicity, this class of potentials is able to fit data from type Ia supernovae \\cite{pav+al02}, galaxy clustering \\cite{ru+se02} and cosmic microwave background \\cite{did+al02}.  The paper is as follows. In Section~\\ref{mode}, the quintessence model is introduced. Section~\\ref{stat} contains the basics of gravitational lensing statistics. Properties of the data sample used in the statistical analysis are illustrated in Sec.~\\ref{data}. Section~\\ref{resu} lists the constraints on the cosmological parameters. Section~\\ref{disc} is devoted to some final considerations. ",
        "conclusions": "\\label{disc}  Exponential potentials in quintessence cosmologies have been recently re-considered and offer interesting perspectives \\cite{ru+sc02}. As an example, an exponential potential can entail acceptable values for the CDM density at the nucleosynthesis epoch.  A suitable choice of the exponent allows, through a transformation of variables suggested by the N\\\"{o}ether symmetry approach \\cite{cap+al96}, an exact integration of the Friedmann equations. This is the main attractive feature of the model we have been considering in this paper.  Scenarios alternative to a cosmological constant have been already tested with gravitational lensing statistics. Flat universes with dark energy with a constant equation of state has been considered \\cite{wa+mi98,cha+al02}. Such an ad hoc modelization of the dark energy helps to discriminate between a cosmological constant and an evolving quintessence but it does not account for a general treatment of the dark energy component. On the other hand, physically motivated scalar fields deserve particular attention. The model described in Section~\\ref{mode} presents, in my opinion, some very interesting properties. In general relativity, it is quite exceptional to deal with exact general solutions. This enables a comprehensive analysis and a full treatment of the model and of its phenomenological properties.  Apart from this general argument in favour of exact solutions in general relativity, the examined exponential potential presents a further appeal from the statistical point of view. Since it has only two free parameters, i.e $\\1$ and $H_0$, following Bayesian arguments, it has the same a priori probability of a flat universe with a cosmological constant and must be preferred to a flat quintessence cosmology with a constant, but yet undetermined, equation of state of the dark energy, $w_X$\\footnote{These models are characterized by three parameters: $H_0$, $\\1$ and $w_X$.}. Despite of its simplicity, the examined exact solution reflects many properties of the observed universe, with the same statistical confidence of a constant $\\Lambda$-term \\cite{did+al02,pav+al02,ru+se02}. It is compatible with the supernovae data, whose analysis gives $\\1 =0.15^{+0.15}_{-0.03}$ at 68\\% confidence level \\cite{pav+al02}. The peculiar velocity field and the perturbation growth of structure can also be accounted for in this scenario \\cite{ru+se02}; data from galaxy clustering suggests a value of $\\1 = 0.18 \\pm 0.05$ at 68\\% confidence level. The considered model, not including radiation, is realistic only in the contemporary or, at most, the recent past regimes of the life of the universe. However, a preliminary analysis of the anisotropies of the cosmic microwave background also gives concordant results \\cite{did+al02}.  In this paper, I have considered complementary constraints from gravitational lensing statistics. This analysis is based on different physics phenomena and on independent observations and can strengthen conclusions on a cosmological model. An analysis of the CLASS data sample provides a best-fit estimate, at the $68.3\\%$ confidence level, of $\\1 = 0.31_{-0.14}^{+0.12}$ in the case of 11 lenses. When 10 lenses are considered, the value of the best estimate increases of $\\Delta \\1 \\simeq 0.04$. In addition to radio-selected sources, optical surveys of quasars have been employed to constrain $\\1$. This analysis slightly prefers higher values of the pressureless matter density, $\\1=0.6 {\\pm}0.2$. The simple fact that the results from the two independent samples are compatible at the 1-$\\sigma$ level is encouraging and we can consider constraints on cosmological parameters from gravitational lens statistics as quite reliable. Furthermore, estimates from lensing statistics are also compatible with results from both supernovae and galaxy clustering data.  A number of systematic effects can plague lens statistics analyses \\cite{cha03}. Whereas most of the possible sources of errors are well controlled, the main uncertainty is connected with the adoption of the galaxy LF. As seen in Sec.~\\ref{resu}, the main contribution to the error in the estimate of $\\1$ derives from the uncertainties in the early-type galaxy LF. Furthermore, there is no consensus among observationally derived results for the early-type LF \\cite{lov+al92,cha03}. Actually, a reliable determination of the LF is needed to firmly draw conclusions on cosmological parameters. Nevertheless, as noted in Chae~\\shortcite{cha03}, whenever the total galaxy LF is reliably determined, conclusions on cosmological parameters are accurate if the partition of the total LF into the specific LF is careful."
    },
    "0410/astro-ph0410488_arXiv.txt": {
        "abstract": "We report on 12 years of observations of PSR~J1713+0747, a pulsar in a 68-day orbit with a white dwarf.  Pulse times of arrival were measured with uncertainties as small as 200\\,ns. The timing data yielded measurements of the relativistic Shapiro delay, perturbations of pulsar orbital elements due to secular and annual motion of the Earth, and the pulsar's parallax, as well as pulse spin-down, astrometric, and Keplerian measurements. The observations constrain the masses of the pulsar and secondary star to be $m_1=1.3\\pm 0.2\\,\\msun$ and $m_2=0.28\\pm 0.03\\,\\msun$, respectively (68\\% confidence). Combining the theoretical orbital period-core mass relation with the observational constraints yields a somewhat higher pulsar mass, $m_1=1.53_{-0.06}^{+0.08}\\,\\msun$.  The parallax is $\\pi=0.89\\pm 0.08$\\,mas, corresponding to a distance of $1.1\\pm 0.1$\\,kpc;  the precision of the parallax measurement is limited by uncertainties in the electron content of the solar wind.  The transverse velocity is unusually small, $33\\pm3$\\,km\\,s$^{-1}$.  We find significant timing noise on time scales of several years, but no more than expected by extrapolating timing noise statistics from the slow pulsar population.  With the orientation of the binary orbit fully measured, we are able to improve on previous tests of equivalence principle violations. ",
        "introduction": "\\label{sec:intro}  High precision timing of a radio pulsar binary system can reveal a wealth of information about the dynamics of the binary, the astrometry of the system, and the natures of the pulsar and secondary stars. PSR~J1713+0747, a 4.6\\,ms pulsar in a 67.8~day orbit with a white dwarf, is among the very best pulsars for high precision timing:  it has a high flux density, shallow spectrum, and sharp pulse peak \\citep{fwc93}.  We observed PSR~J1713+0747 over six years as part of a systematic high precision pulsar timing study at the Arecibo Observatory.  We used data acquisition systems employing coherent dedispersion, allowing substantial improvements in timing precision compared to previous work.  Early observations of this source were reported by \\cite{cfw94} (herein CFW), who analyzed 1.5 years of pulsar timing data.  Previous observations have also been reported by \\cite{vb03}. We combined the data of CFW with our newer observations to produce a single data set spanning twelve years.  The superior timing precision of the new data and the longer total time span of observations yield substantial refinements of all measurements reported by CFW, as well as a number of additional new measurements.  High precision detection of the Shapiro delay allows the pulsar and secondary star masses to be separately measured. The Shapiro delay measurement reported in CFW did not have sufficient precision to independently determine the pulsar and secondary star masses. These quantities are of interest because, while measured double neutron star masses fall into a very narrow range, 1.25 to 1.44\\,\\msun\\ \\citep{lp04}, extended accretion during the formation of pulsar--white dwarf binaries may lead to higher pulsar masses in these systems.  High precision timing of pulsars can be used to place limits on the gravitational wave background at frequencies of $10^{-6}$ to $10^{-5}$ Hz by searching for arrival time variations on time scales of years \\citep{lbsn03,ktr94}, but such measurements require that the timing signal not be contaminated by ``timing noise,'' random variations of the pulsar rotation.  Timing noise is inversely correlated with rotation period derivative \\citep{antt94}, so millisecond pulsars, which have very small period derivatives, are prime candidates for gravitational wave studies.  The 12-year time span of data on PSR~J1713+0747 has revealed significant pulse arrival variations beyond those expected from simple magnetic dipole spin-down.  We quantify this apparent timing noise, but we leave the analysis of the PSR~J1713+0747 signal in the context of the gravitational wave background to another work.  We detected both annual and secular perturbations of orbital elements due to the changing Earth--binary line-of-sight. This allowed us to uniquely determine the inclination and position angle of the orbit.  This is only the second pulsar binary system for which timing observations have fully specified both angles of the orbital orientation \\citep{vbb+01,vb03}.  Preliminary results from this project, using data collected through mid-2002, were published elsewhere \\citep[e.g.,][]{nss03,nss04}. The present paper incorporates two further years of high quality data, resulting in some changes to the best-fit parameters, but within the expected uncertainties.  The plan of this paper is as follows.  In \\S 2 we summarize the observations. In \\S 3 we describe the timing model.  In \\S 4 we present the parameters of the pulsar and binary system derived from fitting the timing model to the observational data.  In \\S 5 we analyze the timing model parameters in the contexts of timing noise, stellar masses, pulsar distance and velocity, and theories of strong field gravitation.  In \\S 6 we summarize the key results. ",
        "conclusions": "PSR~J1713+0747 has lived up to its promise as one of the best pulsars for high precision timing, with 200\\,ns residual pulse arrival times attained on time scales of years, and residuals well under 2$\\mu$s over the full 12-yr data set. Major findings include:  1.  The pulsar mass is constrained to $m_1=1.3 \\pm0.2\\,\\msun$ by the measured Shapiro delay.  If the secondary mass is restricted to values predicted by the theoretical orbital period-core mass relation, the pulsar mass is somewhat higher, $m_1=1.53_{-0.06}^{+0.08}$.  2.  The parallax is $\\pi=0.89\\pm0.08$\\,mas, corresponding to a distance of $1.1\\pm0.1$\\,kpc.  This is consistent with predictions based on the pulsar's dispersion measure.  3.  The orientation of the binary has been fully determined by the combined measurements of Shapiro delay and of perturbations of orbital elements due to relative Earth-pulsar motion. The orientation and very low ellipticity of the orbit lead to an improved constraint on deviations from general relativity.  The timing precision attainable on short time scales appears to be limited by measurement precision.  There is every reason to expect timing precision to improve in the coming years as a new generation of wide bandwidth coherent dedispersion systems are employed at radio telescopes, directly resulting in more precise measurement of the Shapiro delay and hence of the pulsar and white dwarf masses."
    },
    "0410/astro-ph0410349.txt": {
        "abstract": "Current understanding of core collapse and thermonuclear supernovae is reviewed. Recent progress in unveiling the nature of cosmic gamma-ray bursts (GRB) is discussed, with the focus on the apparent link of several GRBs with an energetic subclass of stellar explosions, type Ib/c core-collapse supernovae. This relation provides the strong case that the GRB phenomenon is connected with the final stages of massive star evolution and possibly with the formation of neutron stars and black holes. ",
        "introduction": "An extensive discussion of basic physics of core collapse supernovae can be found in \\cite{Bethe1990}; the evolution of massive stars, core collapse, formation of stellar remnants and supernova nucleosynthesis are reviewed in \\cite{Woosley&al2002}; a recent concise discussion of problems and prospects for core collapse supernovae can be found in \\cite{Mezzacappa2004}.  In the end of thermonuclear evolution, the core of a massive star can lose mechanical stability for various reasons. In the stellar mass range $8M_\\odot<M<20 M_\\odot$ a partially degenerate core with mass close to the Chandrasekhar limit $M_{core}\\sim M_{Ch}$ and high density ($\\rho\\sim 10^9-10^{10}$ g/cm$^3$) appears. Under these physical conditions, the chemical potential of degenerate electrons becomes so high that neutronisation reactions $ e^-+(A,Z)\\to (A,Z-1) +\\nu_e $ become effective even at zero temperature. At  densities $>10^6$ g/cm$^3$ degenerate electrons becomes relativistic so the adiabatic index of matter $\\gamma=d\\log P /d\\log \\rho \\to 4/3$, the critical value for loss of mechanical stability. Neutronisation of matter means its deleptonisation (decrease in the lepton number $Y_e=N_e/N_b$), so the pressure at some moment increases slower than $\\rho^{4/3}$ though $\\gamma$ is formally above $4/3$ calculated at constant $Y_e$, and a catastrophic collapse begins. Temperatures are higher at larger stellar masses  $M>20 M_\\odot$ so the collapse is initiated by  photodissociation of nuclei (here $\\gamma$ really becomes $<4/3$). For most massive stars with $M>60 M_\\odot$ pair creation makes $\\gamma<4/3$.  The collapse occurs on the dynamical (free-fall) time scale $t_{ff}\\sim 1/\\sqrt{G\\rho}\\sim $ a fraction of second. Adiabaticity holds so the entropy per baryon $s=S/k_B\\approx const\\sim 1$ and even can increase due to non-equilibrium beta-processes. Low specific entropy (compared to H-burning phase at the main sequence, $s\\approx 10-15$) prevents dissociation of nuclei until they \"touch\" each other at densities of the order of the nuclear density, $\\rho_n\\sim 2\\times 10^{14}$ g/cm$^3$. The collapse stops (if the core mass is below some $M_{max}$), and bounce of the shock occurs at $\\sim 50$ km from the center.  The bounce shock heats up deleptonized matter and rapidly spends most of its kinetic energy to destroy nuclei and produce plenty of free nucleons (n, p). Modified URCA-processes \\cite{Dicus1972} becomes important: $e^-+p \\to n+\\nu_e$, $e^++n \\to p+\\bar \\nu_e$ and pair-neutrino annihilation takes place: $e^++e^-\\to \\nu_e+\\bar \\nu_e$. At the typical collapse temperatures $T\\sim 10$ MeV a lot of $\\nu$'s is produced \\cite{Zeldovich&Gusseinov1965} However, at densities $\\rho\\sim 10^{12}$ g/cm$^3$ the mean free path of 10 MeV neutrinos is by 5-6 orders of magnitude smaller than the size of the proto neutron star ($R\\sim 50$ km) so the opaque \"neutrinosphere\" forms. Most of the core collapse neutrinos diffuse out of the neutrinosphere on a time scale $\\sim 10$ seconds. First calculations of $\\nu$ spectra in core collapse SN were performed by D.K. Nadyozhin \\cite{Nadyozhin1978, Nadyozhin&Otroshenko1980} We should note that subsequent detailed calculations (e.g. \\cite{Messer&al1998} and references therein) did not change much these spectra. Thus, the modest $\\sim 10\\%$ fraction of the total neutrino energy released in the core collapse ($\\sim 10^{53}$ ergs) would be sufficient to unbind the overlying stellar envelope and produce the phenomenon of type II supernova explosion.  \\subsection{Neutrino-driven explosions}  Thermal SN explosion mechanism was proposed by Colgate and White \\cite{Colgate&White1966}. In this picture, part of the neutrino flux liberated in the core collapse is deposited to the stellar mantle to make it unbound ($\\sim 10^{51}$ ergs is needed). Specific mechanisms include neutrino-driven fluid instabilities, for example convection both above neutrinosphere and inside the proto-NS. Neutrino-driven convection, however, may not be as important as thought before, as follows from recent detailed 2D studies of convection \\cite{Buras&al2003}. Instead, other fluid instabilities such as newly found double-diffusive instability (the so-called \"lepto-entropy fingers\") \\cite{Bruenn&al2004}, may effectively increase neutrino luminosity to help successful explosion. Note here that process $\\nu_e+\\bar \\nu_e\\to e^++e^-\\to \\gamma+\\gamma$ may also be important above $\\nu$-sphere \\cite{Berezinsky&Prilutsky1987}. This process was proposed as the energy source for GRB fireballs.  It is now realized that even detailed 1D-calculations of the core collapse supernovae fail to produce explosion (see e.g. reviews by \\cite{Libendoerfer2004, Mezzacappa2004} and references therein). The main reasons are that the bounce shock rapidly stalls (turns into a pure accretion front) over neutrinosphere at $R_s\\sim 200$ km because of nuclear dissociations, and net neutrino heating of the freshly accreted material immediately after the front is insufficient. Burrows \\cite{Burrows2004} notes that 1D-models are about to produces a successful explosion, only 25\\% to 50\\% increase in the energy deposition rate inside the \"gain region\" is required. The stalled shock is not revived because the fresh fall-back matter is advected more rapidly than heated up by neutrinos. An additional heating by the viscous dissipation in differentially rotating models is considered in  \\cite{Thompson&al2004}. Differential rotation of collapsing core seems quite plausible (see evolutionary calculations of rotating stars \\cite{Heger&al2000, Heger&al2004}). The energy stored in shear can amount to $\\sim 10^{52}$ ergs for reasonably short 6 ms spin periods of proto neutron stars. The magneto-rotational instability (MRI) which arises in differentially rotating magnetized fluids with negative angular velocity gradients $d\\Omega/dr<0$ \\cite{Velikhov1959, Chandrasekhar1960}. MRI is thought to be responsible for turbulence in accretion disks and may operate in stars as well. This viscosity was found \\cite{Thompson&al2004} to be comparable with neutrino-driven convective viscosity and much larger than the neutrino viscosity. The account for additional energy dissipation due to the MRI viscosity in differentially rotating stellar cores allowed \\cite{Thompson&al2004} to obtain a successful explosion in their 1D-calculations.  Transition from 1D to multi-dimensional hydrodynamic calculations has enabled a more accurate treatment of fluid instabilities. Presently, it is recognized that even with detailed account of Boltzmann neutrino transport, additional physical properties of the collapsing core, e.g. rotation, need to be coupled with multi-dimensional calculations to produce successful neutrino-driven explosions (see e.g. \\cite{Janka&al2004a, Mezzacappa2004}).  \\subsection{Asymmetric explosions}  There are increasing observational evidence that supernova explosions are generically asymmetric. The indications are as follows.  \\paragraph{Spectroscopic observations}  The asphericity of the explosion is directly inferred from spectropolarimetric observations \\cite{Leonard&Filippenko2004}. The smaller the hydrogen mass of the presupernova envelope, the higher the degree of linear polarization is observed. In the canonical type II-plateau supernovae with heavy hydrogen envelopes (SN1999em, 2003gd etc.) polarization is small (less than 0.5\\%), indicating low degree of asphericity. In contrast, in the stripped-envelope progenitors giving rise to type IIb (SN1993J), type IIn (SN1998S), type Ic (SN2003gf) or peculiar Ic (SN2002ap, maybe hypernova, see below) supernovae, th polarization degree  is large, of order of 1-2\\%. The degree of polarization tends to increase with time (i.e. as more and more deeper layers of the ejecta become transparent). This suggests the relation of asymmetric supernova explosion to pulsar kicks and, in extreme cases, to GRB explosions (see below).  \\paragraph{Pulsar kick velocities}  Based on measurements of proper motions of $\\sim 100$ radio pulsars, it is well established now \\cite{Lyne&Lorimer1994, Arzoumanian&al2002} that young neutron stars demonstrate space velocities up to several 100 km/s which cannot be explained without introducing additional natal kicks.  \\paragraph{Shape of young galactic SNR}  The asphericity  of SNRs 1987a, Cas A, N132D etc. is directly seen in multiwavelength observations \\cite{Park&al2002, Hwang&al2004}. Recent \\textit{Chandra} X-ray observations of galactic SNR W49B \\cite{W49B} suggest this SNR may be the remnant of an asymmetric hypernova explosion in our Galaxy.  \\paragraph{Strong mixing of Ni-56}  Strong mixing of $Ni-56$ in the outer layers of the ejecta is required to explain the observed SN light curves (especially SN1987a, \\cite{Mitchell&al2001}).  \\vskip\\baselineskip Different mechanisms for the SN asymmetry have been proposed.  \\paragraph{Hydrodynamic mechanism}  Large-scale neutrino-driven overturn between the shock and neutrinosphere can cause the SN anisotropy explosion \\cite{Burrows&al1995}. Pulsar kicks up to 500 km/s from neutrino convection were obtained in recent 2D hydrodynamic calculations with Boltzmann neutrino transport \\cite{Buras&al2003}, while in 3D-calculations (but with less rigorous treatment of neutrino transport) \\cite{Fryer2004} high initial velocities of proto-NS were found to be damped by neutrino emission generally to less than 200 km/s. However, the postbounce stellar core flow is found (both in 2D and 3D) to be subjected to generic non-spherical perturbations of the accretion shock leading to the development of the so-called \"stationary accretion shock instability\" \\cite{Blondin&al2003}. This large-scale ($l=1,2$ modes) fluid instability was also found in multi-dimensional simulations in Ref. \\cite{Janka&al2004b}, and may be responsible for observed bimodal pulsar velocity distribution \\cite{Arzoumanian&al2002}. It is important that the developed accretion-shock instability can result in bipolar explosions even in the absence of rotation.  \\paragraph{Magnetorotational mechanism}  Rotation of the pre-supernova core and magnetic field increase during the collapse, and was proposed in the beginning of 1970s \\cite{BK1970} as an alternative mechanism for core-collapse SNe. The magnetorotational mechanism by its nature is asymmetric and can launch antiparallel jets during the explosion. First simulations were made in Ref. \\cite{LeBlanc&Wilson1970, Symbalisty1984}. Results of recent 2D MHD calculations \\cite{Ardeljan&al2004a, Ardeljan&al2004b} are very encouraging: strong differential rotation in the presence of (maybe initially weak) magnetic field was shown to increase the magnetic pressure and form a MHD shock. As a result, rotational energy of neutron star is converted to the energy of the radial expansion of the envelope. The exponential growth of magnetic field in differentially rotating collapsing stellar cores due to MRI \\cite{Akiyama&al2003} appears to help the magnetorotational explosion as well. The magnetorotational explosion was found in \\cite{Ardeljan&al2004b} to be essentially divided into three stages: linear growth of the toroidal component of the magnetic field due to twisting of the magnetic filed lines, the exponential growth of both toroidal and poloidal field components due to development of MHD instabilities, and the formation of a MHD shock leading to the explosion. These 2D MHD simulations \\cite{Ardeljan&al2004b} have shown that the obtained energy of the magnetorotational explosion $\\sim 0.6\\times 10^{51}$ ergs is sufficient to explain type II and type Ib core-collapse supernovae.  \\paragraph{Formation and subsequent collision of a close binary NS+NS system inside the stellar interiors}  V.S. Imshennik \\cite{Imshennik1992} suggested the following scenario: a rapidly rotating core collapse results in the core fusion due to rotational instability (the original idea goes back to von Weizsaecker \\cite{vonWeizsaecker1947}), then the binary NS system coalesces due to gravitational radiation losses. In this scenario, the lighter NS (with larger radius) first fills its Roche lobe and start losing mass as the orbit shrinks due to gravitational radiation. The NS gets unstable when $M<M_{min}\\approx 0.1M_\\odot$ providing an energy release of $\\sim 10^{50}$ erg in neutrinos \\cite{Blinnikov&al1984, Colpi&al1989, Colpi&al1991}. This mechanism has been further elaborated in Ref. \\cite{Imshennik&Popov1994, Aksenov&al1997}. It was suggested for SN1987a explosion \\cite{Imshennik&Ryazhskaya2004} to explain double LSD neutrino signal separated by $\\sim 5$ hours observed from SN1987a. A point of concern here may be a rapidly rotating pre-supernova core required for this mechanism to operate. Recent calculations by the Geneva group of stellar evolution with rotation but without magnetic field \\cite{Hirschi&al2004} indicate that there are enough angular momentum in the pre-supernova core, while calculations with (even approximate) account for magnetic field \\cite{Heger&al2004} show a severe (30-50 times) reducing of the final rotation of the collapsing iron core in massive stars. The core fragmentation during collapse with rotation was also not found in 3D calculations in Ref. \\cite{Fryer&Warren2004}.  \\paragraph{Non-standard neutrino physics}  To explain high pulsar kicks, neutrino asymmetries in high magnetic fields pertinent to young pulsars have been invoked \\cite{Chugaj1984, BK1996, Lai&al2001}. The neutrino asymmetry is a natural consequence of the asymmetry of basic weak interactions in the presence of a strong magnetic field. Neutrino oscillations requiring a sterile neutrino with mass 2-20 keV and a small mixing with active neutrinos was proposed in Ref. \\cite{Kusenko2004}. However, a critical analysis in Ref. \\cite{Janka&al2004b} suggests that neither rapid rotation nor strong magnetic field of the young neutron star can be unequivocally inferred from observations of core-collapse SNR 1987A and Cas A, so the solution to the problem of high pulsar kicks probably should be looked for within the frame of the conventional hydrodynamic SN explosion mechanism (see above). ",
        "conclusions": "Cosmic explosions including various types of supernovae and GRBs are natural consequence of stellar evolution. The big efforts of different groups to theoretically understand the physical mechanism(s) of core collapse SNe appear to be approaching the final phase. The thermal (neutrino-driven) mechanism for core collapse SNe is mostly elaborated at present but still fails to produce a strong explosion. There is understanding why stars do not explode by this mechanism, and the road map is designed how to obtain the explosion. Recent multidimensional simulations with accurate incorporation of Boltzmann neutrino transport indicate that effectively a modest boost in the neutrino luminosity is still required. This can be due to inclusion of new fluid instabilities and more accurate treatment of neutrino processes and microphysics (the neutron star equation of state) into the full 3D hydrodynamic calculations. The rotation and magnetic fields is not yet fully included into calculations, which is a challenging task. First results of 2D MHD calculations of the magnetorotational supernova explosion seem to be encouraging. The most energetic (hypernova) explosions, however, require an additional to neutrino source of energy, and the rotation and magnetic fields can be the principal ingredients.  An impressive progress has been done in multidimensional calculations of thermonuclear explosions of degenerate dwarfs for type Ia supernovae. It is still however unclear whether pure deflagration or delayed detonation is at work in SN Ia. The important problem is to more precisely determine the initial ignition conditions. Detailed radiation hydrodynamic modeling revealed that SN Ia light curves proved to be very sensitive to the explosion models and thus can be used to check the models.  In contrast to SNe, the nature of cosmic GRBs remains unclear. The most important recent progress in understanding GRBs was in establishing the link of at least part of them to unusually energetic SN Ibc explosions (hypernovae). At present several unequivocal GRB-SN associations are known. The two closest GRBs discovered so far (GRB 989425 and GRB 031203) proved to be intrinsically weak compared to the bulk of other GRBs with measured redshifts. They both show a single-peak smooth gamma-ray light curve with no signs of jet-induced breaks in the afterglows. In the third (most strong) case of the GRB-SN association, GRB 030329/SN 2003dh, the GRB light curve is two-peaked, the afterglows show evidence for jet. Modeling of the underlaid hypernovae light curve and spectra revealed the first two cases to require smaller kinetic energies but higher mass of the ejecta and the amount of the synthesized nickel than SN 2003dh. We tentatively propose that the tendency \"weaker, more spherically symmetric GRB - stronger hypernova\" may indicate the formation of a NS in the case of weak GRBs and of a BH in the case of strong variable GRBs as the final outcome of the core collapse. In the NS case the GRB energy comes from the rotational energy of neutron star and is possibly mediated by the electromagnetic field. When BH is formed the GRB energy source is the gravitational energy released during non-stationary accretion onto the black hole or the black hole rotation. It is not still excluded that the GRB phenomenon signals some new physics underlying the formation of compact stars.  Observations of various types of supernovae in other galaxies and especially of a (long awaited!) galactic event by all available means (including neutrino and gravitational wave detectors) should undoubtedly be crucial for further understanding physical mechanisms of cosmic explosions. We are sure that the increasing statistics of GRB/SNe in the nearest future obtained with new GRB-dedicated space missions like SWIFT will tell us much more on the nature of GRBs and their progenitors."
    },
    "0410/astro-ph0410394_arXiv.txt": {
        "abstract": "We develop the pseudo-$C_\\ell$ method for reconstructing the Cosmic Microwave Background (CMB) temperature and polarization auto- and cross-power spectra, and estimate the pseudo-$C_\\ell$ covariance matrix for a realistic experiment on the cut sky. We calculate the full coupling equations for all six possible CMB power spectra, relating the observed pseudo-$C_\\ell$'s to the underlying all-sky $C_\\ell$'s, and test the reconstruction on both full-sky and cut-sky simulated CMB data sets. In particular we consider the reconstruction of the $C_\\ell$ from upcoming ground-based polarization experiments covering areas of a few hundred deg$^2$ and find that the method is fast, unbiased and performs well over a wide range of multipoles from  $\\ell=2$ to $\\ell=2500$. We then calculate the full covariance matrix between the modes of the pseudo-temperature and polarization power spectra, assuming that the underlying CMB fields are Gaussian randomly distributed. The complexity of the covariance matrix prohibits its rapid calculation, required for parameter estimation. Hence we present an approximation for the covariance matrix in terms of convolutions of the underlying power spectra. The coupling matrices in these expressions can be estimated by fitting to numerical simulations, circumventing direct and slow calculation, and further, inaccurate analytic approximations. We show that these coupling matrices are mostly independent of cosmology, and that the full covariance matrix for all six pseudo-$C_\\ell$ power spectra can be quickly and accurately calculated for any given cosmological model using this method. We compare these semi-analytic covariance matrices against simulations and find good agreement, the accuracy of which depends mainly on survey area and the range of cosmological parameters considered. ",
        "introduction": "The study of the CMB has undergone a revolution in the past few years, culminating in the first-year Wilkinson Microwave Anisotropy Probe (WMAP) observations (Bennet et al. 2003). The WMAP results along with those from large scale galaxy surveys such as the 2-degree Field Galaxy Redshift Survey (2dFGRS; Percival et al. 2001) and the Sloan Digital Sky Survey (SDSS; Tegmark et al. 2004) have revolutionised our understanding of the Universe, yielding measurements of cosmological parameters to unprecedented accuracy and providing robust tests of the standard cosmological model. Furthermore, the pace of progress in observations of the CMB continues unabated. Following on from the exquisite WMAP observations of the CMB temperature field, there is a number of ground-based and balloon-borne experiments designed to measure the polarization of the CMB radiation. Indeed the even-parity E-mode polarization signal has already been detected, initially by the Degree Angular Scale Interferometer (DASI) experiment (Kovac et al. 2002). The first-year WMAP observations then yielded a detection of the cross correlation between the temperature and the E-mode polarization (Kogut et al. 2003) and very recently, the Cosmic Background Imager (CBI) has detected both the E-mode signal and its cross correlation with the temperature field (Readhead et al. 2004).  Over the next few years, one of the central goals of CMB research will be to tie down the polarization power spectra and take advantage of their different cosmological dependencies, particularly their use as a direct probe of inflation (e.g. Dodelson, Kinney \\& Kolb, 1997; Kamionkowski \\& Kosowsky, 1998). With this in mind, a precise measurement of the CMB polarization field would be extremely useful and there are a number of upcoming experiments designed to make such measurements of both the E-mode and, perhaps more importantly, the odd-parity B-mode polarization signal. In addition to planned satellite missions with the capability of large-scale, high-precision polarization measurements (e.g. the Planck, Tauber 2004, and proposed Inflation Probe\\footnote{http://universe.gsfc.nasa.gov/program/inflation.html} experiments), there are a number of funded ground-based experiments designed to measure the polarization field with very high signal-to-noise over a limited region of sky. In particular, the QUaD\\footnote{http://www.astro.cf.ac.uk/groups/instrumentation/projects/} (QUEST\\footnote{$Q$ and $U$ Extragalactic Submillimetre Telescope} and DASI), and BICEP\\footnote{http://bicep.caltech.edu} (Background Imaging of Cosmic Extragalactic Polarization) experiments should begin observations in 2005, while the CLOVER (CMB Polarization Observer; Taylor et al. 2004) experiment is planned for 2008. A detection and measurement of B-mode polarization in the CMB would provide a major step forward in the field and, if a B-mode signal were found to be present on large enough angular scales, this would constitute a direct detection of primordial gravitational waves (Seljak \\& Zaldarriaga 1997; Kamionkowski, Kosowsky \\& Stebbins 1997). Such a measurement would have immense consequences for theories of the early Universe, possibly providing a handle on the energy scale of inflation (e.g. Kinney 1998).  These ground-based experiments will measure the CMB signal over finite areas of sky, typically a few hundred square degrees. There already exist well-established methods for recovering the true underlying CMB temperature power spectrum from observations of such finite areas of sky and these are well tested (G\\'orski 1994; Tegmark 1997; Bond, Jaffe \\& Knox 1998; Dor\\'e, Knox \\& Peel 2001; Szapudi, Prunet \\& Colombi 2001a; Szapudi et al. 2001b; Wandelt, Hivon \\& G\\'orksi 2001; Hansen, G\\'orski \\& Hivon 2002; Hivon et al. 2002). Although these methods, in principle, should be equally applicable to the polarization power spectra, many have not been tested with detailed simulations. Hansen \\& G\\'orski (2003) investigate temperature and polarization power spectra estimation using Gabor transforms while Chon et al. (2004) have extended the analysis of Szapudi et al. (2001a) to construct a polarization power spectrum estimator based on correlation functions.  In this paper, we extend, and test with simulations of the CMB temperature and polarization fields, the pseudo-$C_\\ell$ method for fast estimation of CMB power spectra. Although not optimal, the pseudo-$C_\\ell$ method has the advantage that it is fast and completely applicable to mega-pixel data sets such as those expected from upcoming satellite and ground-based experiments. We have generalised the method first introduced for the temperature field, $T$, by Hivon et al. (2002), and extended to the $E$ and $B$ modes of polarization by Kogut et al. (2003), to the full set of six CMB power spectra: $C_{\\ell}^{TT}$, $C_{\\ell}^{EE}$, $C_{\\ell}^{BB}$, $C_{\\ell}^{TE}$, $C_{\\ell}^{TB}$ and $C_{\\ell}^{EB}$. Note that this method for recovering power spectra is completely applicable to other cosmological data sets such as large-scale weak lensing surveys, via the transformation, $\\{ T, Q, U, E, B\\} \\rightarrow \\{ \\mu, \\gamma_1, \\gamma_2, \\kappa, \\beta\\}$, where $\\mu$ is the magnification, $\\gamma_1$ and $\\gamma_2$ are the two components of the observed shear, $\\kappa$ is the even-parity lensing convergence and $\\beta$ is a non-gravitational odd-parity distortion.  We also present, and test with simulations, a new method for fast and accurate calculation, for any given cosmological model, of the full covariance matrix of all six pseudo-$C_\\ell$ power spectra. This method will be extremely useful in estimating cosmological parameters directly from the observed pseudo-$C_\\ell$ power spectra. We also present the analytic expressions for the covariances (see also the calculations of Hansen \\& G\\'orski 2003 and Challinor \\& Chon 2004). These expressions are not easily simplified and evaluated so we propose fitting the mixing matrices which appear in these expressions from Monte-Carlo (MC) simulations. Provided that these mixing matrices are mostly independent of cosmology, this method allows one to calculate the covariance matrix with relative ease for any given cosmological model.  The paper is organised as follows. In Section 2, we present the calculation of the pseudo-$C_\\ell$ method relating the observed to full-sky power and cross-power spectra for the temperature ($T$) and polarization ($E,B$) fields of the CMB. In Section 3, we calculate the full covariance matrix of the pseudo-$C_\\ell$ estimator and we outline our method for calculating it for a given cosmological model.  In Section 4, we test the pseudo-$C_\\ell$ reconstructions and the covariance matrix fitting method on realistic simulations of CMB temperature and polarization experiments for both full-sky and finite-area data sets, and we also discuss the results and implications of these simulation tests. Our conclusions are presented in Section 5.  \\vspace{-5mm} ",
        "conclusions": "Pseudo-$C_\\ell$ estimators and their use in reconstructing full-sky power spectra have become popular tools in recent years for measuring the temperature and polarization power spectra from mega-pixel observations of the CMB (Netterfield et al. 2002; Ruhl et al. 2002; Beno\\^it et al. 2003; Hinshaw et al. 2003). In this paper, we have investigated in detail the reconstruction of the full set of CMB temperature and polarization full-sky power spectra from the pseudo-$C_\\ell$ estimators. We have presented the calculation of the pseudo-$C_\\ell$ method relating the observed to full-sky power and cross-power spectra for the complete set of CMB temperature and polarization spectra: $C^{TT}_\\ell$, $C^{EE}_\\ell$, $C^{BB}_\\ell$, $C^{TE}_\\ell$, $C^{TB}_\\ell$ and $C^{EB}_\\ell$. These calculations provide one with a mechanism to reconstruct all six CMB power spectra from observations of the CMB over a finite region of sky.  Having calculated the relevant coupling matrices required for the power spectra reconstructions, we have tested these reconstructions, which we have formulated in terms of bandpowers, on simulated CMB temperature and polarization data sets. We have conducted these tests on realistic simulations of two upcoming CMB experiments -- (1) the Planck satellite where we have considered the reconstruction of the power spectra from two sky surveys over 14 months and (2) a combination of the ground-based QUaD polarization experiment with the temperature data from 4-year WMAP observations. In both cases, we have found the estimators to be unbiased over the entire range of multipoles investigated between $\\ell=0$ and $\\ell=2500$. For these tests, we have included primordial gravitational waves and the effects of gravitational lensing in our input cosmological model in order to test the reconstruction of the B-mode polarization power spectrum. Using an input tensor-to-scalar ratio of $T/S=0.05$, we recover a significant detection of the lensing-induced B-mode power from both the Planck and QUaD simulations for multipoles, $\\ell \\le 1000$. The reconstructions of the non-cosmological cross-spectra, $C^{TB}_\\ell$ and $C^{EB}_\\ell$ are consistent with zero on all scales for both experimental scenarios. The reconstruction of these two spectra can be useful as a test for systematic effects in either the data itself or in the data analysis pipeline of CMB experiments.  We have proposed a new method for fast and accurate calculation of the covariance matrix of the pseudo-$C_\\ell$ estimators for any given cosmological model. Having presented the expressions for the full covariance of the pseudo-$C_\\ell$ estimators, we have made a symmetrized approximation, similar to that used by Efstathiou (2004) for the temperature-only case. Our proposed method is based on fitting the mixing matrices of the approximated expressions from the covariances measured from simulated data sets. We have tested this semi-analytic approximation by comparing it with the covariance measured from simulations for a range of cosmological models drawn from within the WMAP 3$\\sigma$ parameter space region. These tests have been carried out for both a Planck-like experiment, with a galactic cut applied and for a QUaD-like experiment covering $\\sim 300$ deg$^2$. We find that the semi-analytic approximation reproduces all the major features of the simulated covariances with varying degrees of accuracy which depend mainly on survey area and the range of cosmological parameters investigated. The effect of our approximation on the estimation of cosmological parameters will be investigated in a forthcoming paper. Comparing the errors measured from the simulations with those predicted by the ``Knox'' formulae, we find that the latter under-estimate the true errors by $\\sim 10$ per cent for a near full-sky experiment and by $\\sim 25$ per cent for a small-scale experiment like QUaD which spans $\\sim 0.78$ per cent of the sky.  In the not-too-distant future, there is the prospect of high-quality full-sky mega-pixel maps of not only the CMB temperature field, but also of the polarization field from satellite missions such as WMAP, Planck and later perhaps, NASA's proposed Inflation Probe experiment. More immediate prospects come from a number of novel ground-based polarization experiments such as QUaD and BICEP, both of which are due to begin observations in 2005. Beyond this, the CLOVER experiment in planned for 2008. These present the astronomical community with the opportunity of continuing the rapid progress in cosmology already achieved with the 2dFGRS, WMAP and SDSS programs. In addition to improving parameter constraints by up to a factor of 6 (e.g. Bowden et al. 2004), such polarization information may also provide an opportunity to detect and perhaps measure the amount of primordial gravitational waves in the Universe through measurements of the B-mode polarization power spectrum. Such a measurement would constitute an important addition to our knowledge of the physics of early Universe, possibly providing a handle on the energy scale, and other parameters, of inflation (e.g. Kinney 1998). If the rapid progress in CMB observations continues as envisaged, then the methods and results presented here can be used to quickly and accurately extract power spectra and cosmological parameters from CMB temperature and polarization data.   \\vspace{-5mm}"
    },
    "0410/astro-ph0410507_arXiv.txt": {
        "abstract": "It is generally assumed that metals detected in the spectra of a few cool white dwarfs cannot be of primordial origin and must be accreted from the interstellar medium. However, the observed abundances of hydrogen, which should also be accreted from the interstellar medium, are lower than expected from metal accretion. Magnetic fields are thought to be the reason for this discrepancy. We have therefore obtained circular polarization spectra of the helium-rich white dwarfs GD\\,40 and L745-46A, which both show strong metal lines as well as hydrogen. Whereas L745-46A might have a magnetic field of about $-$6900~G, which is about two times the field strength of 3000~G necessary to repell hydrogen at the Alf\\'en radius, only an upper limit for the field strength of GD\\,40 of 4000~G (with 99\\% confidence) can be set which is far off the minimum field strength of 144000~G to repell hydrogen. ",
        "introduction": "A few helium-rich white dwarfs at the cool end of the white dwarf sequence show evidence for metal lines in their spectra. Since radiation forces are not strong enough to compete with gravity for temperatures below 40000~K it is expected that metals sink down on time scales which are short compared to the cooling age. If nevertheless metals are observed in the atmospheres, they must have come from outside the star. The most popular mechanism for this is accretion from interstellar matter. The observed metal abundances in helium-rich white dwarfs are in agreement with accretion in solar element proportions (Dupuis et al. 1992, 1993a, 1993b). However, the observed abundances of hydrogen are much too low to be compatible with the accretion rates inferred from metal accretion. The mechanism most widely discussed as reason for this ``hydrogen screening'' is the propeller mechanism adopted to white dwarfs by Wesemael \\& Truran (1982): Metals are accreted in the form of grains onto a slowly rotating, weakly magnetized (10$^3$-10$^5$~G) white dwarf, whereas ionized hydrogen is repelled at the Alfv\\'en radius. ",
        "conclusions": "For L745-46A the predicted magnetic field strength from the weak field approximation amounts to $-$6900~G with a formal 1$\\sigma$ error of 2100~G and a 99\\% statistical confidence range of $\\pm$5550~G. This magnetic field strength exceeds the minimum field strength of 3000~G necessary for the propeller mechanism to work. However, one should keep in mind, that this result is based on the H$\\alpha$ line only. If the detection is confirmed, this means for the first time an indication that magnetic fields may play a role in the accretion of hydrogen and metals onto cool helium-rich white dwarfs.  However, we could not detect signatures of a magnetic field in the circular polarization or flux spectrum of GD\\,40. This does not necessarily mean, that no magnetic field is present, because if we are looking on the magnetic equator of a magnetic dipole ($i$=90$^\\circ$), the components of the magnetic field along the line of sight completely cancel and no circular polarization can be detected. According to Wesemael \\& Truran (1982) a field strength of about 144000~G is needed for GD\\,40 to let the propeller mechanism work. If we now assume that we could detect a field strength of 4000~G (derived upper limit field strength from the fit to all three lines with a confidence of 99\\%) we could estimate an inclination angle of greater than 84$^\\circ$ for which the longitudinal field strength becomes to low to be detected (Brown et al. 1977). Therefore the chance to miss a magnetic field of 144000~G is only about 6\\%. From our observations we must therefore conclude that in GD\\,40 probably other mechanisms than magnetic fields prevent hydrogen from accreting. More details can be found in Friedrich et al. (2004)."
    },
    "0410/astro-ph0410731_arXiv.txt": {
        "abstract": "Many aspects of the explosion mechanism of Type Ia supernovae (SN Ia) still remain unclear -- causing uncertainties in the derived cosmological parameters. Realistic models of the generation and transport of radiation in the ejecta are required which link theoretical explosion models to observations. We aim to construct theoretical spectra and light curves from consistent radiative transfer models that allow to study the dependence of observable features on the physical parameters of the explosion. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410057_arXiv.txt": {
        "abstract": "We present the results of a set of three-dimensional numerical simulations of magnetic flux emergence from below the photosphere and into the corona.  The corona includes a uniform and horizontal magnetic field as a model for a pre-existing large-scale coronal magnetic system. Cases with different relative orientations of the upcoming and coronal fields are studied. Upon contact, a concentrated current sheet with the shape of an arch is formed at the interface that marks the positions of maximum jump in the field vector between the two systems. Relative angles above $90^{^\\circ}$ yield abundant magnetic reconnection and plasma heating. The reconnection is seen to be intrinsically three-dimensional in nature and to be accompanied by marked local heating. It generates collimated high-speed outflows only a short distance from the reconnection site and these propagate along the ambient magnetic field lines as jets.  As a result of the reconnection, magnetic field lines from the magnetized plasma below the surface end up connecting to coronal field lines, thus causing a profound change in the connectivity of the magnetic regions in the corona. The experiments presented here yield a number of features repeatedly observed with the {\\sl TRACE} and {\\sl YOHKOH} satellites, such as the establishment of connectivity between emergent and pre-existing active regions, local heating and high-velocity outflows. ",
        "introduction": "The {\\sl TRACE} satellite has provided impressive movies showing the dramatic phenomena that accompany the emergence of magnetized plasma in the solar atmosphere \\citep{Schrijver99}.  Initially the emerging flux adopts the shape of a small system of loops in the corona, with their ends anchored in two localized flux patches in the photosphere.  As the loop system grows in size it starts feeling the presence of the ambient coronal magnetic field. This eventually leads to profound changes in the geometry and topology of the magnetic flux systems, with the emerging field quickly establishing links to the coronal field.  This process can take place only through magnetic reconnection that, at the same time, is responsible for significant localized energy release. This provides the basic mechanism to create the several million degree plasma structures observed by the {\\sl Yohkoh} soft X-ray telesope above emerging flux regions and the cooler loop structures observed by {\\sl TRACE} to connect emerging regions to the ambient magnetic field.   The first numerical MHD studies of the emergence scenario were two-dimensional models \\citep{Yokoyama+Shibata95,Yokoyama+Shibata96}. These experiments showed the formation of high-temperature jets, which were tentatively identified with X-ray jets.  However, the geometry of the complex structures \\citep{Klimchuk97} resulting from flux emergence is intrinsically three-dimensional.  A number of recent three-dimensional numerical experiments (\\citealt{Fan01,Magara+Longcope01}; \\citeyear{Magara+Longcope03}; \\citealt{Archontis+ea04, Manchester+ea04}) have studied the rise of buoyant magnetic flux tubes initially below (but close to) the photosphere into a non-magnetized corona.  However, the solar corona has a space-filling and dynamically important magnetic field. Hence, the rise and expansion of newly emerging magnetic flux cannot proceed unimpeded \\citep{Fan+Gibson02}.   Here we present first results of three-dimensional MHD emergence experiments that include a large-scale coronal magnetic field modeled, for simplicity, as a uniform horizontal field (see \\citealt{Archontis+ea04} for details regarding the MHD equations and the numerical approach). The experiments reveal a number of strikingly simple features such as the appearance of an arch-shaped high current density sheet containing the reconnection sites; change of connectivity of the field lines that approach the current sheet following regular patterns; formation of high-velocity plasma outflows emerging from the reconnection region; and plasma heating associated with the reconnection site. They are discussed in the following sections. ",
        "conclusions": "The present experiments show that interaction between the emerging flux and the existing coronal magnetic field is unavoidable and that it can cause a highly dynamical evolution. They also show that the efficiency of the interaction between the two flux systems critically depends on the relative orientation between their magnetic field components. Fast magnetic reconnection can be delayed when their orientations are nearly parallel, while it is easily initiated when they are nearly antiparallel. An arch-shaped current sheet is formed; all along its rim the reconnection generates high-velocity outflows, which are seen propagating horizontally away from the emergence site in the ambient coronal field. The experiments also show that nonsymmetric features arise when the field vector of the two flux systems is not contained in the same vertical plane. Furthermore it is found that simple extensions of the magnetic reconnection process in traditional two-dimensional models cannot accomodate the fully three-dimensional features occurring in the reconnection in these experiments.  Scaling the model to solar conditions, the experiment covers a horizontal domain of $20$ x $23$ Mm$^{2}$, corresponding to the size of a small active region. The emergence phase into the corona covers only about 30 minutes. Nonetheless, the reconnection jets obtained in the most favorable cases have high velocities, with peak values of the order of $400$ Km s$^{-1}$. At the same time, the magnetic dissipation heats the plasma significantly in the current sheet, with temperatures in the antiparallel case reaching close to $10^7$ K. Higher numerical resolution - lower effective resistivity - will increase the energy dissipation in the current sheet and with that the peak temperature.  This will be countered by including anisotropic heat conduction and optically thin radiation in the energy equation - the task for a future investigation.  The results obtained in our experiments compare favorably with observations, where typical jet velocities \\citep{Shimojo+ea96} are between $10$ and $1000$ km s$^{-1}$, with an average velocity $200$ km s$^{-1}$ and where plasma temperatures of many million degrees are often reported in connection with fast energy release events, such as flares."
    },
    "0410/astro-ph0410261_arXiv.txt": {
        "abstract": "An analysis of the X-ray variability of the low luminosity Seyfert nucleus \\ngc, based on a long \\xmm\\ observation, is presented.  The power spectrum shows a clear break from a flat  spectrum ($\\alpha \\approx 1$) to a steeper spectrum ($\\alpha \\approx 2$) at a frequency $f_{\\rm br} = 0.5 - 3.0 \\times 10^{-3}$~Hz, comparable to the highest characteristic frequency found previously in a Seyfert galaxy.   This extends the measured $M_{\\rm BH}-f_{\\rm br}$  values to lower  $M_{\\rm BH}$ than previous studies of  Seyfert galaxies, and is consistent with  an inverse scaling of variability frequency with black hole mass.  The variations observed are among the most violent seen in an AGN to date, with the fractional  rms amplitude ($F_{\\rm var}$) exceeding $100$ per cent in the softest band.  The amplitude of the variations seems intrinsically higher in \\ngc\\ than most other Seyferts, even after accounting for the differences in characteristic frequencies. The origin of this difference is not clear, but it is unlikely to be a high accretion rate ($L/L_{\\rm Edd} \\ls 20$ per cent for \\ngc).  The variations clearly follow the linear rms-flux relation, further supporting the idea that this is a ubiquitous characteristics of accreting black holes.   The variations are highly coherent between different energy bands with any frequency-dependent time delay limited to $\\ls 1$ per cent. ",
        "introduction": "\\label{sect:intro}  The nucleus of the dwarf galaxy \\ngc\\ (Filippenko \\& Sargent 1989) lies at the edge of the parameters  space formed by Active Galactic Nuclei (AGN). It is among the nearest ($D \\sim 4.0~\\Mpc$; Thim \\et 2004) and lowest luminosity ($L_{\\rm Bol} \\sim 5 \\times 10^{40}~\\erg~\\ps$; Moran \\et 1999; Lira \\et 1999) Seyfert galaxies known, and harbours a central black hole with a  mass probably in the range $M_{\\rm BH} \\sim 10^4 - 10^5 \\Msun$ (Filippenko \\& Ho 2003), which compares to more luminous Seyfert galaxies with black hole masses typically in the range $M_{\\rm BH} \\sim 10^6 - 10^8 \\Msun$ (Wandel, Peterson \\& Malkan 1999). In addition, it shows some of the strongest X-ray variability of any radio-quiet AGN (Iwasawa \\et 2000; Shih, Iwasawa \\& Fabian 2003) with variations of nearly an order of magnitude in a few thousand seconds. Yet in other respects, such as optical and ultraviolet emission (and absorption) line properties, it resembles its more luminous counterparts (Filippenko \\& Sargent 1989; Crenshaw \\et 2004).  The X-ray variations of Seyfert galaxies are characterised by a `red noise' power spectrum (Lawrence \\et 1987; M$^{\\rm c}$Hardy 1989), approximated by a steep power law  at high frequencies [$\\mathcal{P}(f)\\propto f^{-\\alpha}$] with an index $\\alpha \\approx 2$ breaking to a flatter index $\\approx 1$ below some break frequency $f_{\\rm br}$ (e.g. Edelson \\& Nandra 1999; Nandra \\& Papadakis 2001; Uttley, M$^{\\rm c}$Hardy \\& Papadakis 2002; Vaughan, Fabian \\& Nandra 2003; Markowitz \\et 2003; M$^{\\rm c}$Hardy \\et 2004).  The prevailing hypothesis is that $f_{\\rm br}$ scales  inversely with the mass of the black hole (or, equivalently, the break timescale scales with $M_{\\rm BH}$). Indeed, the limited data to date are consistent with a simple scaling relation for all black holes, from Galactic Black Holes (GBHs) in X-ray binaries ($M_{\\rm BH} \\sim 3-20~\\Msun$; M$^{\\rm c}$Clintock \\& Remillard 2004) to supermassive black holes in AGN ($M_{\\rm BH} \\sim 10^6 - 10^8 \\Msun$). If this relation holds then the break timescales in Seyferts may be used to infer the mass of the central black hole by  comparison with the power spectra of well studied GBHs such as Cygnus X-1 (Belloni \\& Hasinger 1990; Nowak \\et 1999) which contains a $M_{\\rm BH} \\approx 10 \\Msun$ black hole (Herrero \\et 1995). Uncovering the dependence of the characteristic variability frequencies with the fundamental source parameters (mass, luminosity, accretion rate, etc.) may help clarify the origin of the variability itself.  However, this simple hypothesis is difficult to test (see discussions in Markowitz \\et 2003; M$^{\\rm c}$Hardy \\et 2004 and Papadakis 2004). GBHs typically possess masses $M_{\\rm BH} \\sim 10 \\Msun$ whereas the majority of the well studied Seyfert galaxies have masses $M_{\\rm BH} \\gs 10^6 \\Msun$. Thus the $10^2 - 10^6 \\Msun$ region of parameter space remains unexplored. \\ngc, with its unusually low mass black hole, falls in this critical region. Previous studies of its X-ray variability (Shih \\et 2003) using \\asca\\ suggested the expected power spectral break was present (at $f_{\\rm br} \\sim 3 \\times 10^{-4}$~Hz), but this was poorly determined due to the interrupted sampling of the data (caused by the periodic Earth occulations and SAA passages of the low Earth orbit). \\xmm, with its $\\sim 2$ day orbit and high throughput, does not suffer from this problem and has been providing the best constraints on the high frequency power spectra  of Seyferts (Vaughan \\et 2003a; Vaughan \\& Fabian 2003; M$^{\\rm c}$Hardy \\et 2004). The present paper describes a time series analysis of the data from a long $\\sim 100$-ks \\xmm\\ observation of \\ngc\\ designed to confirm the break timescale and test the $M_{\\rm BH}-f_{\\rm br}$ scaling hypothesis. A future paper (Iwasawa \\et in prep.) will discuss other aspects of this observation. ",
        "conclusions": "\\label{sect:disco}  \\subsection{Summary of results}  The high frequency variability properties of \\ngc\\ have been examined using a long ($\\sim 100$~ks) \\xmm\\ observation.  The variations clearly followed the linear rms-flux relation previously observed in both GBHs and other Seyfert galaxies,  further supporting the idea that this is a universal property of the X-ray variations from accreting black holes.  The high frequency power spectrum was steep $\\alpha \\approx 2$ above a break frequency $f_{\\rm br} \\approx 2 \\times 10^{-3}$~Hz, below which it broke to a slope $\\alpha \\approx 1$. The break frequency is the second highest recorded to date for an AGN (the record is presently held by Ark 564; Papadakis \\et 2002). The cross-spectrum between different energy bands shows the variations at frequencies $<f_{\\rm br}$ were highly coherent and any frequency dependent time delay is constrained to be shorter than $\\tau \\ls 0.01 f^{-1}$.  In other Seyferts, such as NGC~7469 (Papadakis \\et 2001), \\mcg\\ (Vaughan \\et 2003a)  and NGC~4051 (M$^{\\rm c}$Hardy \\et 2004), the delays are typically $\\sim 1$ per cent. The limits of the time delays for \\ngc\\ therefore did not rule out typical time delay characteristics. Although the variations in these relatively broad-band light curves were highly coherent, there is a strong energy-dependence to the amplitude of the variations. This will be discussed in a forthcoming paper (Iwasawa \\et in prep).   \\subsection{The $M_{\\rm BH} - f_{\\rm br}$ relation}  Fig.~\\ref{fig:m-f} shows the $M_{\\rm BH} - f_{\\rm br}$ relation for ten Seyfert galaxies using measurements from the literature (see also Markowitz \\et 2003; M$^{\\rm c}$Hardy \\et 2004; Papadakis 2004), together with the new break frequency determined for \\ngc.  The power  spectral break frequencies from the other ten Seyferts were collected from the following papers: Ark~564 (Papadakis \\et 2002); NGC~4051 (M$^{\\rm c}$Hardy \\et 2004); NGC~3516, NGC~3783, NGC~4151, NGC~5548, Fairall~9 (Markowitz \\et 2003); Mrk~766 (Vaughan \\& Fabian 2003); \\mcg\\ (Vaughan \\et 2003a) and NGC~5506 (Uttley \\et 2002). Black hole mass estimates were taken from the  reverberation mapping studies where possible: Ark~564 (upper limit from Collier \\et 2002); NGC~4051 (Shemmer \\et 2003); NGC~3516 (Wandel \\et 1999); NGC~3783, NGC~4151, NGC~5548, Fairall~9  (Kaspi \\et 2000).  For NGC~5506 the mass was estimated using the $M_{\\rm BH}-\\sigma_{\\ast}$ relation (section 2.3 of Papadakis 2004).  For Mrk~766 and \\mcg\\ the mass estimates are very uncertain.  Based on its bulge luminosity, Uttley \\et (2002) argued for $M_{\\rm BH} \\sim 10^6$ in \\mcg. Due to the ambiguities in this method the error bars were taken to span an order of magnitude in each direction. For Mrk~766, Wandel (2002) estimated $M_{\\rm BH} \\sim 10^7 ~ \\Msun$ based on the optical luminosity and line widths, while Woo \\& Urry (2002) estimated $\\sim 3.5 \\times 10^6 ~ \\Msun$ using the same reasoning. For the purpose of comparison a value of $6.8 \\times 10^{6} ~ \\Msun$ (the mean of these two estimates) was used, again assuming an order of magnitude uncertainty in either direction.  Also shown are the typical frequencies of the high frequency break in Cyg X-1 in both the low/hard and high/soft states.  For \\ngc\\ the break frequency $f_{\\rm bf} = 1.9_{-1.5}^{+0.8} \\times 10^{-3}$~Hz was used together with a mass estimate of $M_{\\rm BH} = 10^4 - 10^5 ~ \\Msun$. This is the `best guess' of Filippenko \\& Ho (2003) consistent with the available estimates: they measured a firm upper limit on the mass of $\\ls 6.2 \\times 10^{6} ~ \\Msun$  from stellar velocity measurements. Based on  photoionisation modelling of optical emission lines, Kraemer \\et (1999) derived an estimate of $\\sim 1.2 \\times 10^5 ~ \\Msun$, while Filippenko \\& Ho (2003) estimated $\\sim 1.3 \\times 10^4 ~ \\Msun$ based on the relationship between optical luminosity and broad line region size. These are consistent with the mass estimated from the $M_{\\rm BH}-\\sigma_{\\ast}$ relation (Tremaine \\et 2002): $\\sim 6.6 \\times 10^4 ~ \\Msun$. It is clear that, given the significant uncertainties in the mass estimates, there is no glaring inconsistency between the new \\ngc\\ measurement and the hypothesised $M_{\\rm BH} \\propto f_{\\rm br}^{-1}$ relation. However, the details (such as whether the index of the relation is $-1$ and which, if any, state of Cyg X-1 provides the best comparison) remain to be investigated as more Seyfert power spectra are measured.  \\begin{figure} \\centering \\includegraphics[width=6.40 cm, angle=270]{fig7.ps} \\caption{ The $M_{\\rm BH} - f_{\\rm br}$ relation for $11$ Seyfert galaxies including \\ngc. The masses are from reverberation mapping experiments except for the four objects marked using dotted error bars. Also shown are typical break frequencies for Cyg X-1 in both its low/hard and high/soft states. The dotted lines show example $M_{\\rm BH} \\propto 1/f_{\\rm br}$ relations consistent with the Cyg X-1 points. \\label{fig:m-f}} \\end{figure}  Hayashida \\et (1998) estimated the  black hole masses for Seyferts by directly comparing their X-ray power spectra with that of Cyg X-1. They measured the frequency $f_c$ at which the power spectrum reaches the level $f_c \\times P(f_c) = 10^{-3}$ and compared this to the corresponding frequency in Cyg X-1. The motivation for using this frequency, instead of the break frequency, is that the break is often harder to measure and (in certain states at least) is known to change frequency in Cyg X-1. There are a number of underlying assumptions to this method. The first is that both the slope and the (relative) normalisation of the high frequency power spectra are similar for Seyferts  and Cyg X-1, but the frequencies scale inversely with the black hole mass. Implicit in this is the assumption that we are using comparable energy bands in the two types of source (the power spectrum is energy dependent). Furthermore, the original value of $f_c$ for Cyg X-1 used by Hayashida \\et (1998) was taken from an observation of the source in the low/hard state. This method therefore also assumes that Seyferts most closely resemble GBHs in the low/hard state, which may not be the case (Vaughan \\et 2003a; M$^{\\rm c}$Hardy \\et 2004). In the case of \\ngc\\ the frequency $f_c$ is not directly observed; the power spectrum lies well above $f \\times P(f) = 10^{-3}$ at all measured frequencies. Extrapolating the best-fitting broken power law model to higher frequencies gives an estimate of $f_c \\sim 7.2\\times 10^{-2}$~Hz (after accounting for the factor of two difference in the normalisation used by Hayashida \\et 1998). Scaling from the corresponding frequency in Cyg X-1, as defined by equation 4 of Hayashida \\et (1998) gives an estimate for $M_{\\rm BH} \\sim 6 \\times 10^4 ~ \\Msun$. This is comparable to the estimates made by the other methods although if, as discussed below, the normalisation of the power spectrum of \\ngc\\ is unusually high, this method will not be reliable.  \\subsection{A universal power spectrum?}  The early work on the high frequency power spectra of Seyfert galaxies found in favour of a universal power spectrum (Lawrence \\& Papadakis 1993; Green \\et 1993).  More recent work (Uttley \\et 2002; Markowitz \\et 2003)  suggests that while most Seyferts exhibit a similarly steep ($\\alpha \\approx 2$) power spectrum at high frequencies, the slope below the break may not be universal, and some objects may show a second break (NGC~3783; Markowitz \\et 2003) where others clearly do not (NGC~4051; M$^{\\rm c}$Hardy \\et 2004).  The extraordinary amplitude of the variations in \\ngc\\ suggest that the normalisation of the power spectrum may also vary quite substantially from object to object. A fractional rms in excess of $F_{\\rm var} > 1$ has been noted in only two other Seyferts  prior to \\ngc\\ (these are NGC~4051, Green \\et 1999 and IRAS 13224--3809, Boller \\et 1997).  NGC~4051 is also thought to contain a relatively low mass black hole (Shemmer \\et 2003).  At first sight it remains possible that the high variability amplitudes in both NGC~4051 and \\ngc\\ could be due to their power spectra being shifted to higher frequencies compared to Seyferts with more massive black holes (and, since the spectra are red, more variance is shifted into the observable frequency range) without the need for an intrinsically higher power spectrum normalisation.  The power at the break multiplied by the break frequency, $f_{\\rm br} P_{\\rm br}$, is indicative of the total power in the spectrum.  For \\ngc\\ this is $\\approx 5.4 \\times 10^{-2}$ in units of $({\\rm rms}/{\\rm mean})^2$. The six Seyferts studied by Markowitz \\et (2003) have values  in the narrow range $f_{\\rm br} P_{\\rm br} \\approx 1.6 - 3.0 \\times 10^{-2}$, all lower than \\ngc.  Papadakis (2004) estimated a `universal' value of $f_{\\rm br} P_{\\rm br} \\sim 1.7 \\times 10^{-2}$ for Seyfert galaxies (called ${\\rm PSD}_{\\rm amp}$).  This means that, even after accounting for the different regions of the power spectrum observed in different objects, the normalisation of the  power spectrum of \\ngc\\ is intrinsically high.  What could increase the overall amplitude of the variability (at all frequencies)? A plausible suggestion is that an enhanced accretion rate (relative to Eddington) somehow provides a more unstable accretion flow and hence stronger fluctuations.  The accretion rate, relative to the Eddington limit, can be estimated for \\ngc\\ using the available black hole mass estimates, together with the bolometric luminosity.  Moran \\et (1999) and Lira \\et (1999) estimate similar values of $L_{\\rm Bol} \\sim 5 \\times 10^{40}~\\erg~\\ps$.  Assuming a mass in the range $M_{\\rm BH} = 10^4 - 10^5 \\Msun$ this gives an Eddington fraction of $L/L_{\\rm Edd} \\sim 0.4 - 4$ per cent.  However, the mean unabsorbed $2-10$~keV luminosity during the \\xmm\\ observation was $L_{2-10} \\sim 1.5 \\times 10^{40}~\\erg~\\ps$.  Assuming $L_{\\rm Bol} \\gs 10 L_{2-10}$ would give $L_{\\rm Bol} \\gs 1.5 \\times 10^{41}~\\erg~\\ps$ and  this, less conservative, luminosity estimate gives an Eddington fraction of $L/L_{\\rm Edd} \\sim 1.2 - 12$ per cent\\footnote{ This assumes the Bolometric luminosity to be related to the mean X-ray luminosity by $L_{\\rm Bol} \\gs 10 L_{2-10}$. If however the peak X-ray luminosity was more representative of the overall  energy distribution then the luminosity could be revised upwards by a further factor of $\\sim 3$.  }.  Thus \\ngc\\ is most likely accreting at a `normal' rate; super-Eddington accretion seems unlikely to be the origin of the unusually large variability amplitude. Other possibilities include the black hole spin parameter, $a$, and the inclination of the putative accretion disc-black hole system, $i$. Unfortunately, these are more difficult to test."
    },
    "0410/astro-ph0410111_arXiv.txt": {
        "abstract": "We present a detailed $(U-V)$ color-magnitude relation (CMR) analysis for three local ($z<0.06$) clusters of galaxies. From square-degree imaging of the Abell clusters A85, A496, and A754, we select 637 galaxies down to $L\\sim0.1L^{\\star}_V$ with spectroscopic membership, which minimizes uncertain field contamination corrections.  To characterize the degree of CMR uniformity among nearby clusters, we use a maximum likelihood technique to quantify the CMR properties of the red galaxies in each cluster. We find that these clusters have similar CMRs with a mean color of $(U-V)=1.40$ at $M_V=-20+5\\log_{10}h$, and narrow limits of intrinsic color scatter [0.047,0.079] and slope [-0.094,-0.075]. If we restrict our analysis to the core cluster population of red galaxies, the resultant CMRs are in close agreement with that of the Coma Cluster, the only local cluster with a rest-frame $(U-V)$ CMR determination of comparable precision. Therefore, the CMR uniformity of present-day clusters spans a fairly wide range of cluster masses. We test how sensitive the CMR uniformity is to variations in color aperture size and sample selection, and we find at most slightly wider limits of scatter [0.047,0.112] and slope [-0.104,-0.054].  This upper limit for $(U-V)$ scatter is consistent with the bulk of the stellar populations of red cluster galaxies forming before $z=1.2$, with a maximum age spread of 5.2 Gyr.  In addition, we note that using colors from apertures containing equal fractions of galactic light does not remove the CMR slope and that none of the slopes exhibit a break as claimed by Metcalf et al.  Our findings expand the single Coma data point and provide a much-needed $z=0$ baseline for comparisons to high redshift cluster CMRs at the same rest-frame wavelengths.  The range in CMR scatter that we find among nearby clusters is consistent with the values reported for clusters at higher redshifts, further suggesting that there has been little evolution in the stellar populations of red-sequence cluster galaxies out to at least $z\\sim1$. To identify the most recently accreted cluster galaxies, we divide each cluster's membership into three galaxy populations based on $(U-V)$ color relative to the well-defined CMR. Blue and moderately blue galaxies make up $18-23\\%$ by number of each cluster population more luminous than $0.1L^{\\star}$. Our color-magnitude division should represent a rough time since cluster accretion. In testing this hypothesis, we find that blue galaxies are spatially, kinematically, and morphologically distinct from red cluster galaxies. Even in projection, the blue galaxies reside towards the outskirts of the cluster and appear to avoid the inner half megaparsec, in contrast with the increasing density of red sequence galaxies towards the cluster center.  In addition, the blue galaxies have velocity distributions relative to the cluster rest frame that are flatter, and some have different means, compared to the roughly Gaussian distribution of red member velocities. Members with the bluest colors tend to be disk-like or irregular in appearance compared to the red galaxies, which have mostly early-type (E/S0) morphologies.  Moderately blue cluster galaxies may be an intermediate mix with a fraction of small bulge-to-disk ratio S0s, yet these require closer scrutiny.  The spatial, kinematic, and morphological distinctions between blue and red cluster galaxies provide further evidence that CMR-relative color is related to time since cluster infall, and that bluer members are indeed the most recently accreted field spirals as expected in an hierarchical universe. ",
        "introduction": "\\subsection{Cluster Imaging} We observed clusters A85, A496, and A754 during two runs at the Kitt Peak National Observatory (KPNO) 0.9-meter Telescope with the NOAO Mosaic Wide-Field Imager \\citep{boroson94,muller98}.  We used the engineering-grade Mosaic for the $V$ data in November 1997, and the science-grade Mosaic with thinned CCDs for the $U$-band in January 2000.  We imaged each cluster in each passband cluster with a minimum of five exposures dithered a few arcminutes on the sky for the removal of inter-chip gaps, cosmic rays and other defects. The dither pattern we used ensured at least $80\\%$ of the maximal exposure for all regions of a final combined image from five exposures. The initial exposure of a dithered set was centered on the brightest cluster galaxy's position given in Table~1.  The Mosaic imagers provide a $59\\arcmin \\times 59\\arcmin$ field of view with an unbinned pixel scale of $0\\farcs423$ (15 $\\micron$ pixels).  This wide-field coverage spans physical scales of 1.6 (A496) and 2.7\\hmpc (A85 and A754). We had at least one photometric night per each Mosaic run during which we observed hundreds of \\citet{landolt92} standard stars for calibrating our cluster galaxy data to \\citet{johnson66} $U$ and $V$ magnitudes (see \\S \\ref{PhotSys}).  Our cluster observations are summarized in Table~2.  The KPNO 0.9-meter telescope is a Richey-Chretien design with an $f/3.3$ Cer-Vit primary and an $f/7.5$ secondary.  This $f/7.5$ system produces a curved focal plane; therefore, a 2-element fused silica field corrector is required for wide-field imaging with the Mosaic camera.  This imager employs large (5.75 inches square) par-focal $U$ and $V$ filters.  The Mosaic imaging CCD system has eight $2048$\\,$\\times$\\,$4096$ pixel Loral chips.  The $V$-band data were taken with the original unthinned, engineering-grade chips (hereafter MOS1), which result in lower quantum efficiency (QE) as shown in Figure~\\ref{FiltResp}.  For these observations the readout noise, dark current, and average single chip gain were $12.25 e^-$, $\\sim100$~to~$250  e^-$~hour$^{-1}$, and $6.64 e^-$~ADU$^{-1}$, respectively. The $U$-band data were obtained following an upgrade \\citep{wolfe98} in which the original CCDs were replaced with thinned, science-grade (AR coated SITe) chips (hereafter MOS2). Even with the improved sensitivity the QE of the science-grade detector falls off rapidly blueward of 4000\\AA\\ (Figure~\\ref{FiltResp}); therefore, the spectral response in the $U$-band is not a perfect match to standard \\citet{johnson66} $U$-band, which was obtained with a blue sensitive photomultiplier.  Nevertheless, the use of photometric standards with a limited range of $(U-V)$ colors allows calibration to a standard photometric system (as described in \\S \\ref{PhotSys}). The read noise, dark current, and average single chip gain for the $U$-band observations were $5.66 e^-$, $\\sim15 e^-$~hour$^{-1}$, and $2.86  e^-$~ADU$^{-1}$, respectively.  \\begin{figure} \\includegraphics[scale=.95, angle=0]{mcintosh_fig1.ps} \\caption{Average Mosaic response in the $U$ and $V$ passbands (solid lines). The dotted lines represent the average quantum efficiency (QE) of the eight Mosaic Loral CCDs for the thinned, science-grade ($U$-band, MOS2) and unthinned, engineering-grade ($V$-band, MOS1) cases.  The transmission of each filter is given by the dashed lines.  The total response is the combination of the filter transmission and QE.  The $U$-band response of our observations drops off rapidly blueward of $\\sim3750$~\\AA, similar to the Coma observations of \\protect\\citet{bower92a}. \\label{FiltResp} } \\vspace{-0.2cm} \\end{figure}  \\subsection{Data Reduction} \\label{Redux} For each cluster in our study we want to construct deep integration, wide-field $U$ and $V$ images devoid of bad pixels, cosmic rays, and gaps between CCDs. Achieving such high quality images requires stacking dithered Mosaic frames, which places high demands on the initial data reduction steps\\footnote{NOAO CCD Mosaic Imager User Manual (hereafter MosManual), Version Sept. 15, 2000, G. Jacoby; http://www.noao.edu/kpno/mosaic/manual/index.html.}. In particular, the data must be well-flattened and carefully corrected for photometric effects of the variable pixel scale. To this end we employ a customized reduction pipeline that uses the IRAF\\footnote{IRAF is distributed by the National Optical Astronomical Observatories, which are operated by AURA, Inc. under contract to the NSF.} environment and adheres to standard image reduction techniques.  We perform basic reduction of each individual Mosaic frame using the IRAF {\\sc mscred} package.  This software allows image processing to be performed on a multi-chip exposure as if it was a single CCD frame. For each single Mosaic exposure, we fit a first order Legendre polynomial to the overscan regions of the individual chip images, and then use the fits to correct (``debias'') each chip image.  Following debiasing we trim the chip images of their overscan regions.  We include a small correction ($<0.1\\%$ MOS1, $<0.3\\%$ MOS2) that is necessary due to cross talk between pairs of adjacent chips sharing the same electronics in the Mosaic detectors; i.e., for each $(i,j)$ pixel of chips \\#1,3,6,8 a predetermined fraction of the $(i,j)^{\\rm th}$ pixel value from its corresponding adjacent chip (i.e. \\#2,4,5,7) is subtracted.  We then subtract the average bias (using seven zero exposures) level from each Mosaic frame.  The $V$-band images from MOS1 require a dark correction due to significant dark current ($2-5$ ADU per chip per exposure). Roughly $3.1\\%$ ($0.4\\%$) of the MOS1 (MOS2) pixels are bad.  We flag bad pixels and include them in an image mask used during the final image combination. In addition, a large portion ($1000\\times1000$~pixel) of the upper northeast corner of MOS1 chip\\#5 appears to be adversely affected by many hot pixels with variable dark current.  We remove this region as we cannot correct for it adequately.  We note that the removal of this bad region represents only $1.6\\%$ of the total area covered by the $V$-band imaging, resulting in the loss of roughly 2-4 galaxies per cluster at large cluster-centric distance.  This loss does not affect our results.  An important step towards achieving precise photometry is the determination and removal of the response function of the individual CCDs -- i.e. flat fielding the data. Traditionally, dividing each exposure by a uniformally illuminated blank frame will produce an image that has a uniform and flat appearance.  Yet, the Mosaic imagers have pixel scales that decrease roughly quadratically such that an individual pixel in a field corner is 6\\% smaller, and contains only 92\\% of the flux, compared to a pixel at field center (see MosManual for details).  Therefore, although an individual star anywhere on the image will have the same number of photons within the point-spread-function (PSF), the variable pixel scale causes the photometric zero point to vary by 8\\% over the field of view.  We correct for this photometric effect following the recommendations given in the MosManual. Briefly, we flatten each image with a flat-field frame that has {\\it not} been corrected for the variable pixel scale.  Then, following our astrometric calibration and prior to stacking multiple exposures, we re-grid each frame to a tangent-plane projection with pixels of constant angular scale. We note that during the re-gridding, we {\\it do not} scale each pixel photometrically by the amount each pixel area has changed.  In this manner we account for the variable pixel scale and produce uniform images over the entire field of view.  We construct flat-field frames for each passband using a combination of twilight and night-sky flats.  First we make a normalized flat with high S/N by averaging a set of twilight illuminated exposures. This component accounts for both the small scale, high frequency (pixel-to-pixel) variations in response, and the spatial variations over large fractions of an image.  We then fit a smooth surface to a night-sky flat produced by median combining a set of unflattened cluster frames with all objects masked out, and we multiply this smooth surface by the twilight flat frame to produce a high S/N flat with the same spectral response as our science observations. We iterate the flat field construction twice to optimize the night-sky flat object masking.  Thus, the resultant ``super flat'' for each passband is spatially and spectrally flat. We divide each cluster and standard star image by its super flat to achieve $\\lesssim1\\%$ deviations from global flatness over the eight chip array.  Good astrometry is required to register and stack each dithered set of exposures.  In addition, accurate celestial coordinates are necessary to identify galaxies with measured redshifts. All Mosaic images have an initial default world coordinate system (WCS) loaded in their header at the time of observation.  The WCS maps the image pixel space onto celestial coordinate axes (RA and Dec).  However, effects such as global pointing offsets, instrument rotations, and differential atmospheric refraction produce the need for corrections to this astrometric calibration. We use the {\\sc msccmatch} package in IRAF to interactively derive accurate (RMS~$\\lesssim 0.3\\arcsec$) astrometry for each dithered exposure by registering $\\sim300$ stars (with $12.0<V<16.5$) distributed evenly over each image to their epoch J2000.0 equatorial coordinates in the USNO-v2.0 system \\citep{monet96}. We then remap the eight chip exposures for each Mosaic frame onto a single image by rebinning the pixels to a tangent-plane projection, thus producing pixels of constant angular size (as described above). At this point our standard star frames are fully reduced.  Before combining a dithered sequence of registered exposures into a single high S/N image, we must account for the effects of differing transparency between individual frames.  Transparency differences are due to time varying effects during a dither sequence. The most common effect is the changing airmass over an hour long dithered exposure.  Other effects, resulting in variable photometric conditions (e.g. thin cirrus, clouds) can change the atmospheric transparency. First we subtract a constant flux level equal to the mode sky value from each image in a dithered set.  Higher order terms are unnecessary due to the better than 1\\% global flat fielding.  Next all exposures of a dithered sequence are scaled to a reference image with the lowest airmass and/or best photometric conditions. We calculate the multiplicative scale factor for each image by comparing simple aperture flux measurements for the set of $\\sim300$ astrometric calibration stars common between each image and its reference image.  Scaling factors have typical values within a few percent of unity. Therefore, the set of precombined cluster frames are scaled to the same effective airmass and exposure time. We note that variable airmass conditions can affect color, but we find the relative airmass and $(U-V)$ color dependence to be negligible (see \\S~\\ref{PhotSys}).  Finally, we median combine each dithered set of scaled and sky-subtracted cluster images.  We use median combination, rather than a cosmic ray rejection algorithm, to produce final images free of most cosmic rays.  We restore the sky level to the average background calculated from the modes of each image multiplied by their corresponding transparency scale factors.  In Table~3 we tabulate the relevant properties of the final cluster images.  \\subsection{Photometric System} \\label{PhotSys} We transform our cluster galaxy photometry to the \\citet{landolt92} system.  We selected photometric flux standards with colors similar to the range of colors observed in typical cluster galaxies at low redshift $0.0\\leq (U-V) \\leq2.1$, and which have been observed more than twice by \\citet{landolt92} to ensure photometric consistency. During nights with apparently photometric conditions we observed $\\sim100$~to~$200$ photometric standard stars in each passband. To achieve good photometric calibration, we observed standard fields at least three times during the night at relative airmass readings in the range $1.2\\leq X \\leq2.0$.  The photometric solution is based on a simple, linear transformation equation that relates standard magnitudes to instrumental magnitudes in each band pass. The total flux $f_{\\rm ap}$ from an exposure of duration $t_{\\rm exp}$ within a circular aperture of $14\\farcs4$ diameter, similar to that used by \\citet{landolt92}, is converted to an instrumental magnitude \\begin{equation} m_{\\rm instr} = -2.5 (\\log_{10}{f_{\\rm ap}} - \\log_{10}{t_{\\rm exp}}) . \\label{InsMag} \\end{equation} Then we use IRAF's {\\sc photcal} package on the selected set of standard star $m_{\\rm instr}$ measurements to find the best-fit solutions to the following transformation equations: \\begin{equation} V_{\\rm L92} = m_{\\rm instr} + zp_V + \\alpha_V X_V + \\beta_V (U-V)_{\\rm L92} \\label{VTrnsfEqn} \\end{equation} \\begin{equation} U_{\\rm L92} = m_{\\rm instr} + zp_U + \\alpha_U X_U + \\beta_U (U-V)_{\\rm L92} . \\label{UTrnsfEqn} \\end{equation} The photometric system is defined by the coefficients ($zp$, $\\alpha$ and $\\beta$) of the transformation equation in each passband.  The photometric zero point $zp$ quantifies the gain and the total sensitivity of the telescope plus detector.  The airmass term $\\alpha$ is a measure of the atmospheric extinction as a function of telescope altitude. The color term $\\beta$ shows how well the instrumental system matches the \\citet{landolt92} system.  A color dependence on airmass is known to exist especially at bluer wavelengths; however, through testing we find no need for an airmass-color cross term in the transformation equation for our photometric calibrations.  The observed standards primarily fall on the central chips (i.e. \\#2,3,6,7); however, recall that we obtain the instrumental magnitudes from fully reduced images that have been corrected to have a globally uniform zero point (see \\S \\ref{Redux}).  We solve the corresponding transformation equation for each nightly set of standards in a given passband and present the coefficients and their formal errors in Table~4. We note that the apparent systematic difference of 0.10 mag between the two $V$-band calibration zero points in fact amounts only to a 0.02 mag difference when the airmass terms are included. The photometric accuracy ($\\sim3\\%$ error in $zp$) of this data is close to optimum given the remapping of the eight chips.  In this case, the dominant source of the photometric uncertainty is due to CCD-to-CCD sensitivity differences (MosManual).  With the above photometric system and a flux measurement $f$, the magnitude for each galaxy is \\begin{equation} m = zp_{\\rm eff} - 2.5\\log_{10}{f} , \\label{mageq} \\end{equation} where the effective zero point $zp_{\\rm eff} = zp + \\alpha \\tilde{X} + 2.5\\log_{10}{\\tilde{t}_{\\rm exp}}$ accounts for the effective exposure time and airmass of each final image as given in Table~3.  \\subsection {Selection of Cluster Members}  \\subsubsection{Source Detection and Deblending} Given the large numbers of sources in each square-degree image, we use the source detection and extraction software SExtractor \\citep{bertin96} to automatically construct catalogs of accurate source positions from the $U$ and $V$ final frames. In addition, SExtractor provides instrumental magnitudes \\\\ (MAG\\_BEST) and other photometric parameters for each detected source.  We configure SExtractor to detect objects comprised of a minimum of 5 connected pixels (DEBLEND\\_MINAREA) above a background threshold of $3\\sigma_{\\rm bkg}$ (DETECT\\_THRESH). Overlapping sources are deblended into multiple objects if the contrast between flux peaks associated with each object is $\\geq0.05$ (DEBLEND\\_MINCONT). These parameters provide our working definition of an imaged source.  We confirm that these parameters provide good source detection and deblending by visually inspecting random regions from each image.  We reject sources flagged (FLAGS$\\ge4$) as saturated, or otherwise bad, either of which will corrupt the deblending and extraction routines of SExtractor. In addition, we exclude sources within 140 pixels (1 arcmin) of image edges; these regions are lower in S/N due to the dithering of the individual frames. For each catalog we determine an empirical magnitude limit $m_{\\rm min}$ where the source counts distribution flattens and begins to fall off.  We summarize the magnitude limits and number counts for our final $U$ and $V$ cluster source catalogs in Table~5.  \\subsubsection{Star/Galaxy Separation} \\label{stargal} Owing to the variable PSF as a function of image position, stellar profiles can be extended especially towards the image edges.  Moreover, small circular galaxies are unresolved especially at faint magnitudes. Therefore, we do not rely on the SExtractor star/galaxy classifier (CLASS\\_STAR) to separate stars and galaxies.  Instead, we use GIM2D \\citep{simard02} to fit a PSF-convolved, bulge$+$disk model to the two-dimensional (2D) surface brightness profile of all sources in the $V$-band catalog.  We select the $V$-band catalog due to its higher S/N and better seeing characteristics.  GIM2D provides the model half-light (or effective) radius \\rhl and the total flux $f_{\\rm tot}$, among other structural parameters.  In Paper~3 we present a detailed description of our galaxy profile fitting technique and the complete catalog of structural properties measured from our cluster data.  In this paper we use the fluxes to calculate total $V$-band magnitudes and the half-light size to separate stars and galaxies. Though stars are not fit well by a two component model, the PSF convolution produces $r_{\\rm hl}<1$~pixel model profiles for most stars.  In contrast, galaxies brighter than $V=19$ have \\rhl significantly larger than the seeing size. Thus, our method provides robust star/galaxy separation to $V\\leq19$, one magnitude fainter than the limit where our $U,V$ photometry and redshift matching are quite ($\\ge95\\%$) complete. With this method we find a total of 4320 extended sources (i.e. galaxies) among the three clusters, of which 1469 are brighter than  $V=19$ mag. We cross-correlate the $V$-band selected galaxies with the magnitude-limited $U$ sources to produce a catalog of 1315 total galaxies with $U,V$ photometric data and accurate coordinates from the three cluster fields.  We note that our cross-correlation with $U$-band data at similar limiting depth as the $V$-band creates a color bias such that at the faintest magnitudes only increasingly bluer sources are matched. Therefore, as shown in Figure \\ref{zuvCts} we find $U$-band detections for 1093 (74\\%) of the $V\\leq19$ galaxies; while we successfully match 95\\% (727/765) brighter than $V=18$ mag. We give the source counts breakdown for each cluster in Table~5.  \\subsubsection{Total Magnitudes} \\label{totmags} We use GIM2D model fluxes from profile fits to the $V$-band imaging to calculate the total magnitudes (equation~\\ref{mageq}) for all $V$-selected extended sources.  We present the $V_{\\rm tot}$ distribution for the combined cluster (A85$+$A496$+$A754) galaxy sample in Figure~\\ref{zuvCts}. GIM2D estimates parameter uncertainties through full Monte-Carlo propagation of the parameter probability distributions during each best-fit model determination.  Typical mean errors in $V_{\\rm tot}$ are 0.04 mag, which is significantly larger than the $V$-band Poisson errors of cluster members ($\\le0.01$ mag). In Figure~\\ref{magcomp} we compare the GIM2D-based total magnitudes against the SExtractor (MAG\\_BEST)\\footnote{SExtractor automatically selects between several photometric measurements based on degree of influence from neighboring sources.} derived values for all galaxies with $U$ and $V$ detections. GIM2D $V_{\\rm tot}$ magnitudes are 10\\% brighter on average than SExtractor determinations.  The GIM2D total flux measurement includes light from the LSB regions at large galactic radii, whereas SExtractor magnitudes are aperture-based and thus underestimate the total flux systematically when large isophotal thresholds, such as $3\\sigma_{\\rm bkg}$, are used \\citep{bertin96}.  We emphasize here that we do not use SExtractor magnitudes in any of our analysis.  Throughout this paper we use GIM2D-derived $V_{\\rm tot}$ magnitudes to calculate rest-frame absolute magnitudes \\begin{equation} M_V - 5\\log_{10}h = V_{\\rm tot} - DM - A_V + k_V + C , \\label{TotMagEqn} \\end{equation} where $DM$ is the cosmological distance modulus if $h=1$, $A_V=3.315\\times E(B-V)$ is the Galactic extinction correction following \\citet[][hereafter SFD98]{schlegel98}, $k_V$ is the $k$-correction from \\citet{poggianti97}, and $C$ is a minor photometric color correction (see \\S~\\ref{colors}) amounting to roughly $+0.01$ mag on average. These magnitudes have random uncertainties of $\\approx0.02-0.07$ mag, with the dominant source of random error being the uncertainty in the GIM2D model flux.  We present $M_V$ measurements with formal random errors for a small sample of cluster members in Table~6; the full catalog is available electronically. In addition, sources of systematic error in $M_V$ include uncertainty in the overall photometric zero point (0.02 mag), the relative extinction zero point (0.07 mag), and $k$-correction (0.01 mag).  We use the mean cluster redshift $z_{\\rm clust}$ to calculate the distance moduli for individual cluster members, which removes the possibility of introducing magnitude uncertainties due to galaxy motions through the cluster.  Assuming each cluster is a sphere with a characteristic radius of $1.5$~\\hmpc (consistent with an Abell radius), we estimate a $\\sim2\\%$ error in our adopted distance moduli due to the uncertainty of front-to-back cluster distance. We present each cluster's mean redshift and resulting cosmological distance modulus in Table~7.  We correct the galaxy magnitudes for the effects of Galactic extinction along the line of sight using the dust maps of SFD98, which provide reddening $E(B-V)$ values as a function of Galactic coordinates $(l,b)$. The average corrections are 0.12 (A85), 0.42 (A496), and 0.21 (A754) $V$ magnitudes. The correction for A496 is particularly large due to the presence of a Galactic molecular cloud along the line of sight (Finkbeiner, D.\\ 2001, private communication). SFD98 adopt a systematic uncertainty in extinction correction of 0.020~mag in $E(B-V)$ (i.e. $\\sim0.07$ mag in $M_V$), which represents the median difference between their reddening estimates from $100\\micron$ dust maps and the 21cm gas maps of \\citet{burstein82}.  In addition, SFD98 give a formal uncertainty of $10\\%$ in $E(B-V)$, which is 0.01 (A85), 0.04 (A496), and 0.02 (A754) in $M_V$.  Even at the low redshifts ($z<0.06$) of our cluster galaxies, the observed shift of the rest-frame spectral energy distributions is a significant effect, making galaxies appear redder and dimmer, requiring a $k$-correction. We use the elliptical (E-model) $k$-corrections from \\citet{poggianti97}, which amount to mean corrections of -0.06 mag (A496) and -0.10 mag (A85,A754) to the rest-frame $V$-band. At $z<0.06$, the three basic Poggianti models (E,Sa,Sc) give $k$-corrections that differ by $<1\\%$.  We adopt 0.01~mag from the model-dependent difference as an estimate for the systematic uncertainty in the \\citet{poggianti97} $k$-corrections.  \\subsubsection{Confirmed Member Galaxy Catalogs} \\label{confmem} By analyzing only known cluster members we remove the need for uncertain foreground/background statistical corrections. \\citet{christlein03} analyzed 1486 recessional velocities along the line of sight towards the three clusters we study here. These redshifts were obtained from galaxy spectra in $1.5\\times 1.5\\arcdeg$ multi-fiber fields with central coordinates similar to our cluster images; therefore, about 35\\% of the spectroscopic data are outside of our imaging.  We select member galaxies based on mean recessional velocity $\\left<cz\\right>_{\\rm clus}$ and internal velocity dispersion $\\sigma_{\\rm clus}$ measurements for each cluster from \\citet{christlein03}. We define cluster members as those galaxies with recessional velocities $cz_i=\\left<cz\\right>_{\\rm clus} \\pm 3\\sigma_{\\rm clus}$, where the cluster parameters are given in Table~7. There are 971 galaxies with redshifts that have coordinates within the regions bounded by our cluster imaging, of which 721 are cluster members and 250 fall outside of our membership definition. We cross-correlate the coordinates of galaxies from the redshift data with our $U,V$ source positions from our imaging catalog and achieve a total of 793 image/redshift matches (637 members and 156 non-members). We define image/redshift matches to be the nearest within $5\\arcsec$ and find the mean coordinate separation is $<2\\arcsec$. The image/redshift membership breaks down into 180 (A85), 146 (A496), and 311 (A754) galaxies; we present the relevant information for each cluster membership in Table~7. The final sample contains a total of 637 spectroscopic cluster member galaxies with a large range of absolute magnitudes ($-16.5 < M_V < -23.0$).  The $U,V$ C-M data for A85, A496 and A754 are comparable in depth and in membership ($N=275$) to the most comprehensive study of the Coma cluster by \\citet{terlevich01}.  Within the region of our imaging, 84 spectroscopic members have no $U,V$ matched image counterparts, of which only 12 are brighter than $V=18$ mag and 72 are fainter than this limit.  Therefore, the completeness for our galaxy image/redshift cross correlation for cluster members is 97.6\\% (494/506) at $V\\leq18$ mag.  Nearly half (5/12) of the incompleteness at $V\\le18$ is due to $>5\\arcsec$ coordinate mismatch between the spectroscopy and imaging catalogs, while two more have imaging that is contaminated by nearby saturated stars.  The remaining five are comprised of two E/S0s and three spirals.  The fainter ($V>18$) unmatched redshifts are missed for several reasons:  24\\% (17/72) with $\\theta_{\\rm sep}>5\\arcsec$; 7\\% (5/72) contaminated by nearby saturated stars; 49\\% (35/72) are even fainter ($V>19.5$ mag) and/or are spheroidal in appearance and likely red in color, which in turn drops them out of our $U$-band limited catalog of $U,V$ matches; and 21\\% (15/72) have disk morphologies in the $V$-band, either LSB or edge-on, in each case also not detected in our $U$-band imaging.  Additionally, there are 94 non-cluster redshifts not detected in both $U$ and $V$ frames. All but five are background sources, thus, we assume these have been missed due to red $(U-V)$ color making them fainter than our $U$-band cut.  \\subsubsection{Sample Completeness} For the combined cluster sample in Figure~\\ref{zuvCts}, we plot the distribution of galaxy counts per $V$ magnitude for sources with spectroscopic identification (members and non-members) in comparison with our imaging catalog counts for $U,V$ matched and $V$-band extended sources. This illustrates the completeness of our matched $U$ and $V$ extended source detections (95.0\\% for $V\\le18$).  We have shown in the previous section that for $V\\le18$ spectroscopic members we have a 97.6\\% complete sample of $U$ and $V$ image sources. Finally, we have hundreds of unidentified (i.e. no spectroscopic counterparts) extended sources with $U,V$ photometry towards each cluster.  At $V\\le18$ there are 112 imaged extended sources without redshifts. These represent the expected fractional incompleteness of the spectroscopic sample \\citep{christlein03} and provide a direct measure of the completeness of our cluster galaxy observations, 84.6\\% (615/727) for $V\\le18$.  \\subsubsection{Morphological Classification} \\label{mmakeup} Using the $V$-band images we classify cluster members into three coarse visual types: (1) elliptical, (2) S0, and (3) spiral or irregular.  For the bulk of our analysis we will concentrate on the subset of 546 members brighter than $M_V=-18.1+5\\log_{10}h$ ($0.1 L^{\\star}$). This magnitude limit is within 0.1~mag of the relatively complete cutoff of $V=18.0$ for A85 and A754, and represents an extremely complete sample for A496 at $V=16.9$ mag. We use the isophotal contours of each galaxy to distinguish between ellipticals, which have smooth radial profiles, and lenticular (S0) galaxies with separate bulge and disk morphological components characterized by an intensity discontinuity \\citep{dressler80b}.  Most bright disk galaxies belonging to these clusters have smooth appearances with no spiral structure, as such we classify all featureless disks as S0s, regardless of bulge-to-disk ratio\\footnote{Yet, in the Hubble sequence S0 galaxies have more prominent bulges relative to spirals \\citep[e.g.][]{dressler80}.}.  We note that the lack of spiral structure may be a resolution issue.  Our typical $V$-band seeing is $1.2\\arcsec$ FWHM corresponding to physical resolutions of 910 pc (A85, A754) and 540 pc (A496); therefore, while we resolve the cluster members, galaxies with weak spiral arms may look smooth at the distances of these clusters.  We account for this possible bias in our quantitative analysis of morphology in Paper~1. Our purpose here is to provide basic visual morphology information for our analysis of different cluster populations based on well-defined color criteria, and we present many examples in Figures \\ref{stamps85}-\\ref{stamps754}. We discuss the relative ratios of different morphological types as a function of galaxy color in \\S \\ref{CMmorphs}.  In Paper~1, we address differences in quantitative measures of morphology, such as bulge-to-disk ratio and disk substructure, among different color-based galaxy populations from these clusters.  In column (7) of Table~7 we give the fraction of basic morphological types among cluster galaxies brighter than $0.1L^{\\star}$.  These morphological ratios are in agreement with the bulk of those from global cluster populations studied by \\citet{dressler80}. The overall numbers of bright E, S0, and S$+$Irr galaxies in these clusters are $163\\pm33$, $311\\pm87$, and $72\\pm14$, respectively, which correspond to the relative fractions of 0.30/0.57/0.13 (E/S0/S+Irr) morphological types from the total bright cluster membership.  \\begin{figure}[t] \\center{\\includegraphics[scale=1, angle=0]{mcintosh_fig2.ps}} \\caption{Extended source counts in our cluster images as a function of total magnitude $V_{\\rm tot}$ from GIM2D fits to the $V$-band imaging.  For the combined cluster sample we plot the number counts distribution for 4320 extended sources from our $V$-band limited catalog (solid outline), for 1315 sources with $U$- and $V$-band photometry (binned histogram), and for 793 galaxies with spectroscopic redshifts matched to extended source coordinates from our $U,V$ imaging (solid histogram).  Among the 793 $U,V$ sources with redshifts, 637 are cluster members and 156 are non-member, mostly background, galaxies.  The $V$-band extended source distribution has been clipped to resolve the level of completeness in the imaging and spectroscopy.  At $V\\le18$ we have $U$-band photometry for 95\\% (727/765) of our $V$-band detections, and we find that 85\\% (615/727) of these have redshifts.  Furthermore, within our imaging 97.6\\% (494/506) of cluster member redshifts brighter than $V=18$ have $U,V$ source counterparts. \\label{zuvCts} } \\vspace{-0.2cm} \\end{figure}  \\begin{figure*} \\center{\\includegraphics[scale=.8, angle=0]{mcintosh_fig3.ps}} \\caption{Comparison of GIM2D total $V$-band magnitudes with SExtractor MAG\\_BEST measurements for the complete sample of cluster galaxies with $U$ and $V$ detections.  The total number of galaxies from each cluster image are noted in each panel (see Table~5).  The lower panels show the magnitude difference (GIM2D-SExtractor). The bold line represents the mean magnitude differences for each cluster: $-0.087\\pm0.004$ (A85), $-0.083\\pm0.009$ (A496), and $-0.118\\pm0.004$ (A754). The scatter in $\\Delta_V$ for each cluster is, from left to right: 0.08, 0.16, and 0.09 mag. \\label{magcomp} } \\vspace{-0.2cm} \\end{figure*}  \\subsection{Rest-frame $(U-V)$ Galaxy Colors} \\label{colors} We require precise $(U-V)$ color measurements for every cluster galaxy. With these we will construct C-M diagrams and determine empirical CMRs for each cluster.  In this section we describe our cluster galaxy color measurements including our choice of aperture size, and the necessary photometric corrections and their corresponding uncertainties.  The effects of the variable PSF over each individual Mosaic image, plus differing seeing conditions between $U$ and $V$ observations, combine to adversely affect the image quality of each individual galaxy by different amounts.  We account for varying image quality by degrading all $U$ and $V$ galaxy imaging to a common PSF size for each cluster: 2.5 (A85), 2.0 (A496), and 2.5 (A754) arcseconds (FWHM). This correction insures that we measure galaxy aperture magnitudes for colors in a consistent manner. We use {\\sc daophot} to construct a variable PSF model for each cluster mosaic from hundreds of stars distributed evenly over the image. We obtain the PSF size (FWHM) at each galaxy's location in our imaging from the PSF model.  The image quality characteristics of each combined cluster frame varying only a few percent except near the outer image edges where focal plane distortions start to dominate. The typical PSF size for the images is roughly $1\\farcs3$ FWHM, with a tail of larger values.  We degrade each image to the largest PSF size using a Gaussian smoothing kernel.  While some information may be smeared out in the bulk of galaxies found away from mosaic edges, our procedure produces uniform galaxy images for consistent aperture color measurements free of systematic biases due to different PSF sizes.  We note that the total integrated flux we use to derive absolute magnitudes are derived from PSF-convolved fits to surface brightness profiles, which require no image quality correction.  To calculate the $(U-V)$ color of each galaxy we measure aperture magnitudes from the $U$ and $V$ images smoothed to a common PSF size. We select a fixed aperture size to directly compare our $(U-V)$ color data with that of Coma from \\citet[][hereafter BLE92]{bower92a,bower92b}. For Coma ($cz = 6500$~\\kms), BLE92 used an $11\\arcsec$ aperture diameter\\footnote{This size corresponds to their $13\\arcsec$ aperture with a small correction applied to match the same physical size as the $60\\arcsec$ aperture they used for Virgo cluster galaxies \\citep[see][]{bower92a}.}.  We scale the BLE92 aperture by the ratio of the line-of-sight comoving distance to Coma and to each of our clusters to obtain comparable fixed aperture diameters of $4.45\\arcsec$ (A85), $7.42\\arcsec$ (A496), and $4.50\\arcsec$ (A754). Therefore, the fixed aperture we use to measure member $(U-V)$ colors in each cluster samples the same physical diameter of $3.34$~\\hkpc in a $\\Lambda$-CDM ($\\Omega_{\\rm M}=0.3$, $\\Omega_{\\rm \\Lambda}=0.7$) cosmology.  We measure aperture fluxes using the {\\sc apphot} package in IRAF. This package is the standard for performing aperture photometry on uncrowded digital images. To ensure the sky is measured well outside each galaxy we use an annulus with a width of 5 pixels and an inner radius at least six times larger than the galaxy's half-light radius. Our typical sky annulus contains $1000-5000$ pixels, with a minimum of $\\sim300$ pixels for the smallest galaxy images.  The local sky in our images is quite flat given the steps we have taken to make the images globally flat (\\S \\ref{Redux}). For each final cluster image we calculate the mean sky level in 48 separate $30\\times 30$ pixel boxes, free of celestial objects and bad pixels, and find flat-field uncertainties of $0.8-1.1\\%$ in $V$, and $0.9-1.6\\%$ in the $U$-band.  We transform the aperture magnitudes to the \\citet{landolt92} photometric system by solving the following color-correction equations iteratively: $V_i = V + \\beta_V (U_{i-1} - V_{i-1})$ and $U_i = U + \\beta_U (U_{i-1} - V_{i-1})$, where $U$ and $V$ are the uncorrected aperture magnitudes from equation~(\\ref{mageq}), and the $\\beta$ coefficients are color calibration terms given in Table~4.  For the initial iteration we set $(U_{i-1} - V_{i-1}) = (U-V)$.  We perform these calculations until the magnitude difference between successive iterations is $\\delta m \\leq 0.001$~mag.  The resultant average color corrections are -0.04 (A85), -0.02 (A496), and -0.04 (A754) mag.  Using Galactic reddening values from SFD98 we apply a $(U-V)$ extinction correction of $A_{U-V}=2.119\\times E(B-V)$ to each cluster galaxy $(U-V)$ color. The mean corrections are 0.08 mag (A85), 0.27 mag (A496), and 0.14 mag (A754). Finally, we apply $k$-corrections of -0.05 mag (A496) and -0.08 mag (A85,A754), on average, from \\citet{poggianti97} to produce the final rest-frame $(U-V)$ colors.  We note that biases are associated with the choice of aperture size used to measure the colors of galaxies.  Fixed circular apertures do not take into account the different physical sizes of galaxies, which may affect the measured CMR because brighter galaxies have larger sizes \\citep{scodeggio01} and color gradients are anticipated in most types of galaxies \\citep[see][and references therein]{peletier90,dejong96}. Therefore, we test the effects that three different aperture sizes (i.e. fixed vs. multiples of the half-light radius $r_{\\rm hl}$) have on the measured CMR properties of each cluster in Appendix \\ref{apertures}.  Our $(U-V)$ color measurements have formal random errors of $\\approx0.02-0.05$ mag.  The Poisson noise in the aperture magnitudes is the dominant contribution, while additional sources include the color correction, flat fielding, and dust ($10\\%$ in $E(B-V)$) uncertainties. All cluster members are bright enough that errors due to dark current and read noise are not significant. We present galaxy colors from the three aperture choices for a small sample of cluster members in Table~6; the full catalog is available electronically. In addition, there are $(U-V)$ systematic error sources from the overall photometric zero point (0.04 mag), the relative extinction zero point (0.04 mag), and the $k$-correction (0.01 mag).  \\begin{figure}[hp] \\center{\\includegraphics[scale=.73, angle=0]{mcintosh_fig4.ps}} \\caption {Rest-frame $U,V$ C-M diagrams for the total sample of 637 spectroscopic members in clusters A85, A496, and A754.  The colors and total absolute magnitudes are fully corrected (see text for details). The colors are from fixed aperture magnitudes with $3.34$~\\hkpc physical diameters.  The systematic photometric uncertainties are given at lower left, and the random errors for each galaxy are shown with $\\pm1\\sigma$ errorbars.  For each cluster we plot the fit to the CMR of E/S0s in Coma from BLE92 (thin solid line), and we include the $\\pm1\\sigma$ intrinsic scatter of 0.050 mag (dashed lines). The red-sequence galaxies in clusters A85 and A754 are well-fit by the BLE92 result.  Cluster A496 is also consistent with the BLE92 fit when the 0.11 mag systematic error (large bold error bar at $U-V=1.60$) in CMR zero point is taken into consideration. \\label{TheCMDs_fixap} } \\vspace{-0.2cm} \\end{figure} ",
        "conclusions": "We present a large data base of precise $U,V$ photometry for 637 present-day cluster galaxies, most more luminous than $M^{\\star}_V+3$.  The sample is drawn from square-degree imaging of three nearby ($z<0.06$) Abell clusters: A85, A496, and A754.  All galaxies have cluster membership confirmed using spectroscopic redshifts from \\citet{christlein03}, which removes the need for uncertain field corrections.  To determine the degree to which cluster $(U-V)$ CMRs are uniform in the local universe and thereby constrain galaxy evolution models, we employ a maximum-likelihood technique to analyze the CMR properties of the red galaxies within a projected radius of $1/3R_{200}$ for each cluster in our sample. Using a variety of red galaxy sample selections and $(U-V)$ color measurements, we find that CMRs (defined by slope, intrinsic scatter, and zero point at fixed absolute magnitude $M_V=-20+5\\log_{10}h$) remain quite uniform, with maximum variations in these properties as follows: 0.047 to 0.112 mag (intrinsic scatter); -0.104 to -0.043 (slope); and 1.34 to 1.45 mag (zero point).  We note that the range in zero point color can be explained by systematic photometric error, thus, our results are consistent with homogeneous CMR colors in agreement with \\citet{andreon03}. The range of CMR scatter allows a maximum spread in the formation times of the stellar populations of 5.2 Gyr, with the bulk of the stars forming before $z=1.2$, following the analysis of \\citet{bower98}.  If we limit our analysis to the core red galaxy populations and adopt the same color aperture as used for the Coma Cluster by BLE92, we find that local clusters spanning a range of cluster masses and richnesses have universal CMRs within tighter limits of intrinsic scatter [0.047,0.079] and slope [-0.094,-0.075]. We find that the CMR universality is robust to changes in sample size and sampling radii. Moreover, we show that CMR uniformity extends to $M^{\\star}_V+3$ and we find no marked change in the slope of the $(U-V)$ relation in contrast with the study of \\citet{metcalfe94}.  Last, we find that CMRs based on colors measured in apertures containing the same fraction of light are inconsistent with flat (i.e. zero slope) as claimed by \\citet{scodeggio01}.  The CMR results that we present here provide an useful benchmark for comparisons with CMRs studied in high redshift clusters using similar rest-frame colors. The variation in CMR scatter we have observed among four nearby Abell clusters \\citep[the three in our sample plus Coma; BLE92][]{terlevich01} is consistent with the values reported in the literature at higher redshifts. This lack of evidence for evolution of the $(U-V)$ CMR scatter suggests, as the narrow CMR scatter among our three clusters does, that the bulk of stars in the red-sequence galaxies formed above $z\\sim1$.  We use the CMR, spatial, kinematic, and morphological data to identify the population of galaxies most recently accreted by the clusters. We first divide each cluster sample into three populations based on $(U-V)$ color relative to the well-defined CMR. The colors of cluster galaxies are likely to depend on their time since arrival \\citep{balogh00}; thus, bluer relative color should correspond to later accretion. We establish the existence of significant numbers ($18-23\\%$) of blue and moderately blue galaxies among the luminous ($\\ge0.1L^{\\star}$) residents of these Abell clusters.  The blue members have spatial, kinematic, and morphological properties that are inconsistent with their red cluster neighbors.  In general, we find bluer galaxies prefer the cluster outskirts, avoid the inner regions, have flat and off-center velocity distributions, and have late-type morphologies.  These properties provide compelling evidence confirming the more recent arrival nature of members with very blue colors relative to the well-defined CMR.  Moreover, the members with intermediate colors are possibly objects in ``transition'' between more recent arrival VBGs and long-member RSGs. The existence of late cluster arrivals provides the best present-day location to look for evidence of environmental-induced galaxy evolution in cluster environments.  To further comment on how the most recent arrival populations fit into the hierarchical picture of cluster galaxy evolution requires more detailed scrutiny (profile model fitting with proper accounting of selection effects) of the blue membership and direct comparison with blue field galaxies, all of which we address in Paper 1."
    },
    "0410/astro-ph0410327_arXiv.txt": {
        "abstract": "Using Monte Carlo codes, we follow the collisional evolution of clusters in a variety of scenarios. We consider the conditions under which a cluster of main sequence stars may undergo rapid core collapse due to mass segregation, thus entering a phase of runaway collisions, forming a very massive star (VMS, $M_\\ast > 1000\\,{\\rm M}_\\odot$) through repeated collisions between single stars. Although collisional mass loss is accounted for realistically, we find that a VMS forms even in proto-galactic nuclei models with a high velocity dispersion (many 100\\,km\\,s$^{-1}$). Such a VMS may be a progenitor for an intermediate-mass black hole ($M_\\bullet\\ge 100\\,{\\rm M}_\\odot$). In contrast, in galactic nuclei hosting a central massive black hole, collisions are found to be disruptive. The stars which are subject to collisions are progressively ground down by high-velocity collisions and a merger sequence appears impossible. ",
        "introduction": "In a stellar cluster, if the velocity dispersion is smaller than about 300\\,km\\,s$^{-1}$, the collision cross section is dominated by gravitational focusing. Under these conditions the (local) collision time, i.e., the average time after which all stars have experienced one collision, is given by \\begin{equation} t_{\\rm coll} \\simeq 2.1\\times 10^{12}\\,{\\rm yr}\\,\\frac{10^6\\,{\\rm pc}^{-3}}{n} \\frac{\\sigma_v}{30\\,{\\rm km\\,s}^{-1}} \\frac{{\\rm R}_\\odot}{R_\\ast} \\frac{{\\rm M}_\\odot}{M_\\ast}, \\label{eq:t_rlx} \\end{equation} where $n$ is the stellar density, $\\sigma_v$ the 1D velocity dispersion, and $R_\\ast$ and $M_\\ast$ are the typical values of the radius and the mass of a star, respectively. In view of this relation, one may think that direct collisions never play a role in cluster dynamics unless they are mediated by binary interactions, in which case the cross section is larger by a factor $a/R_\\ast$ where $a$ is the typical binary separation. The role of binaries in cluster dynamics is presented by Fregeau, McMillan and Portegies-Zwart in these proceedings. In contrast, here we focus on two cases in which the collisions between single stars are likely to play an important role: (1) fast core collapse followed by runaway collisions in young dense clusters and (2) stellar dynamics in the vicinity of a massive black hole (MBH) at the center of a galaxy. In case (1), if there are no primordial binaries, the merger sequence is likely to start before the first binary forms dynamically through 3-body interactions \\citep*{FGR04}; furthermore, primordial binaries probably foster collisions rather than preventing them by stalling the core collapse. In case (2), interaction rate peaks in the most central parts where densities exceed $10^7\\,{\\rm pc}^{-3}$. However, at this location owing to the MBH's presence, the velocity dispersion is also very high; as a result, most binaries are soft and should be dynamically dissociated.  To simulate the evolution of stellar clusters containing a large number of stars over time scales comparable to the relaxation (or collision) time, we use two Monte Carlo (MC) codes based on the pioneering ideas of \\citet{Henon71a,Henon71b}. These codes are described in detail in the literature \\citep{JRPZ00,JNR01,FB01a,FB02b}. The MC method assumes that the cluster is spherical, in dynamical equilibrium and that 2-body relaxation can be treated in the Fokker-Planck approximation. The cluster is represented as a set of spherical shells (``particles'') each of which may stand for a single star or a fixed number of stars which share the same orbital and stellar properties\\footnote{When one or a few particles start playing a special role, as in the collisional runaway process, using a a single particle to represent multiple stars is clearly questionable.}. Typical particle numbers used in the simulations range between $10^5$ and $10^7$. The stellar orbital motion is not followed explicitly; instead one tracks the long-term evolution of orbits (and stellar properties) subject to relaxation, direct stellar collisions, stellar evolution and possibly interactions with a central MBH. Even though the effect of binaries may be accounted for in the MC method \\citep{FregeauEtAl03}, here we consider only single stars. Various prescriptions can be used for the outcome of stellar collisions, ranging from sticky spheres to inter/extrapolation from a large database of SPH simulation results \\citep{FB04}. ",
        "conclusions": ""
    },
    "0410/astro-ph0410333_arXiv.txt": {
        "abstract": "We present an \\xmmn\\ observation of the Seyfert-LINER galaxy NGC 7213. The RGS soft X-ray spectrum is well fitted with a power law plus soft X-ray collisionally ionised thermal plasma ($kT = 0.18^{+0.03}_{-0.01}$ keV). We confirm the presence of Fe I, XXV and XXVI K$\\alpha$ emission in the EPIC spectrum and set tighter constraints on their equivalent widths of 82$^{+10}_{-13}$, 24$^{+9}_{-11}$ and 24$^{+10}_{-13}$ eV respectively. We compare the observed properties together with the inferred mass accretion rate of NGC~7213, to those of other Seyfert and LINER galaxies. We find that NGC~7213 has intermediate X-ray spectral properties lying between those of the weak AGN found in the LINER M\\,81 and higher luminosity Seyfert galaxies. There appears to be a continuous sequence of X-ray properties from the Galactic Centre through LINER galaxies to Seyferts, likely determined by the amount of material available for accretion in the central regions. ",
        "introduction": "Low ionisation nuclear emission line region (LINER) galaxies are characterised by optical emission line ratios which indicate a low level of ionisation \\citep{Heckman}. The origin of these emission lines is still the subject of debate: the lines are attributed either to shock heating (Baldwin, Phillips \\& Terlevich 1981) or to photoionisation by a central AGN (Ferland \\& Netzer 1983; Halpern \\& Steiner 1983). For example, Dopita \\& Sutherland (1996) have shown that shock heating can produce the optical line ratios observed in LINER galaxies, with shock velocities of 150 - 500 km s$^{-1}$. Such shocks can originate from stellar processes (e.g. supernovae), and can also be produced as a result of a jet or a wind from an AGN. On the other hand, when the optical line ratios of LINERs were compared to those measured in Seyfert galaxies, Ferland \\& Netzer (1983) found that they could explain the lines in both types of object using the same photoionising AGN continuum shape, but with a systematic variation of ionisation parameter, from low values in LINERs to high values in Seyferts.  NGC 7213 is a nearby ($z$ = 0.005977) S0 galaxy with AGN and LINER characteristics. It is clear that there is an AGN in this source, classified as a Seyfert 1 from its H$\\alpha$ line width (full width at zero intensity $\\sim$ 13000 km s$^{-1}$, Phillips 1979). A variety of optical emission lines are observed in this galaxy with velocities ranging from 200 to 2000 km s$^{-1}$ FWHM (Filippenko \\& Halpern 1984, hereafter FH84). FH84 argue that photoionisation of clouds spanning a range of densities and velocities by a non-stellar continuum is likely to be the mechanism creating the optical line emission. By invoking high density rather than high temperature clouds, FH84 eliminate the need for shock heating.  Since its discovery as a low luminosity X-ray source (Marshall \\etal 1978) NGC 7213 has been observed with several X-ray missions. The presence of a soft X-ray excess in NGC~7213 was implied by the results of an \\exosat\\ spectral survey of AGN \\citep{exosat} when the measured absorbing column for a single power law fit was found to be significantly lower than the Galactic value. In addition, the UV flux measured by Wu, Boggess \\& Gull (1983) was higher than would be expected from an extrapolation of the optical flux indicating that NGC~7213 may have a big blue bump (BBB), although weak compared to most Seyfert galaxies. The BBB is often interpreted as thermal emission from an accretion disc; if accretion discs are present in LINERs it is important to determine their properties if we are to understand the underlying emission mechanisms.  Turner \\& Pounds (1989) and Pounds \\etal (1994) reported possible Fe K$\\alpha$ emission in the X-ray spectrum observed with \\exosat\\ and \\ginga, but there was not sufficient resolution to unambiguously detect these features. Recently Bianchi \\etal (2003) combined data from the pn camera on \\xmmn\\ with simultaneous data from the PDS instrument on \\emph{BeppoSAX} to investigate the iron line complex and reflection hump. They found that a neutral Fe K$\\alpha$ line is present, with excess emission at higher energies best explained as weak narrow emission lines from highly ionised iron. The data do not show the presence of a significant reflection hump, suggesting that the neutral Fe K$\\alpha$ line originates in Compton-thin material. They also suggest that the highly ionised Fe emission may arise in material photoionised by the AGN power law continuum.  Here we present an analysis of the high resolution RGS spectra, 0.3 - 2 keV EPIC spectra and OM photometry taken during the same \\xmmn\\ pointing. We revisit the iron line complex with an analysis of the combined pn and MOS1 spectrum in which it is possible to obtain better constraints on individual line features. Identifying the physical mechanisms producing the X-ray emission may provide some clues to the origin of the optical emission lines where at present neither shock heating nor photoionisation by the AGN can be ruled out. We also compare the X-ray spectral properties of NGC~7213 with those of the nearest LINER galaxy M\\,81, which has also been studied in detail using \\xmmn\\ (Page \\etal 2003, 2004), and discuss the relationship between Seyfert galaxies and LINERs. ",
        "conclusions": "We present an \\xmmn\\ observation of the LINER galaxy NGC~7213 with the EPIC and RGS X-ray instruments and UV photometry with the OM. The X-ray spectrum is typical of a Seyfert galaxy, although with a much lower luminosity ($L_{2-10 keV}=1.7\\times 10^{42}$ erg s$^{-1}$). A collisionally ionised thermal plasma of temperature 0.18$^{+0.03}_{-0.01}$ keV is required in the RGS soft X-ray region where several low ionisation emission lines are detected, and we find no evidence for a significant blackbody-like soft-excess component. We suggest that this low ionisation thermal plasma may be a LINER characteristic. We constrain the equivalent widths of Fe I, XXV and XXVI K$\\alpha$ emission lines to 82$^{+10}_{-13}$, 24$^{+9}_{-11}$ and 24$^{+10}_{-13}$ eV respectively.  The nuclear radiation at X-ray wavelengths in NGC~7213 is substantially more Seyfert-like than the AGN component in the nearby LINER galaxy M\\,81, meaning NGC 7213 has the properties of a galaxy somewhere in between a typical Seyfert and a LINER galaxy. This supports the notion of continuity between the LINER and Seyfert classes, dictated by the luminosity of the central AGN and its accretion rate, both increasing from LINERs to Seyferts. In turn, the mass accretion rates depend critically on the amount of material present in the central regions surrounding the AGN. This continuous distribution can be extended to include the Galactic Centre, which shares some properties with LINER galaxies."
    },
    "0410/astro-ph0410105_arXiv.txt": {
        "abstract": "We report on the observations of five recent galactic classical novae, performed with XMM-Newton. Some of them show emission in the whole range from 0.2 to 8.0 keV, up to 3 years after the explosion. Nova Oph 1998 shows clear evidence of resumed accretion less than 3 years after its explosion, indicating a very fast recovery of the cataclysmic variable after the nova outburst. In another case, Nova Sgr 1998, it is difficult to disentangle if the emission is produced in the ejecta and/or in the accretion stream. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410275_arXiv.txt": {
        "abstract": "I review the unwrapping story the SN~1987A explosion event, and the  main discoveries associated with it. I will show that, although this supernova is somewhat peculiar, the study of SN~1987A has clarified quite a number of important aspects of the nature and the properties of supernovae in general. ",
        "introduction": "Supernova 1987A was discovered on February 24, 1987 by Ian Shelton (Kunkel \\& Madore 1987) in the Large Magellanic Cloud.  SN~1987A is the first supernova to reach naked eye visibility after the one studied by Kepler in 1604 AD and is undoubtedly the supernova event best studied ever by the astronomers. Actually, even if SN~1987A  has been more than hundred times fainter than its illustrious predecessors in the last millennium, it has been observed in such a detail and with such an accuracy to make this event a {\\it first} under many aspects (\\eg neutrino flux, progenitor identification, gamma ray flux) and definitely the {\\it best} studied of all.  The early evolution of SN~1987A has been highly unusual and completely at variance with the {\\it wisest} expectations. It brightened much faster than any other known supernova: in about one day it jumped from 12th up to 5th magnitude at optical wavelengths, corresponding to an increase of about a factor of thousand in luminosity.  However, equally soon its rise leveled off and took a much slower pace indicating that this supernova would have never reached those high peaks in luminosity as the astronomers were expecting.  Similarly, in the ultraviolet, the flux initially was very high, even higher than in the optical.  But since the very first observation, made with the International Ultraviolet Explorer (IUE in short) satellite less than fourteen hours after the discovery(Kirshner \\etal 1987, Wamsteker \\etal 1987), the ultraviolet flux declined very quickly, by almost a factor of ten per day for several days.  It looked as if it was going to be a quite disappointing event and, for sure, quite peculiar, thus not suited to provide any useful information  about {\\it ``normal\"} supernova explosions.  But, fortunately, this proved not to be the case and soon it became apparent that SN~1987A  has been the most valuable probe to test our ideas about the explosion of supernovae.  Reviews of both early and recent observations and their implications can be found in Arnett \\etal (1989), McCray (1993, 2003, 2004),  Gilmozzi and Panagia (1999), and Panagia (2003). In the following, I summarize some of the most  interesting findings from SN~1987A studies.   \\begin{figure}[!ht] \\plotfiddle{panagia_fig1_BW_small.eps}{6cm}{0}{30}{30}{-90}{-30} \\caption{True color picture ({\\it HST-WFPC2}) of SN 1987A, its companion stars, and the circumstellar rings (Courtesy of Peter Challis)).} \\end{figure}   \\section {Neutrino Emission from SN~1987A} For the first time ever, particle emission from a supernova was directly  measured from Earth: on February 23, around 7:36 Greenwich time, the neutrino telescope \"Kamiokande II\" (a big cylindrical ``tub\" of water, 16~m in diameter and 17~m in height, containing about 3300 m$^3$ of water, located in the Kamioka mine in Japan, about 1000~m underground) recorded the arrival of 9 neutrinos within an interval of 2 seconds and 3 more 9 to 13 seconds after the first one. Simultaneously, the same event was revealed by the IMB detector (located in the Morton-Thiokol salt mine near Faiport, Ohio) and by the  ``Baksan\" neutrino telescope (located in the North Caucasus Mountains, under Mount Andyrchi) which recorded 8 and 5 neutrinos, respectively,  within few seconds from each other.  This makes a total of 25 neutrinos from an explosion that allegedly produces 10 billions of billions of billions of billions of billions of billions of them! But a little more than two dozens neutrinos was enough to verify and confirm the theoretical predictions made for the core collapse of a massive star (\\eg Arnett \\etal~ 1989 and references therein). This process was believed to be the cause of the explosion of massive stars at the end of their lives, and SN 1987A has provided the experimental proof that the theoretical model was sound and correct, promoting it from a nice theory to the description of the truth.  \\section {SN~1987A Progenitor Star}  From both the presence of hydrogen in the ejected matter and the conspicuous flux of neutrinos, it was clear that the star which had exploded was quite massive, about twenty times more than our Sun. And all of the disappointing peculiarities were due to the fact that just before the explosion the supernova progenitor was a blue supergiant star instead of being a red supergiant as common wisdom was predicting. There is no doubt about this explanation because SN~1987A is exactly at the same position as that of a well known blue supergiant, Sk~$-69^{\\circ}$~202. And the IUE observations indicated that such a star was not shining any more after the explosion: the blue supergiant star  unambiguously was the SN progenitor. This heretic possibility was first suggested in Panagia \\etal (1987) and subsequently confirmed by the more detailed analyses presented by Gilmozzi \\etal (1987) and Sonneborn, Altner \\& Kirshner (1987).  On the other hand,  the presence of narrow emission lines of highly ionized species, detected in SN~1987A short wavelength spectrum since late May 1987,  has provided evidence for the progenitor having been a red supergiant before coming back toward the blue side of the HR diagram (Fransson \\etal 1989). Also, the detection of early radio emission that decayed in a few weeks (Turtle \\etal 1987) indicated that the ejecta were expanding within a circumstellar environment whose properties were a perfect match to the expected wind of a blue supergiant progenitor (Chevalier \\& Dwarkadas 1995).  Such an evolution for an $\\sim20~M_\\odot$ star was not expected, and theorists struggled quite a bit to find a plausible explanation for it.  As summarized by Podsiadloswki (1992), in order to explain all characteristics of SN~1987A, rotation  has to play a crucial role, thus limiting the possibilities to models involving either a rapidly rotating single star (Langer 1991), or a stellar merger in a massive binary system (Podsiadloswki 1992).  \\begin{figure}[!ht] \\plotfiddle{panagia_fig2_small.eps}{5cm}{-90}{30}{30}{-135}{170} \\caption{Evolution of the UV and optical spectrum of SN~1987A (Pun et al. 1995).} \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410043_arXiv.txt": {
        "abstract": "We compute the rates $P$ at which  energy is injected into the p modes for a set of 3D simulations of  outer layers of stars. We found that $P_{\\rm max}$ -~ the maximum in $P$ ~ - scales as $(L/M)^{s}$ where $s$ is the slope of the power law,  L and M are the luminosity and the mass of the 1D stellar models associated with the simulations. The slope is found to depend significantly on the adopted representation for the turbulent eddy-time correlation function, $\\chi_k$. According to the expected performances of COROT, it will likely be possible to measure $P_{\\rm max}$  as a function of $L/M$ and to constrain the properties of stellar turbulence as the turbulent eddy time-correlation. ",
        "introduction": "Stars with masses $  M \\lesssim 2 M_\\odot$ have upper convective zones where the stochastic excitation of p~modes takes place. Models of stochastic excitation yield the rate $P$ at which p modes are excited by turbulent convection but require an accurate knowledge of the time averaged and -~ above all ~ - the dynamic properties of turbulent convection. The latter are represented in the approach of Samadi \\& Goupil (2001) by $\\chi_k$, the frequency component of the kinetic energy spectrum, also designed as the convective eddy time-correlations.  A Gaussian   is usually assumed for $\\chi_k$ , {\\it i.e.}  in other word it is usually assumed that two distant points in the turbulent medium are uncorrelated.  Here, we investigate two different forms for $\\chi_k$:  the Gaussian and the Lorentzian.  We consider a set of 3D simulations of solar-like oscillating stars. We compute internal structure of 1D stellar models and associated eigenfunctions consistent with the 3D simulations. The   3D simulations provide constraints on  the turbulent kinetic spectrum $E(k)$. Theses constraints enable us to compute $P_{\\rm max}$, the maximum in $P$, for the 1D internal models consistent with the 3D simulations (see details in Samadi et al, 2004). ",
        "conclusions": "\\begin{figure}[h] \\centering \\includegraphics[width=9cm]{samadi_fig1.ps} \\vspace{-0.5cm} \\caption{Computed ${P}_{\\rm max}$ versus $L/M$: The crosses assume a Gaussian for  $\\chi_k$ and the squares a Lorentzian respectively. The lines are the results of fitting each set of dots with a power law of the form $(L/M)^{s}$ where $s$ is the slope of the power law. } \\label{figure_mafig} \\end{figure}   Results of the calculations of ${P}_{\\rm max}$  are shown in Fig.~(1). We find that $P_{\\rm max}$  scales as $(L/M)^s$ where $s$  is the slope of the scaling law, $L$ is the   luminosity and $M$ is the mass of the 1D~models associated with the 3D simulations. The value found for $s$ is  sensitive to the choice for $\\chi_k$: Indeed the Gaussian results in  $s=3.3$  and the Lorentzian in $s=2.6$. Solar seismic observations can clearly distinguish between the two forms for $\\chi_k$: they  favor the Lorentzian (see Samadi et al, 2003). For other stars for which stochastically excited p modes are expected, the Lorentzian is likely to remain the correct form for  $\\chi_k$, but this needs to be confirmed. The high-quality data of the COROT instrument will enable us to derive $P_{\\rm max}$ as a function of L and M for the different target stars. These data will very likely enable us to  determine the appropriate model for the eddy time-correlations  ($\\chi_k$)."
    },
    "0410/hep-ph0410095_arXiv.txt": {
        "abstract": "{The renormalization group (RG) approach to cosmology is an efficient method to study the possible evolution of the cosmological parameters from the point of view of quantum field theory (QFT) in curved space-time. In this work we continue our previous investigations of the RG method based on potential low-energy effects induced from physics at very high energy scales $M_X\\lesssim M_P$. In the present instance we assume that both the Newton constant, $G$, and the cosmological term, $\\CC$, can be functions of a scale parameter $\\mu$. It turns out that $G(\\mu)$ evolves according to a logarithmic law which may lead to asymptotic freedom of gravity, similar to the gauge coupling in QCD. At the same time $\\Lambda(\\mu)$ evolves quadratically with $\\mu$. We study the consistency and cosmological consequences of these laws when $\\mu\\simeq H$. Furthermore, we propose to extend this method to the astrophysical domain after identifying the local RG scale at the galactic level. It turns out that Kepler's third law of celestial mechanics receives quantum corrections that may help to explain the flat rotation curves of the galaxies without introducing the dark matter hypothesis. The origin of these effects (cosmological and astrophysical) could be linked, in our framework, to physics at $M_X\\sim 10^{16-17}\\,GeV$.} ",
        "introduction": "\\label{sect:Intro}  The recent developments in physical cosmology have provided an astonishingly accurate picture of our universe, within the canons of the Friedmann-Lema\\^\\i tre-Robertson-Walker (FLRW) paradigm\\,\\cite{Peebles}. From these achievements emerged what has been called the {\\em cosmological concordance model}, characterized by an essentially zero value of the spatial curvature parameter and a non-vanishing positive value of the cosmological constant (CC) term, $\\CC$, in Einstein's equations -- or in general of some form of dark energy (DE) that cannot be attributed to any form of matter and radiation. This new standard cosmological model supersedes and refutes (by many standard deviations) the old Einstein-de Sitter (or critical-density) cosmological model, which has been the standard cosmological model until recently. The latter is also a spatially flat model, but it is characterized by a zero value of the cosmological constant. Evidence for the values of the cosmological parameters within the new concordance model comes both from tracing the rate of expansion of the universe with high-z Type Ia supernovae (SNe Ia) experiments \\cite{R98P99,R04P03} and from the precise measurement of the anisotropies in the cosmic microwave background (CMB) radiation\\,\\cite{WMAP2,WMAP1}. These experiments indicate that the value of the cosmological constant is around\\footnote{The original CC term, $\\lambda$, in Einstein's equations is related to our $\\CC$ by $\\lambda=8\\pi\\,G\\,\\CC$, where $G$ is Newton's constant. Notice that $\\CC$, in our notation, has dimensions of energy density, and to emphasize this fact it will sometimes be denoted $\\CC\\equiv\\rL$.} \\begin{equation}\\label{CCvalue} \\CC_0=\\OL^0\\,\\rc^0\\simeq 6\\,h_0^2\\times 10^{-47}\\,GeV^{4}\\,, \\end{equation} where $\\OL^0\\sim 70\\%$. {Here $\\rc^0 \\simeq \\left( 3.0\\,\\sqrt{h_{0}}\\times 10^{-12}\\,GeV\\right) ^{4}$ is the present value of the critical density and $H_0 = 100\\,h_0$ km/s/Mpc, with $h_{0}= 0.72\\pm 0.08$, is the most precisely known value of Hubble's constant\\,\\cite{HST03}. In the context of the Standard Model (SM) of electroweak interactions, the physical (measured) value (\\ref{CCvalue}) should be the sum of the vacuum CC term in Einstein's equations, $\\CC _{vac}$, and the induced contribution from the vacuum expectation value of the Higgs effective potential, $\\CC_{ind}=\\langle V_{\\rm eff} \\rangle$, namely $\\CC=\\CC _{vac}+\\CC _{ind}\\,.$ Other contributions should be added within the SM, like those from QCD, but they are much smaller. It is only the combined parameter $\\CC$ that makes physical sense, whereas both $\\CC _{vac}$ and $\\CC _{ind}$ remain individually unobservable-- see e.g. \\cite{JHEPCC1} for a more detailed explanation. From the current LEP numerical bound on the Higgs boson mass, $\\,M_{H}> 114.4\\,GeV$ (at the $95\\%$ C.L.) \\cite{LEPHWG}, one finds $\\,\\,\\,\\left| \\CC _{ind}\\right|=(1/8)M_H^2\\,v^2> 1.0\\times 10^{8}\\,GeV^{4}$, where $v\\simeq 246\\,GeV$ is the vacuum expectation value of the Higgs field. Clearly, $\\left| \\CC _{ind}\\right|$ is $55$ orders of magnitude larger than the observed CC value (\\ref{CCvalue}). Such discrepancy, the so-called ``old'' CC problem \\cite{weinRMP,CCRev}, manifests itself in the necessity of enforcing an unnaturally exact fine tuning of the original cosmological term $\\CC _{vac}$ in the vacuum action that has to cancel the induced counterpart $\\,\\CC _{ind}$ within a precision (in the SM) of one part in $\\,10^{55}$. But the CC problem seems to be quite more general and it seems to plague all physical theories (QFT and string theory) making use of the concept of vacuum energy\\, \\cite{weinRMP,CCRev,wittenDM}. Moreover, all attempts to deduce the small value of the cosmological constant from a sound theoretical idea involve some explicit or implicit form of severe fine-tuning among the parameters of the model. This became clear already from the first attempts to treat the dark energy component as a dynamical scalar field \\cite{Dolgov,PSW,Aguila}. More recently this approach took the popular form of a ``quintessence'' field slow--rolling down its potential\\,\\cite{Caldwell} and has adopted many different forms \\cite{PeebRat,Moreq}. The main motivation for the quintessence models is that they could provide a variable dark energy. Although there is no compelling reason for a time evolving DE at the moment -- and it is difficult to get an experimental handle on it \\,\\cite{CCvariableTest}--  it is even more difficult to accept (at least from the theoretical point of view) that the DE can be described by a static CC term in Einstein's equations with the small value (\\ref{CCvalue}) throughout most of the history of the universe, namely after the electroweak phase transition took place.   An alternative  possibility is to think of the DE as a time-evolving CC. This has been put forward from a purely phenomenological point of view in many places of the literature\\,\\cite{CCphenom,Overduin}. However, a treatment of an evolving CC term from the more fundamental point of view of the renormalization group in QFT, with a direct link to Particle Physics, has been suggested only recently\\,\\cite{cosm,JHEPCC1,Babic1,RGTypeIa1}. The possible experimental consequences of this idea in the light of high-z SNe Ia data have also been tested in great detail in \\,\\cite{RGTypeIa2} (see also the framework and analysis of \\cite{Reuter03a,Bentivegna}). In the renormalization group method one takes the point of view that $\\CC$ should not be constant, because the quantum effects may shift away the prescribed value, even if the latter is assumed to be zero\\,\\cite{cosm}. In this approach the equation of state for the DE is the one for the ``true'' cosmological term $p_\\Lambda = - \\rho_\\Lambda$. Hence it is fundamentally different from all kinds of quintessence-like proposals, where the CC is assumed to vanish precisely and the role of the DE is played by some new entity with a nontrivial equation of state. In the RG approach the CC becomes time-evolving, or redshift dependent: $\\CC=\\CC (z)$. Although we do not have a quantum theory of gravity where the running of the gravitational and cosmological constants could ultimately be substantiated, a semiclassical description within the well established formalism of QFT in curved space-time (see e.g. \\cite{birdav,book}) should be a good starting point. From the RG point of view, $\\CC$ becomes a scaling parameter $\\CC (\\mu)$ whose value should be sensitive to the entire energy history of the universe -- in a manner not essentially different to, say, the electromagnetic coupling constant in QED or the gauge coupling constant in QCD.  While the RG method to tackle the CC problem has been used earlier \\cite{Nelson,Polyakov,lam},  the decoupling effects of the massive particles at low energies may change significantly the structure of the renormalization group equations (RGE) with important phenomenological consequences. This feature has been pointed out recently by several authors from various interesting points of view \\cite{Babic1,JHEPCC1}. However, it is not easy to achieve a RG cosmological model where $\\CC$ runs smoothly without fine tuning at the present epoch. In Ref.\\,\\,\\cite{RGTypeIa1,IRGA03} a successful attempt in this direction has been made, which is based on possible quantum effects near a high mass scale $M_X$, not far away from the Planck scale $M_P$, that are transferred (via soft decoupling) to the low energy physics. At the same time, the approximate coincidence of the observed CC and the matter density, $\\ \\OL^0\\sim \\OM^0\\,$ (i.e. the ``new'' CC problem, or ``time coincidence problem'' \\cite{weinRMP,CCRev}) can be alleviated in this framework\\,\\cite{IRGA03}. Thus the RG cosmologies constitute a healthy alternative to the \\textit{ad hoc} quintessence models and are worth exploring\\footnote{See \\cite{AHEP03} for a summarized presentation of this RG framework and an exposition of the basic results. Recently some applications of the RG method to inhomogeneous cosmologies have also been proposed\\,\\cite{Ibanez}.}. In the present paper we generalize the RG approach developed in \\cite{JHEPCC1,Babic1,RGTypeIa1,IRGA03} by including the possibility to have a running Newton's constant as well\\,\\footnote{Phenomenological models with variable $G$ and $\\CC$, without reference to the RG, have been studied in the literature since long ago\\,\\cite{CCphenom, Overduin}. See also the more recent papers\\,\\cite{Stefancic}. }. A fundamental question that we address here is the possibility to maintain the phenomenologically successful quadratic evolution law $\\delta\\CC\\sim H^2$ (emerging from an RGE of the form $d\\CC/d\\ln H\\propto H^2$) as proposed in\\,\\cite{JHEPCC1,RGTypeIa1,IRGA03}, also in a $G$-running framework. This law is based on the assumption that the physical RG scale is defined by $\\mu=H$, the Hubble parameter. We will see that this Ansatz is consistent with the systematic RG scale-setting procedure proposed recently in\\,\\cite{Babic2}. In addition, the $\\delta\\CC\\sim H^2$ law is the only one among those that are consistent with this procedure which could have measurable phenomenological consequences. Therefore, it is of foremost importance to explore its phenomenological viability, in spite of the insurmountable technical difficulties encountered at present for an explicit calculation in QFT on a curved background\\,\\cite{apco}. In this sense it is remarkable that the quadratic law $\\delta\\CC\\sim H^2$ can also be motivated from very general arguments in QFT\\,\\cite{PadmanabhanH2} which may go beyond the RG. No less appealing is that this quadratic law can further be claimed, at least heuristically, on the grounds of the so-called Holographic Principle (HP) applied to effective field theories\\,\\cite{Horvat}. In fact, by linking IR and UV scales through entropy bounds, the HP may provide essential ingredients to relate the microscopic world of Particle Physics and the large scale structure of the Universe\\,\\footnote{See e.g. \\cite{Bousso} for a general review of the HP and its applications.}.  A main result of this paper is that the adoption of that quadratic law will lead us to an alternative cosmological RG framework to that presented in Ref.\\,\\cite{RGTypeIa1,IRGA03}, keeping though all the nice phenomenological features of the latter, however with the distinctive property of having a (logarithmically) running gravitational constant. This last property endows the new cosmological model with potentially far-reaching phenomenological implications both at the cosmological and astrophysical level which will be explored in some detail in this paper.  The structure of the paper is as follows. In the next section we introduce a few general features of the cosmological models based on variable $G$ and $\\CC$. In Section \\ref{sect:RGmodel1} we provide some general discussion of the RG cosmological models. The consistency of the RG  scale setting in cosmology is discussed in Section \\ref{sect:RGmodel3}. In Section \\ref{sect:RGmodel2} we focus on the cosmological model presented here, characterized by a logarithmic running of $G$ and a quadratic running of $\\CC$. The numerical analysis of this model is presented in Section \\ref{sect:numanalysis}, where we also show the predicted deviations from the standard FLRW expectations. In Section \\ref{sect:Kepler} we narrow down the applicability of this RG cosmological model to the astrophysical domain by considering the possibility to attribute the origin of the flat rotation curves of the galaxies to a RG quantum correction to Kepler's third law of celestial mechanics. We close with some final discussion in the last section, where we draw our conclusions. ",
        "conclusions": "\\label{sect:conclusions}  In this paper we have further elaborated on the formulation of the renormalization group cosmologies along the lines of \\,\\cite{JHEPCC1, cosm, Babic1,RGTypeIa1,RGTypeIa2,IRGA03,AHEP03,Babic2}. The main aim of these RG cosmologies is to provide a fundamental answer, within QFT, to the meaning of the most basic cosmological parameters such as $H_0$, $G$, $\\rM^0$ and $\\CC_0\\equiv\\rL^0$, and at the same time to try to understand their presently measured values and possible interrelations among them. In doing so we have found that the classical FLRW cosmologies become ``renormalization group improved'', and therefore they become slightly modified by quantum corrections associated to renormalization group equations that govern the evolution of the parameters $\\CC$ and $G$. This evolution behaviour, primarily a scaling one within the framework of the RG,  can be inter-converted into a temporal evolution (i.e. in terms of the cosmic time or, if desired, in terms of the redshift), after an appropriate identification of the scale parameter $\\mu$ of the RG. We have found that $\\mu\\sim H(t)$ is a consistent identification for the class of RG cosmologies in which $\\CC$ evolves quadratically, and $G$ evolves logarithmically, with $\\mu$. And we have shown that this combination of running laws is a most promising one from the point of view of phenomenology. The scaling (and redshift) variation of these two laws is very simple and it is controlled by a single parameter, $\\nu$ -- cf. Eq.\\,(\\ref{nu}). For $\\nu>0$ the running law for the gravitational constant exhibits (logarithmic) asymptotic freedom, much as the strong coupling constant in QCD. For $\\nu\\ll 1$ the quadratic evolution law for $\\CC$ turns out to be very similar to the one previously found in \\cite{RGTypeIa1,RGTypeIa2} and therefore it can be tested in the future experiments by SNAP\\,\\cite{SNAP}, provided $\\nu$ is not much smaller than $\\nu_0\\sim 10^{-2}$. This value is presently consistent with data from high-z SNe Ia, CMB and LSS. Furthermore, from the point of view of the cosmological constant problem, we have seen that this RG framework can shed some light to the coincidence cosmological problem ($\\rL^0\\simeq\\rM^0$ ?) because it provides a natural explanation of the puzzle of why the energy scale associated to the cosmological constant is at present the geometric mean of the Hubble parameter and the Planck mass: $E_{\\CC}(t_0)\\simeq\\sqrt{M_P\\,H_0}$.   It is important to remark that the quadratic law (\\ref{CCH}) associated to our RG framework is not just a proportionality law $\\CC\\sim H^2$, but a variation law: $\\delta \\CC\\sim H^2$. In contradistinction to the former, the latter actually introduces an additive term. This is because it originally comes from a RGE of the form $d\\CC/d\\ln H\\propto H^2$, which has to be integrated, see Eq.\\,(\\ref{RG01}). Therefore the ensuing cosmologies are very different from the $\\CC\\sim H^2$ ones. Only the latter have been amply discussed in the past in many places of the literature\\,\\cite{CCphenom,Overduin}, including some recently raised criticisms\\,\\cite{Wang,Hsu}. In contrast the $\\delta \\CC\\sim H^2$ law is, to the best of our knowledge, relatively new in the literature\\,\\cite{JHEPCC1}, and its motivation within QFT gives it a credit at a much more fundamental level. Let us mention that some variation of this law has been proposed very recently (on pure phenomenological grounds) in \\cite{Wang}, in the form $d\\CC/da\\propto dR/da$, where $R$ is the curvature scalar.  However, since $\\mu=R^{1/2}$ is, as mentioned above (and in \\cite{JHEPCC1}), a covariant generalization of $\\mu=H$, it is essentially the same law and so the conclusions are similar. Therefore the RG cosmologies associated to $\\delta \\CC\\sim H^2$ can also be formulated covariantly, without committing to the particular FLRW metric. But from the practical point of view (and also to better assess its physical meaning and potential applications\\,\\cite{RGTypeIa2,AHEP03}) it is better to formulate this law  making use of the Hubble parameter, instead of curvature.  The last part of our paper is more speculative. We have tantalized the possibility of extending the RG formalism to local cosmological domains such as the galactic level. After arguing that the local RG scale could naturally be $\\mu\\sim 1/r$, we have solved the modified Schwarzschild metric in the limit $\\nu\\ll 1$ and derived a RG correction to Kepler's third law of celestial mechanics. We have hypothesized that this quantum correction could furnish an explanation for the flat rotation curves of the galaxies, provided $\\nu\\sim 10^{-6}$. Obviously, this possibility would exclude measurable effects on the cosmology side, where a larger $\\nu$ is needed, but it might hit square on one of the biggest astrophysical conundrums of the last thirty years. This result is curious enough and rather intriguing, and we have found interesting to include it as an additional potential application of our RG framework.  All in all, we deem this last part of our work more speculative than the rest, because we have had to descend from the cosmological level to the (more complicated) local domain where cosmology is entangled with astrophysics, and other factors may come into play.  Thus, more detailed theoretical considerations are needed to formally establish the RG effects at the galactic level,  which could be modified by other factors such as the modeling of the galaxy and the details of the RG framework. For this reason, even if $\\nu$ is of order of $\\nu_0$, and therefore much larger than (\\ref{nuvalue}), the corresponding RG predictions at the galactic level are not necessary in contradiction with the experimental situation. The actual effects could be much smaller due to, say, the need to introduce some significant modification in the Ansatz (\\ref{newscale}), as noted in the last section. Finally, even if accepting that $\\nu$ is small enough for the RG to be able to provide some explanation of the flat rotation curves, it does not preclude the possibility that galactic dark matter is present in the haloes of the galaxies. It only shows that other, even more subtle, effects could creep in. Whether or not these effects are directly related to the RG must be further investigated, but already at the pure phenomenological level the possibility that a logarithmic law of the type (\\ref{GH2}) could be behind the flat rotation curves looks rather enticing. Obviously, even if it is not the sole explanation, it may help to reduce the total amount of dark matter needed to account for the flat rotation curves and also to explain some features related to the deviation from universality of the flat velocities. In any case it is clear that further experimental work will also be necessary to decide on the ultimate origin of the rotation curves of the galaxies\\,\\cite{Prada}. In general, even if letting aside the possibility of these local effects, our cosmological RG framework provides a sound alternative to explain the nature of the cosmological parameters and their potential connection to quantum field theory. It also offers a link between the low-energy physics characteristic of large scale Cosmology with the high-energy interactions of Particle Physics. Last, but not least, it constitutes a ready testable framework for the next generation of cosmological experiments. \\\\   {\\bf Note Added:}  Shortly after the e-print Archive version of this paper was submitted as  \\texttt{hep-ph/0410095}, some other works were submitted to the Archive -- see \\cite{ReuterWeyer,Moffat}-- which overlap significantly with the content of Section 7. These works also use the RG to study the flat rotation curves of the galaxies and arrive to conclusions very similar to the ones presented here. Since, however, their approach is different from ours we think it can be useful for the reader if we comment on the similarities and also the main differences between the two methods.  First of all we recall that two Renormalization Group formulations for gravity are present in the literature, namely the perturbative RG for curved space-time and the Wilsonian non-perturbative RG for quantum gravity. The present work is based on the first of these RG approaches (see \\cite{birdav,book} for an introduction and references). The object of our interest is the quantum theory of matter fields on classical curved background. It is very important that this theory is renormalizable and therefore application of the perturbative RG to this case is a consistent procedure. Unfortunately, the consistency does not mean that this approach does not meet real difficulties. The existing well-established formalism is essentially restricted by the $\\overline{\\mbox{MS}}$-scheme of renormalization. Within this approach one can learn how the observables dependent on the dimensional parameter $\\,\\mu$, but the physical interpretation of this parameter always requires a special additional effort. Furthermore, the $\\overline{\\mbox{MS}}$-scheme is reliable only in the UV limit and is not appropriate for the investigation of the relatively weak IR scaling, that is for the study of decoupling. An alternative mass-dependent scheme has been recently applied in curved space-time, but only in the particular case of linearized gravity \\cite{apco}. In this case we can observe the Appelquist\\&Carazzone decoupling theorem\\,\\cite{AC} to hold in the higher derivative sectors of the gravitational action. The present and previous \\cite{JHEPCC1,Babic1,RGTypeIa1,RGTypeIa2} works are based on the phenomenological assumption that a similar kind of decoupling behaviour takes place for the cosmological constant. One can find an extended discussion of this issue in the Appendix A of \\cite{RGTypeIa2}.  On the other hand, the alternative works \\cite{ReuterWeyer} (exactly as the previous ones \\cite{Reuter03a,Bentivegna}, see further references therein) deal with the IR effects of quantum gravity. It is well known that quantum gravity is inconsistent within the perturbative approach and hence the appropriate version of the RG is necessarily non-perturbative. Ref.\\,\\cite{ReuterWeyer} is based on a simplified version of the Wilsonian approach to this non-perturbative RG framework. The objective is the derivation of the functional of the flow effective action, depending on the cut-off parameter $\\,k$. This parameter can be viewed as a direct analog of the parameter $\\,\\mu\\,$ for the perturbative RG in the $\\overline{\\mbox{MS}}$-scheme. The main similarity between the two parameters is that one also needs to properly identify $\\,k\\,$ with certain physical quantity (e.g. the energy of the gravitational quanta or the inverse of the cosmic time $\\,t$). Without such identification the theory has no physical output. The choice of the scaling parameter in both RG approaches is not constrained by the RG formalism itself, and in certain (sufficiently simple) situations it can be guessed on the basis of physical (e.g. dimensional) arguments. This is the case when the system has just one relevant physical scale. For instance, in the astrophysical case under consideration in Section \\ref{sect:Kepler} the galaxy is assumed to have spherical symmetry and then the single relevant physical scale is the radial distance $r$ from its center. This leads to the successful scale choice (\\ref{newscale}). On the other hand, in the cosmological context, a possible way of constraining the scale setting is the procedure applied in Section \\ref{sect:RGmodel3}. This leads immediately to our choice $\\mu=H$ -- cf. Eq.\\,(\\ref{muH}) -- roughly consistent with the scale setting $k\\sim 1/t$ advocated in the alternative RG approach of \\cite{ReuterWeyer}. However, even this regular procedure does not obviously provide universal solution for all physical situations (for example, in multiple scale systems or in the transition from the cosmological to the astrophysical domain). Therefore, the link between the RG running and the $\\,r$-dependence in the astrophysical setting is an essentially phenomenological assumption for both RG formulations, and needs further investigation.  As we already mentioned above, in our RG approach we assume the standard form of decoupling for the cosmological constant. Similarly, the alternative RG approach -- the Wilsonian non-perturbative formulation of the RG for quantum gravity -- requires also of strong phenomenological input. In the Wilsonian version of the RG used in \\cite{ReuterWeyer} the effective action depends on infinitely many couplings associated to powers of curvature and non-local invariants. Solving the exact functional RG flow equation is virtually impossible at present and in practice one truncates this effective action into a linear combination of powers of curvature, the simplest possibility being the powers one and zero, i.e. the Einstein-Hilbert term and the cosmological term. The coefficients of these terms define the flow trajectory $(G(k),\\Lambda(k))$ projected onto this simplified finite-dimensional space.  However, by virtue of this simplification the flow trajectory $G=G(k)\\,,\\Lambda=\\Lambda(k)$ cannot be derived from the RG flow equation, and therefore an Ansatz on the trajectory is to be made in general. For the galactic problem at hand, one typically dismisses the role played by the cosmological term and is left with proposing an (\\textit{ad hoc}) modeling for the $G$ flow. The model trajectory explored in \\,\\cite{ReuterWeyer} reads $G(k)=1/k^q$ (with $q$ an arbitrary parameter). In order to describe the flat rotation curves the authors of \\cite{ReuterWeyer} have to take for $q$ a very small value of order $10^{-6}$. In contrast, in a previous paper\\,\\cite{Bentivegna} it was assumed the particular value $q=2$ in order to describe a RG cosmology with a (hypothetical) IR fixed point. Obviously this trial and error strategy shows that the state of the art in the RG flow effective action approach is not good enough at present as to allow for determining the flow trajectory $G=G(k)$ in a fully dynamical way. It is nevertheless remarkable that after accepting the value $10^{-6}$ the previous power law trajectory becomes effectively a logarithmic running law, which is entirely similar to the one we have obtained in Eq.\\,(\\ref{GH2}) on a very different basis.  Comparing the merits of the two RG approaches we can indicate two advantages (from our point of view) of the one considered in the present paper. {\\it First}, we are dealing with the quantum effects of matter fields, which definitely take place. The necessary phenomenological implication concerns only the technical (albeit serious) difficulty of evaluating the vacuum effective action of massive fields on a curved background. In contrast, the relevance of the quantum gravity effects and the possibility of their description without solving the fundamental problem of quantum gravity looks rather uncertain. This concerns both the method of Ref. \\cite{ReuterWeyer} and the one of the old works \\cite{Oldpapers} where the perturbative higher derivative quantum gravity effects were considered (and similar results were obtained). {\\it Second}, the present approach has at least the virtue of introducing a parameter, $\\nu$, whose definition (namely, $\\nu$ is the ratio of the masses squared of the heavy matter fields versus the Planck mass squared, see Eq.\\,(\\ref{nu})), naturally fits in the right order of magnitude to explain the effect. At then end of the day, since both approaches are essentially based on phenomenological input, the preference of one of them remains a matter of taste. However, the net outcome is that the running law (\\ref{GH2}) and the one found in \\cite{ReuterWeyer} are identical if one identifies $q\\equiv 2\\,\\nu$ in the limit of very small value of these parameters (the only relevant limit for the galactic dynamics).  Let us also mention a recent modification of the RG approach, the so-called metric-skew-tensor-gravity theory proposed in \\cite{Moffat}. Here a skew-symmetric field tensor $F_{\\mu\\nu\\lambda\\sigma}$ is introduced and coupled to matter. This model introduces a new particle, but the main contribution to the mass of the galaxy is not from the skew field energy density, but from the RG flow running of $G$ (induced by the coupling of this field to ordinary matter), which leads to a modified Newtonian acceleration law that can be successfully fitted to the galaxy rotation curves. Again this is achieved for special RG trajectories for which gravity becomes a (slowly varying) ``confining force''. In this case the trajectory is again the $q$-dependent one mentioned above, but for the original value $q=2$. The price of returning to this value, which is associated to a (hypothetical) IR fixed point, is to postulate the existence of a new particle in the galaxy halo. The physical consequence is again similar to our running law  for $G$, Eq.\\, (\\ref{GH2}), which entails anti-screening for $\\nu>0$, much in the same way as the QCD force at small distances.  We find suggestive that all these RG approaches, even though they differ in the particular details, seem to converge to the same basic idea that the flat rotation curves of the galaxies could perhaps be understood without using the hypothesis of CDM. At the moment these RG frameworks cannot account for subleading effects, such as the Tully-Fisher law. However, they may constitute a promising framework to eventually construct the theory of the flat rotation curves of the galaxies from the first principles of Quantum Field Theory. The necessary modifications to account for subleading effects can be a matter of improving the technicalities of the RG framework and of the scale-setting procedure, including a more accurate modeling of the galaxy.  \\vspace{0.3cm}   {\\bf Acknowledgments:}  The work of I.Sh. has been supported in part by CNPq, FAPEMIG and ICTP. J.S. has been supported in part by MECYT and FEDER under project FPA2001-3598. J.S. is also thankful to the Max Planck Institut f\\\"ur Physik in Munich for the warmth hospitality and the financial support provided. The work of H. \\v{S}tefan\\v{c}i\\'{c} is supported by the Ministry of Science, Education and Sport of the Republic of Croatia under the contract  No. 0098002."
    },
    "0410/astro-ph0410725_arXiv.txt": {
        "abstract": "Deep $J$, $H$ and $K_s$ photometry of the low Galactic latitude Norma cluster has been obtained with the 1.4-m InfraRed Survey Facility (IRSF) and with SOFI on the NTT. Together with spectroscopy taken at the Anglo-Australian Observatory using 2dF, the near-infrared Fundamental Plane of the Norma cluster will be used to determine the distance to this cluster and to assess its location within the Great Attractor overdensity. Potential systematic uncertainties related to the determination of {peculiar velocities in the Zone of Avoidance} are discussed in some detail. In particular the effects (uncertainties, systematics and scatter) of extinction and star-crowding on the derived magnitudes at low latitudes are investigated in great detail. We present a first look at the emerging $K_s$-band Fundamental Plane of the Norma cluster. We expect an uncertainty of $\\sim$3\\% ($\\sim 150$ km s$^{-1}$ at the distance of the Norma cluster) in our final Fundemental Plane distance estimate based on 76 galaxies. ",
        "introduction": "The {Norma cluster} ({ACO 3627}, Abell et al.~1989) is a rich and nearby galaxy cluster at low Galactic latitude ($\\ell, b, cz$) = ($325.3^{\\circ}, -7.2^{\\circ}, 4844$ km s$^{-1}$). Kraan-Korteweg et al. (1996) and Woudt (1998) showed it to be the nearest massive rich cluster in the local Universe and comparable in its properties to the well-studied {Coma cluster}. Were it not for the  obscuration by the Milky Way (A$_{\\rm B}$ $\\sim$ 1\\fm0) the  prominence of Norma would have been apparent a century ago and it might have taken the first place as the best-studied cluster in  the local Universe. Following the recognition as a nearby rich cluster, however, the Norma cluster has been subjected to dedicated multi-wavelength observations: detailed studies of its X-ray properties have now been made by B\\\"ohringer et al. (1996) and Tamura et al. (1998), the effects of the  intercluster medium on the gas contents of spiral galaxies in the  Norma cluster by Vollmer et al. (2001), substructure and dynamics by Woudt (1998), and a derivation of its optical and near infrared (NIR) luminosity function is in preparation (Woudt et al., in prep.).  A further major point of interest lies in its location at (or close to) the bottom of the potential well of the {Great Attractor}. Many years after the initial discovery of the GA (Dressler et al.~1987; Lynden-Bell et al.~1988) there still remains some ambiguity regarding the extent and nature of the GA -- as was evident by various lively discussions at this meeting upon the simple question `What is the Great Attractor?'.  One of our aims in studying the Norma cluster is a highly acccurate distance determination through a NIR {Fundamental Plane} (FP) analysis with which we want to pin down the position of the the Norma cluster within the GA overdensity, determine its -- and possibly the whole GA's -- movement with respect to the Cosmic Microwave Background (CMB), and see whether the Shapley Concentration (Scaramella et al.~1989; Allen et al.~1990) at ($\\ell, b, cz$) $\\sim$ ($306^{\\circ}, +30^{\\circ}, 15\\,000$ km s$^{-1}$) exerts a non-negligible attraction on the GA (see, e.g., Kocevski et al.~and Lucey, Radburn-Smith \\& Hudson, these proceedings). ",
        "conclusions": "Following careful considerations of issues regarding the Galactic foreground extinction, star-crowding, and substructure within the Norma cluster, an accurate distance determination (accurate to $\\sim 150$ km s$^{-1}$) to the Norma cluster will follow from our near-infrared FP analysis of 76 early-type galaxies.  It appears that for low-latitude clusters subjected to a modest amount of extinction and star-crowding, the near-infrared FP can be successfully used to arrive at unbiased distance estimates. It will therefore be of great interest and importance to select suitable clusters for near-infrared FP studies to supplement, for example, the 6dF $v$-survey (Jones et al., these proceedings) at low Galactic latitudes ($|b| \\le 10^{\\circ}$). Of particular interest are clusters in and beyond the {Great Attractor} region, in order to weigh the relative contribution of the GA and the {Shapley Concentration} (see also Lucey et al. and Ebeling et al., these proceedings)."
    },
    "0410/astro-ph0410039_arXiv.txt": {
        "abstract": "We have imaged emission from the millimeter lines of eight molecules---C$_2$H, C$^{34}$S, N$_2$H$^+$, CH$_3$OH, HNCO, HNC, HC$_3$N, and SO---in the central half kiloparsec of the nearby spiral galaxy IC 342. The 5\\arcsec\\ ($\\sim 50$ pc) resolution images were made with the Owens Valley Millimeter Array (OVRO).  Using these and previously published CO and HCN images we obtain a picture of the chemistry within the nuclear region on the sizescales of individual giant molecular clouds.  Bright emission is detected from all but SO. There are marked differences in morphology for the different molecules.  A principal component analysis is performed to quantify similarities and differences among the images.  This analysis reveals that while all molecules are to zeroth order correlated, that is, they are all found in dense molecular clouds, there are three distinct groups of molecules distinguished by the location of their emission within the nuclear region. N$_2$H$^+$, C$^{18}$O, HNC and HCN are widespread and bright, good overall tracers of dense molecular gas. C$_2$H and C$^{34}$S, tracers of Photo-Dissociation Region chemistry, originate exclusively from the central 50-100 pc region, where radiation fields are high. The third group of molecules, CH$_3$OH and HNCO, correlates well with the expected locations of bar-induced orbital shocks.  The good correlation of HNCO with the established shock tracer molecule CH$_3$OH is evidence that this molecule, whose chemistry has been uncertain, is indeed produced by processing of grains.  HC$_{3}$N is observed to correlate tightly with 3mm continuum emission, demonstrating that the young starbursts are the sites of the warmest and densest molecular gas.  We compare our HNC images with the HCN images of Downes et al. (1992) to produce the first high resolution, extragalactic HCN/HNC map: the HNC/HCN ratio is near unity across the nucleus and the correlation of both of these gas tracers with the star formation is excellent.  The ratio exhibits no obvious correlation with gas temperature or star formation strength. ",
        "introduction": "}  The bright and abundant molecule CO has dominated the study of molecular clouds in external galaxies.  The millimeter-wave transitions of CO and its isotopomers are powerful probes of diffuse molecular gas \\citep[eg.][]{YS91}.  However, emission from the optically thick and easily excited CO tends to favor the warmer, radiatively lit and diffuse surface layers of clouds \\citetext{eg., \\citealp{THH93}; Meier, Turner \\& Hurt 2000, hereafter \\citeauthor{MTH00}}.  Since dense gas is the component most closely connected to star formation \\citep*[eg.,][]{GS04}, to understand the links between molecular clouds and star formation in different galactic environments, we need to study tracer molecules appropriate to a cooler, dense component that may not be well traced by CO.  Each molecular line traces a distinct regime of density and temperature within a molecular cloud, while different molecules can trace different gas chemistries \\citep*[see][for a review of Galactic astrochemistry]{VB98}.  Emission from high density tracers such as HCN, HCO$^{+}$ and CS show that large amounts of dense ($\\gsim ~ 10^{4}$ cm$^{-3}$) molecular gas are present in the centers of galaxies \\citep*[eg.][]{MH89,MHWH89,NNJ89,NJHTM92,SDR92,HB93}, and that the dense molecular interstellar medium (ISM) can vary significantly on scales of a few tens of parsecs \\citep*[eg.][]{Dow92,BS93,HB97,SFB97,SNKN02}.  Numerous molecules have been detected in nearby starbursts \\citep*[eg.,][]{HJMMS87,HSM88,NHJM91,MHWSW91,PB92,SZ95, MHC95,MMMGH03,UGFMR04,WHCWHMM04}.  These studies find that the chemical differentiation seen within Galactic molecular clouds survives to scales of hundreds of parsecs, and even to galaxy-wide differences.  However, most of these observations have been done with single dish telescopes, at resolutions of $\\sim$30\\arcsec, or a few hundred pc on the galaxy, and thus average many giant molecular clouds (GMCs) together into one beam.  With millimeter interferometers it is now possible to resolve individual GMCs in the nearest galaxies with sufficient sensitivity to allow the study of selected chemical species \\citep*[eg.,][]{GMFN00,GMFN01,GMFUN02}.  Because of the low ($\\sim 50-100$ pc) physical resolution, extragalactic observations are insensitive to small-scale ($\\sim$tenths of pc) chemistry typically studied in the Galaxy, but such maps can resolve the chemical properties operating over the bulk of a GMC and between individual GMCs.  These observations can provide insights regarding the extent to which distinct large-scale properties, such as starbursts, shocks, bars, spiral arms, tidal forces and AGN influence the chemistry of molecular clouds and how these influences are transmitted from the GMC-scale to the galaxy as a whole.  We have surveyed the nearby Scd galaxy, IC 342, in transitions from eight astrochemically important molecules. C$_{2}$H, HNCO, HNC, HC$_{3}$N, N$_{2}$H$^{+}$, C$^{34}$S, CH$_{3}$OH, and SO were imaged at $\\sim$50 pc resolution with the Owens Valley Millimeter Array (OVRO).  These are the first published interferometer maps of these lines in an external galaxy.  IC 342 is the nearest (D $\\sim$ 2 Mpc)\\footnote{Due to its location behind the Galactic plane, IC 342's distance is still a matter of debate. Recent studies have argued for a distance of $\\sim$3 Mpc \\citep*[eg.,][]{SCH02,KSDG03}, but we adopt the shorter distance to keep consistent with our previous work.}, gas-rich spiral with active star formation in its nucleus \\citep*[][]{Bec80,TH83}, and the first galaxy to be mapped with millimeter interferometers \\citep{L84}.  There is widespread resolvable molecular gas distributed in both dense clouds, with masses similar to SgrB2, and a diffuse medium about the size of the central molecular zone \\citep{MS96} in the Milky Way.  Detailed knowledge of the H$_{2}$ column densities, excitation temperatures and densities in IC 342 exists from CO and its isotopomers as a basis for comparison with these new lines \\citetext{\\citealp{I90,Dow92,TH92,WITHL93,THH93}; \\citeauthor{MTH00}; Meier \\& Turner 2001, hereafter \\citeauthor{MT01}}. ",
        "conclusions": "}  We have detected emission from a number of molecules in the central region of IC 342 at resolution of 5\\arcsec\\ using OVRO.  Significant differences are seen in the spatial distributions of these molecules, and we have analyzed the causes of these differences.  Main results include:  1. A principal component analysis of the maps of the seven detected molecules, $^{12}$CO(1-0), $\\rm C^{18}O$, HCN(1-0) and 3mm continuum, reveals that the maps all have some degree of correlation, and that the dominant common factor (PC1 axis) is the density-weighted mean column density. Molecules with the largest projections on PC1 are $\\rm N_2H^+$, $\\rm C^{18}O$, and HNC. These molecules, and primarily $\\rm N_2H^+$, give the best overall representation of the total molecular gas distribution.  2. Differences in excitation and critical density are not the dominant influences shaping the spatial differences among the different molecules in IC342, since the excitation energies of the transitions are similar, and cloud densities are high enough that most molecules are collisionally excited.  3. The spatial differences among the molecular line emission are not driven by variations in chemical timescale: all clouds in this region are likely to have early-time chemistry due to the short orbital timescales of the nucleus.  4. The second principal component axis, PC2, splits into two anti-correlated groups of molecules: molecules that peak within the central 75 pc  near GMC A and the nuclear star cluster, and those that peak up along the prominent northeastern arm, located 100-200 pc away from the nucleus.  5. The GMC A peakers, $\\rm C_2H$ and C$^{34}$S, are enhanced in the center of IC 342 by high radiation fields. GMC A is a PDR cloud. The fact that these molecules tend to favor the parts of the clouds towards the center of IC 342 rather than toward the slightly off-center $\\rm L_{IR}\\sim10^8~L_\\odot$ radio/IR source suggests that the source of the radiation field affecting the chemistry of GMC A is the $\\sim$60 Myr central star cluster and not the actively star forming radio/IR sources.  6. Molecules that peak along the north-east arm of the molecular mini-spiral are HNCO and $\\rm CH_3OH$.  Large scale shocks due to the changes in gas velocity at the arms of the nuclear minispiral are probably the cause of enhanced HNCO and $\\rm CH_3OH$ here.  7. HCN and HNC are the best tracers of the dense gas and HC$_{3}$N of warm, dense gas. Their correlation with the 3mm continuum is excellent. The HCN/HNC ratio is 1-2 across the nuclear region, and appears to be unrelated to kinetic temperature.  8. The brightness and extent of emission from $N_2H^+$ suggests that the cosmic ray ionization rate is higher than $\\zeta ~> ~ 10^{-17}$ s$^{-1}$ in the nucleus of IC~342."
    },
    "0410/astro-ph0410380_arXiv.txt": {
        "abstract": "We make use of the on-going All Imaging Survey of the UV GALEX satellite cross-correlated with the IRAS all sky survey to build samples of galaxies truly selected in far-infrared or in ultraviolet. We discuss the amount of dust attenuation and the star formation rates for these galaxies and compare the properties of the galaxies selected in FIR or in UV. ",
        "introduction": "We have worked on 600 deg$^2$ observed by GALEX in NUV( $\\lambda = 231$ nm and FUV ($\\lambda =  153 $ nm)  to build two samples of galaxies. The first one, called the {\\sl NUV selected sample} includes all the galaxies brighter than m(NUV) = 16 mag (AB scale), among the 88 selected galaxies (excluding ellipticals and active galaxies) only 3 are not detected by IRAS at 60 $\\mu m$. The second sample, called the {\\sl FIR selected sample} is based on the IRAS PSCz (Saunders et al. \\cite{saundersetal}): 118 galaxies from this catalog lie within our GALEX fields, only 1 is not detected in NUV. ",
        "conclusions": ""
    },
    "0410/astro-ph0410663_arXiv.txt": {
        "abstract": "{We present new, simultaneous, multifrequency observations of 45 out of the 55 candidate High Frequency Peakers (HFP) selected by Dallacasa et al. (2000), carried out 3 to 4 years after a first set of observations. Our sub-sample consists of 10 galaxies, 28 stellar objects (``quasars'') and 7 unidentified sources. Both sets of observations are sensitive enough to allow the detection of variability at the 10\\% level or lower. While galaxies do not show significant variability, most quasars do. Seven of them no longer show the convex spectrum which is the defining property of Gigahertz Peaked Spectrum (GPS)/HFP sources and are interpreted as blazars caught by Dallacasa et al. (2000) during a flare, when a highly self-absorbed component dominated the emission. In general, the variability properties (amplitude, timescales, correlation between peak luminosity and peak frequency of the flaring component) of the quasar sub-sample resemble those of blazars. We thus conclude that most HFP candidates identified with quasars may well be flaring blazars. ",
        "introduction": "The Compact Steep Spectrum (CSS) and Gigahertz Peaked Spectrum sources (GPS) are two classes of intrinsically compact objects (linear size $<10$--1 kpc) defined on the basis of their spectral properties: the overall shape is convex with turnover frequencies between 100 MHz (CSS) and a few GHz (GPS) and the spectral index at high frequencies is steep (see O'Dea 1998 for a complete review).  The currently accepted model (the {\\textit{youth scenario}}) relates the small linear size of GPS/CSS sources to their age, implying that they are the progenitors of extended radio sources. Both kinematic studies (Owsianik \\& Conway 1998; Owsianik, Conway \\& Polatidis 1998) and spectral analysis (Murgia et al. 1999) indicate ages of $10^3$--$10^5$ yr. On the other hand, various observations across the electromagnetic spectrum (e.g. Fanti et al. 2000) failed to reveal the presence of a particularly dense medium that could prevent the expansion of the radio lobes in the  intergalactic medium, as proposed in the {\\textit{frustration model}} (van Breugel 1984; Baum et al. 1990). So far no evidence of differences in the properties of the host galaxies between GPS/CSS and FRII radio sources has been reported (e.g. see Fanti et al. 2000).  Observational studies of the population of GPS/CSS sources have led to the discovery  of an anti-correlation between the radio turnover frequency, $\\nu_p$, and the projected angular size, $\\theta$ (O'Dea \\& Baum 1997; Fanti et al. 2002). This relationship is expected if the peaked component is due to synchrotron self-absorption, where $\\theta^2\\propto \\nu_p^{-5/2}$ , although free-free absorption could also play a role (Bicknell, Dopita \\& O'Dea 1997). In the youth scenario this correlation means that the youngest objects have the highest turnover frequencies and that the peak frequency is expected to move towards lower frequencies as the source expands and the energy density decreases. In principle, the highest the turnover frequency, the youngest the radio source is.  The shape of the radio spectrum and the position of the turnover have been used as selection tools for this class of objects. Dallacasa et al. (2000) selected a sample of ``bright'' radio sources with turnover frequencies above 5 GHz and called them High Frequency Peakers (HFPs). An early summary of their properties can be found in Dallacasa (2003). Simultaneous multifrequency VLA observations were made to define the spectral shape of all candidates. The comparison with literature data showed that flux density variability is not uncommon among the candidates and that also beamed radio sources like blazars can meet the selection criteria of the sample, e.g. when a flaring, strongly self-absorbed synchrotron component dominates the emission spectrum. A substantial contamination of the HFP sample by blazars may obviously lead astray analyses of HFP properties.  In this paper we present the results of a second set of simultaneous multifrequency VLA observations of the sample of bright HFPs, aimed at discriminating between bona-fide HFPs and blazars, and at studying the spectral evolution of the sources. At 1.465 and 1.665 GHz, longer observations were carried out to search for extended emission around the point-like dominant component. It is known that a significant fraction ($\\sim 10\\%$) of GPS sources do have extended emission (e.g. Stanghellini 2003). Also polarization data have been taken and will be presented elsewhere.  This paper is structured in the following way: Section 2 describes the radio properties of the sample, the second-epoch VLA observations and the data reduction; Section 3 presents the analysis of simultaneous radio spectra; Section 4 deals with the extended emission; in Section 5 the variability properties are investigated; the main conclusions are summarized and discussed in Section 6. Throughout this paper, we adopt $H_0=50 km s^{-1} Mpc^{-1}$ and $\\Omega=1$   \\begin{figure} \\centering \\includegraphics[width=8 cm]{dist_z_55.ps} \\caption{Distribution of measured redshifts for the sample by Dallacasa et al. (2000).} \\label{fig:dist_redshift} \\end{figure} ",
        "conclusions": "The spectral selection adopted by Dallacasa et al. (2000) is independent of optical identification and redshift. As a result, both galaxies and star-like objects (``quasars'') are present in the original sample of bright HFP candidates. Quasars are almost three times more numerous than galaxies, consistent with the previously noted decrease of the galaxy to quasar ratio in GPS samples selected at increasing turnover frequency (Fanti et al. 1990; O'Dea 1998; Stanghellini et al. 1998, 2003).  However, our new multifrequency observations, with a time lag of 3 to 4 years, indicate that the sample of candidate HFP quasars is likely to be strongly contaminated by beamed objects, with variability properties consistent with those of blazars.  Seven out of the 28 star-like sources (25\\%) no longer have a peaked spectrum, and 5 of them have extended emission. Such sources are therefore classified as blazars, caught by Dallacasa et al. (2000) during a flaring phase of a strongly self-absorbed component.  Seven additional star-like sources show a variation of the peak frequency, significant at a $>3\\sigma$ level, larger than expected if HFP/GPS sources expand at mildly relativistic velocities ($\\Delta \\nu_p/[\\nu_p \\Delta t({\\rm yr})]> 0.05$), suggesting the presence of relativistic beaming effects. The only galaxy (1407+2827) with a peak frequency variation significant at $>3\\sigma$ level has $\\Delta \\nu_p/[\\nu_p \\Delta t({\\rm yr})]= 0.017$.  Another criterion to discriminate between truly young sources and blazar is the variability index $V$ [Eq. (\\ref{eq:V})]. The distribution of $V$ for the less variable sources is close to the $\\chi^2$ distribution with two degrees of freedom, as may be expected since variations of spectra of young sources are controlled by two parameters, $S_p$ and $\\nu_p$. If so, sources with $V>9$ have a probability $<0.01$ of being extracted from such a distribution and are therefore likely blazars. We have 20 sources with $V>9$, 17 of which are star-like (including the 7 with flat spectrum at the second epoch and 4 of the 7 with $\\Delta \\nu_p/[\\nu_p \\Delta t({\\rm yr})]> 0.05$) and 3 unidentified sources. The median r.m.s. variations for these sources are $\\simeq 30\\%$ ($27\\%$ for quasars only), close to those found for blazars ($\\simeq 32\\%$; Ciaramella et al. 2004).  The variability timescales in the frequency range 5--15 GHz of bright blazars, derived from their structure function, are $\\simeq 2\\,$yr in the source frame (Hughes et al. 1992; Ciaramella et al. 2004). To convert the time-lag between the two sets of observations into the mean time interval in the source frame, we need to divide it by $(1+z_{\\rm mean})$. Only 11 of the 17 starlike sources with $V>9$ have measured redshift, with a median value of 1.7. We may thus expect that $\\sim 50\\%$ of blazars flaring at the moment of the first set of observations have recovered their baseline flat spectrum by the epoch when they where re-observed. For comparison, 7 out of 17 (i.e. $\\simeq 41\\%$) were indeed found to have a flat spectrum in the second observing run.  There is evidence of a positive correlation between $V$ and $\\nu_p$ (probability of no correlation $\\sim 0.01$), indicating that the fraction of blazars is higher among objects with higher peak frequency. This is not surprising because young objects become increasingly rare since their lifetimes decrease with increasing $\\nu_p$.  We have found a statistically significant positive correlation between the peak luminosity and rest-frame peak frequency of quasars, after having removed the effect of the redshift distribution (Kendall's partial correlation test). A relationship among these two quantities is expected both in the case of GPS/HFP sources and of blazars, but with a much steeper slope for the latter sources. While previous studies enlightened a $S_{p}$--$\\nu_{p}$ correlation for a few individual blazars (e.g., Stevens et al. 1996) our analysis indicates that it holds for flaring blazars as a class.  There are thus various pieces of evidence, albeit circumstantial, that most quasar HFP candidates are actually objects where the effects of beaming are relevant (blazar like) even if this does not necessarily require that we are seeing evolved/old objects.  While sources with $V>9$ are likely blazars there is no guarantee that those with $V<9$ are truly young sources. Some of them may really be blazars in a relatively quiescent phase. The 3 such objects with $\\Delta \\nu_p/[\\nu_p \\Delta t({\\rm yr})]> 0.05$ may belong to this category.  As for galaxies, all 10 HFP candidates have $V<9$. Their peak luminosities and rest-frame peak frequencies do not show evidence of a correlation, not surprisingly given the smaller range spanned by both quantities (compared to the case of quasars) and the weaker dependence of $L_p$ on $\\nu_p$ expected for GPS/HFP sources, easily swamped by intrinsic dispersions of both quantities.  Two galaxies show extended emission compatible with recurrent activity. According to Begelman's model the synchrotron emission decreases with source age $t$ as $L\\propto t^{-\\eta}$ with $\\eta=(n+4)/[4(5-n)]$. In the case of recurrent activity, the HFP to extended emission age ratio can then be estimated as $t_{\\rm HFP}/t_{\\rm ext}=(L_{\\rm ext}/L_{\\rm HFP})^{1/\\eta}$, where $L_{\\rm ext}$ and $L_{\\rm HFP}$ are the total {\\it emitted} radio luminosities, that can be taken as proportional to the {\\it emitted} flux densities at 1.4 GHz. In the case of the HFP component, the latter can be estimated extrapolating to 1.4 GHz the peak flux density. The age ratio derived from Begelman's model depends on the slope $n$ of the density profile of the surrounding medium. Adopting $n=2$ we get, for J$0111+3906$, $t_{\\rm ext}/t_{\\rm HFP}\\simeq 9\\times 10^4$, and for J$0428+3259$ $t_{\\rm ext}/t_{\\rm HFP}\\simeq 2\\times10^4$, consistent with the intermittency timescales suggested by Reynolds \\& Begelman (1997). Although any conclusion at this stage is premature, this illustrates the potential of deeper low frequency observations to shed light on the nature of radio activity.  In conclusion, while the variability properties of candidate HFP galaxies may be consistent with expectations for truly young, possibly recurrent, sources, those of most quasars are consistent with those of blazars. If we set the boundary between HFP candidates and blazars at $V=3$ (see Sect.~\\ref{analysis}), 3 galaxies would host a blazar nucleus, and 22 out of the 28 starlike sources and 4 out of the 7 unidentified sources would be blazar candidates. The median variability amplitude of sources with $V>3$ is $\\simeq 22\\%$, slightly lower than, but still not far from that found for blazar samples. Milli-arcsec morphology at different frequencies, polarization measurements and possibly the study of the spectral index distributions in the optically thin regime can help distinguishing between flaring blazars and truly young objects.  High resolution VLBA observations at two frequencies in the optically thin part of the spectrum (between 8.4 and 43.2 GHz) of 50 objects of the original sample have already been acquired. The analysis of the data is in progress."
    },
    "0410/astro-ph0410455_arXiv.txt": {
        "abstract": "{The extremely hot and tenuous accretion flow in the immediate vicinity of \\sgra\\ is believed to be invisible (too dim) in the X-ray band, except for short X-ray flares. Here we point out that during pericenter passages, close brightest stars irradiate the inner region of the accretion flow, providing a plenty of optical/UV photons. These seed photons are Compton up-scattered by the hot electrons of the accretion flow to higher frequencies, some into the X-ray band, potentially making the innermost accretion flow much brighter in X-rays than usual. We propose to use coordinated near infra-red and X-ray observations of close star passages to put constraints onto \\sgra\\ accretion theories.  The absence of a noticeable change in the steady emission of \\sgra\\ as observed by {\\em Chandra} in the year 2002, when the star named S2 passed through a pericenter of its orbit, already rules out the hotter of the ``standard'' Advection-Dominated Accretion Flows. The less dense accretion flows, in particular the model of \\cite{Yuan03}, passes the test and is constrained to accretion rates no larger than $\\sim$~few$\\times 10^{-7} \\msun$~year$^{-1}$. ",
        "introduction": "\\sgra\\ is the super massive black hole (SMBH) in the center of our Galaxy \\citep{Reid99, Schoedel02, Ghez03a}. The gas from the winds of the surrounding young massive stars should be able to maintain the accretion rate on the SMBH at a sizable fraction of the Bondi accretion rate {\\em if the SMBH can accept the gas at this rate}. The Bondi accretion rate is estimated to be $\\sim 10^{-6} \\msun$ year$^{-1}$ \\citep{Baganoff03a}, whereas the linear polarization measurements at a range of radio wavelengths constrain the accretion rate onto the SMBH to a value of $\\sim 10^{-8}-10^{-7} \\msun$~year$^{-1}$ \\citep[e.g.,][]{Bower03}. The latter estimate is based on somewhat uncertain assumptions about the magnetic field equipartition and absence of many magnetic field reversals \\citep{Ruszkowski02}.  It would certainly be beneficial to have another physically independent method to constrain the properties of the accretion flow near the event horizon. Unfortunately, in the X-ray band, it appears that the {\\em quiescent} \\sgra\\ emission is resolved \\citep{Baganoff03a} and can be explained by thermal bremsstrahlung from the capture radius region \\citep{Quataert02}. The inner flow may thus be simply invisible to the observer unless there is a sizable X-ray emission from a jet \\citep[e.g.,][]{Yuan02}, or during strong X-ray flares \\citep{Baganoff01}.  In this Letter we note that comptonization of the radiation of the closest of the observed bright stars by the hot accretion flow may be used to this effect. These stars produce as much as $L_*\\sim 10^5 L_\\odot \\sim 4 \\times 10^{38}$ erg/sec in optical-UV radiation, the amount far greater than the total bolometric luminosity of \\sgra, $L_t \\sim 10^{36}$ erg/sec. In a single Compton scattering, the photon energy on average increases by a factor $\\gamma^2$, where $\\gamma$ is the Lorenz $\\gamma$-factor of the electron taking part in the up-scattering. Averaged over a relativistic Maxwellian electron distribution with temperature $T_e$, the average boost factor becomes $16 (k_B T_e/m_e c^2)^2$, where $k_B$ is the Boltzmann constant, and $m_e$ is the electron mass. The typical frequency of a photon emitted by the star is around $\\nu = 3 k_B T_*/h \\approx 2 \\times 10^{15}$ Hz, where $T_* = 30,000$ Kelvin is the stellar temperature, and $h$ is the Planck constant. {\\em Chandra} frequency band is centered around $\\nu \\sim 10^{18}$ Hz.  Therefore electrons with temperature of $T_e \\approx 3\\times 10^{10} K$ are the most efficient in scattering of S2 radiation into the {\\em Chandra} band.  Such high temperature electrons are present in both the canonical ADAF (without winds) model, e.g. \\cite{Narayan02}, and even more so in the less dense but hotter Non-Radiative Accretion Flows (NRAFs) that incorporate effects of gas outflows. Since \\sgra\\ is very dim in X-rays in its normal ``quiescent'' state, e.g. $L_X \\sim 10^{33}$~erg~s$^{-1}$, there is then a hope of detecting a comptonized ``echo'' of a close star passage. ",
        "conclusions": "\\label{sec:discussion}  We have shown that up-comptonization of the optical/UV radiation of the star S2, which famously passed through the pericenter of its orbit only $\\sim 2000$ \\schw radii away from \\sgra\\ \\citep{Schoedel02}, constrains the properties of the accretion flow in the immediate vicinity of the SMBH. These constraints are completely independent from those due to the rotation measure limits in the radio wavelengths \\citep{Bower03}, and are thus complimentary to those results. We urge the observers to search the {\\em Chandra} 2002 data for signatures of the transient brightening with the lightcurve shown in Figure \\ref{fig:fig2}. Such an analysis may perhaps yield much tighter constraints on the amplitude of the effect than Figure \\ref{fig:fig1} indicates.  Future observations of close star passages may be more constraining that that from the year 2002 if (i) the star approaching the SMBH is brighter and/or closer to \\sgra\\ than $2000 R_S$; (ii) if more sensitive and well planned X-ray observations of the close stellar passage are made. In addition, interaction of stellar winds with the surrounding hot accretion flow may yield observable signatures in a range of wavelengths depending on the stellar mass loss rate and the outflow velocity and (unfortunately) not well constrained details of presumably collisionless wind/accretion flow shocks (e.g., see \\cite{Loeb04}).  The idea of using the close star passages to constrain the properties of the accretion flow around \\sgra\\ is not new. \\cite{Chang03} noted that the inverse Compton effect will provide an additional cooling of the hot electrons. However the corresponding relative change in the electron temperature is only $\\simlt 10^{-4}$ and appears very hard to observe in any wavelength. Here we showed that the up-comptonized S2 radiation itself, on the over hand, is much easier to observe since \\sgra\\ quiescent X-ray luminosity is very small.  In addition, \\cite{NS03} and \\cite{Cuadra03} have shown that passages of close stars could also be used to constrain the properties of a cold geometrically thin ``inactive'' accretion disk around \\sgra. The fact that no near infra-red ``echoes'' or stellar eclipses were observed up to date probably indicates that such a disk does not presently exist in \\sgra.  Concluding, we note that close stellar passages are of significant value for solving the riddle of \\sgra\\ and collisionless gas accretion onto SMBHs in general through a variety of physical processes.    \\del{Density of the wind: \\begin{equation} n_w \\sim M_w/4\\pi R^2 v_w m_p \\sim 5\\times 10^4 cm^{-3} \\dot M_{-6} r_3^{-2} v_8^{-1} \\end{equation} Mass of the wind: \\begin{equation} M \\sim 10^{-6} \\msun/year \\times 0.1 year = 2 \\times 10^{26} g \\end{equation} Mass of the non-radiative flow $\\sim 5 (10^{15})^3 \\times 2 10^4 \\times 2 10^{-24} = 2 \\times 10^{26}$ g also.}   \\acknowledgement The author acknowledges useful discussions with Fred Baganoff and Henk Spruit."
    },
    "0410/hep-ph0410356_arXiv.txt": {
        "abstract": "We have studied the cosmological evolution of domain wall networks in two, three and four spatial dimensions using high-resolution field theory simulations. The dynamical range and number of our simulations is larger than in previous works, but does not allow us to exclude previous hints of deviations to the naively expected scale-invariant evolution. These results therefore suggest that the approach of domain wall networks to linear scaling is a much slower process than that of cosmic strings, which has been previously characterized in detail. ",
        "introduction": "Introduction}  The cosmological consequences of primordial phase transitions that are thought to have happened in the early universe have been the subject of many studies in recent years. One of these inevitable consequences is the formation of topological defects \\cite{Kibble,Book}. While most such studies have focused on cosmic strings, domain walls can be of interest too, although observational constraints rule them out if their symmetry breaking scale is $\\eta\\ge 1MeV$ \\cite{Zeldovich}. However, note that this constraint is on the assumption that they are topologically stable and that their cosmological evolution is scale-invariant (\\textit{id est}, that the network correlation length is proportional to the horizon size, or equivalently that the total area density is inversely proportional to it).  In this work we want to highlight the fact that, in spite of some simple analytic arguments in favor of it \\cite{Hindmarsh}, this hypothetical scale-invariant solution has not yet received any strong backing from numerical simulations. In fact, in previous simulations of cosmological domain wall evolution \\cite{Press,Coulson,Larsson,Fossils,Garagounis} some hints for deviations to a scale-invariant evolution were found. The question we wish to address presently is whether these deviations remain if one performs larger runs.  We report on a large set of high-resolution simulations of domain walls in two, three and four spatial dimensions, using the standard Press-Ryden-Spergel (PRS) algorithm \\cite{Press}. Our boxes are typically larger (hence can be evolved for a longer dynamic range) that those in previous works. Gains in statistical accuracy also come from performing large numbers of runs. We highlight the differences between running a large number of small runs, or a small number of larger ones. We do find some statistical evidence for deviations to the scale-invariant behaviour, though there is some dependency on dimensionality and box sizes. While 2D and 3D numerical simulations have been previously carried out (though usually in smaller boxes or with smaller dynamic range) by a number of authors \\cite{Press,Coulson,Larsson,Fossils,Garagounis}, we have also carried out for the first time a series of 4D simulations, which in a simple phenomenological way may be of interest to brane world scenarios \\cite{tye0,Brax}.  In what follows we start with a brief description of the PRS algorithm and then proceed to present some tests of our code and the main results of our simulations. We also compare these with the work of previous authors, and highlight issues having to do with the slightly different methods used in each case. Finally, we briefly comment on the cosmological implications of our results. ",
        "conclusions": ""
    },
    "0410/hep-ph0410026_arXiv.txt": {
        "abstract": "Bright bursts of gamma rays from outer space have been puzzling Astronomers for more than thirty years and there is still no conceptually complete model for the phenomenon within the standard model of particle physics.  Is it time to consider a supersymmetric (SUSY) origin for these bursts to add to the astronomical indications of supersymmetry from dark matter? ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410740_arXiv.txt": {
        "abstract": "{ We investigate the effect of  a density gradient on Chandrasekhar's dynamical friction formula based on the method of 2-body encounters in the local approximation. We apply these generalizations to the orbit evolution of satellite galaxies in Dark Matter haloes. We find from the analysis that the main influence occurs through a position-dependent maximum impact parameter in the Coulomb logarithm, which is determined by the local scale-length of the density distribution. We also show that for eccentric orbits the explicit dependence of the Coulomb logarithm on position yields significant differences for the standard {\\it homogeneous} force. Including the velocity dependence of the Coulomb logarithm yields ambigous results. The orbital fits in the first few periods are further improved, but the deviations at later times are much larger. The additional force induced by the density gradient, the {\\it inhomogeneous} force, is not antiparallel to the satellite motion and can exceed 10\\% of the homogeneous friction force in magnitude. However, due to the symmetry properties of the inhomogeneous force, there is a deformation and no secular effect on the orbit at the first order. Therefore the inhomogeneous force can be safely neglected for the orbital evolution of satellite galaxies. For the homogeneous force we compare numerical N-body calculations with semi-analytical orbits to determine quantitatively the accuracy of the generalized formulae of the Coulomb logarithm in the Chandresekhar approach. With the local scale-length as the maximum impact parameter we find a significant improvement of the orbital fits and a better interpretation of the quantitative value of the Coulomb logarithm. ",
        "introduction": "}  Dynamical friction has a wide range of applications in stellar-dynamical systems. It appears as the first order term in the Fokker-Planck approximation (see  e.g. Binney \\& Tremaine \\cite{bin87}, hereafter BT, Eq. (8-53)) competing with the second order diffusion terms to allow for a Maxwellian as the equilibrium distribution. In the field of relaxation processes (dissolution of star clusters, gravo-thermal collapse, mass segregation) the physical constraints and dominating aspects of dynamical friction are very different to those of the dynamical evolution of a single massive body in a sea of light background particles. In the latter case the friction force is independent of the mass of the background particles, and as long as the massive body is far from energy equipartition the diffusion can be neglected. Its main applications are the satellite galaxies in the dark matter halo (DMH) of their parent galaxy, and super-massive black holes (SMBH) or compact star clusters in the central bulge region of galaxies. Even in a collisional gas, dynamical friction occurs, but corrections due to pressure forces must be applied (Just et al. \\cite{jus86}; Just \\& Kegel \\cite{jus90}; Ostriker \\cite{ost99} using mode analysis, and S\\'anchez-Salcedo \\& Brandenburg \\cite{san99,san01} using numerical calculations).  There are different approaches to analyse the parameter dependence and to calculate the magnitude of the friction force, most of them based on perturbative methods (2-body collisions: Chandrasekhar \\cite{cha43}; Binney \\cite{bin77}; Spitzer \\cite{spi87}; BT; mode analysis: Marochnik \\cite{mar68}; Kalnajs \\cite{kal72}; Tremaine \\& Weinberg \\cite{tre84}; theory of linear response: Colpi et al. \\cite{col99}; statistical correlations in  the fluctuating field: Maoz \\cite{mao93}). Due to its complicated structure, dynamical friction shows many different features. In a recent theoretical investigation Nelson \\& Tremaine (\\cite{nel99}) connected many physical aspects of the fluctuating gravitational field on the basis of the linear response theory. For practical applications one basic question is whether dynamical friction acting on a massive object is dominated by local processes. There is a longstanding discussion on the contribution of global modes to dynamical friction, which are known to be excited and enhanced by the self-gravity of the perturbation (Weinberg \\cite{wei86,wei89,wei93}; Hernquist \\& Weinberg \\cite{her89}). These strong resonances may reduce the dissipation rate considerably (Kalnajs \\cite{kal72}; Tremaine \\& Weinberg \\cite{tre84}; Weinberg \\cite{wei89}), but see also Rauch \\& Tremaine (\\cite{rau96}) for the opposite effect. For interacting galaxies with moderate mass ratios,  the global deformation of the main galaxy and the strong tidal forces on the perturber can enhance the energy and momentum loss by considerable factors, which cannot be described by the Chandrasekhar dynamical friction form (Prugniel \\& Combes \\cite{pru92}; Leeuwin \\& Combes \\cite{lee97} and the discussion therein). The mass of the perturber should not exceed $\\sim 5\\%$ of the mass of the host object in order to neglect both the global contribution to the dynamical friction force and the effect of deformation.  This paper is concerned with the local contribution to dynamical friction. We restrict our analysis to the low mass regime with $M<0.05\\,M_{\\rm h}$, where $M$ and $M_{\\rm h}$ are the mass of the perturber and the host object, respectively. For clearness we formulate most equations in terms of the orbital motion of satellite galaxies in the DMH of the parent galaxy and give additionally the explicit application to a singular isothermal halo. This case will also be used for the numerical analysis. Nevertheless the results can be used for more general background distributions with isotropic distribution functions and also for compact objects. As a second  application we give the corresponding equations for massive Black Holes moving in galactic centres.  The basic formula of Chandrasekhar relies on an infinite homogeneous background with an isotropic velocity distribution function. For a Maxwellian velocity distribution function and a massive object with $M\\gg m$ it is given by \\begin{eqnarray} \\vec{\\dot{v}_{\\rm hom}} &=& -C_{\\rm Ch} G_{\\rm hom}(X) \\vec{e_{\\rm v_{\\rm M}}} \\quad \\mbox{with} \\label{Cha0}\\\\ G_{\\rm hom}(X)&= & \\frac{\\ln\\Lambda}{X^2} \\left[ erf(X) -\\frac{2X}{\\sqrt{\\pi}}\\exp(-X^2)\\right]\\quad ,\\nonumber\\\\ C_{\\rm Ch} &= & \\frac{2\\pi G^2 \\rho_0 M}{\\sigma^2} \\quad\\mbox{and}\\quad X=\\frac{|v_{\\rm M}|}{\\sqrt{2}\\sigma} \\label{Cc} \\end{eqnarray} (BT Eqs. (7-14), (7-18)). Here $M$ and $v_{\\rm M}$ are mass and velocity of the massive body; $m$, $\\rho_0$, and $\\sigma$ are mass, mass density, and velocity dispersion of the ambient particles, respectively. In this formula,  only  2-body interactions between the massive particle and the background are included, adopting a straight-line motion in the unperturbed potential. Essentially all uncertainties of the relevant regime in phase space for these 2-body encounters are hidden in the Coulomb logarithm $\\ln\\Lambda$. Different perturbation methods lead to different interpretations of $\\Lambda$: (1) 2-body encounters: $\\Lambda=b_{\\rm max}/b_{\\rm min}$ with maximum and minimum impact parameter (BT) (2) Mode analysis:  $\\Lambda=k_{\\rm max}/k_{\\rm min}$ with maximum and minimum relevant wavenumber (Kalnajs \\cite{kal72}) (3) Theory of linear response: $\\Lambda=\\tau_{\\rm max}/\\tau_{\\rm min}$ with maximum and minimum time lag of the correlations (Nelson \\& Tremaine \\cite{nel99}). The most important correction to Eq. \\ref{Cha0} is due to an anisotropic velocity distribution (Binney \\cite{bin77}; Statler \\cite{sta88,sta91}). A numerical investigation of the anisotropy is discussed in another paper for satellite motion in flattened Dark Matter haloes (Pe\\~narrubia \\cite{pen03}; Pe\\~narrubia et al. \\cite{pen04}). In different approaches the inhomogeneity of the background distribution is generally included (Kalnajs \\cite{kal72}; Tremaine \\& Weinberg \\cite{tre84}; Colpi et al. \\cite{col99}; Nelson \\& Tremaine \\cite{nel99}) but a useful scheme for applications is not yet available. An explicit formula for corrections due to the local density gradient is given by Binney (\\cite{bin77}), who has shown that  the additional inhomogeneous force can be neglected for the dynamics of galaxies in galaxy clusters. Maoz (\\cite{mao93}) and follow-up papers ( S\\'anchez-Salcedo \\cite{san99a}; Dominguez-Teneiro \\& G\\'omez-Flechoso \\cite{dom98}) used the fluctuation-dissipation theory with approximations similar to those presented here to derive the energy loss due to dynamical friction. Del Popolo \\& Gambera (\\cite{pop99}) and Del Popolo (\\cite{pop03}) compute the corrections due to a density enhancement spherically symmetric with respect to the massive body, which may be used to estimate the shielding effect on the friction force.  Numerical calculations suffered for a long time from resolution problems (e.g. Lin \\& Tremaine \\cite{lin83}; White \\cite{whi83}; Bontekoe \\& van Albada \\cite{bon87}; Weinberg \\cite{wei93}) and, even in recent years, computations with particle numbers below $10^5$ are frequently used (Cora et al. \\cite{cor97}; van den Bosch et al. \\cite{bos99}; Jiang \\& Binney \\cite{jia00}; Hashimoto et al. \\cite{has03}). Low resolution raises two major problems for the determination of orbital decay times and quantitative numbers for the Coulomb logarithm $\\ln\\Lambda$. Due to noise, the numerically determined Coulomb factors  $\\Lambda$ vary significantly for the same orbits. With different seeds and evolution times for relaxation of the unperturbed galaxy using 5,000 particles Cora et al. (\\cite{cor97}) found a variation of fitted values for $\\Lambda$ by almost a factor of 3. Even in the 20,000 particle runs the single orbits differ by up to a factor of 2 in the satellite energy. Secondly, increasing the resolution should lead to systematically larger values for the Coulomb logarithm, because the range of impact parameters, where the perturbation of the motion of the background particles is resolved, increases. This is explicitely shown in Spinnato et al. (\\cite{spi03}). Therefore even an average over many low resolution computations underestimates the friction force and leads to too large decay times for the satellite galaxies. With the new generation of computers it is now possible to perform a larger number of self-consistent calculations with particle numbers above $10^6$ and the corresponding high grid resolutions in particle-mesh codes (Vel\\'azquez \\& White \\cite{vel99}; Pe\\~narrubia et al. \\cite{pen02}; Spinnato et al. \\cite{spi03}). In these calculations one can hope to suppress noise sufficiently and to also resolve the small scale perturbations to model the full range for the Coulomb logarithm. Bertin et al. (\\cite{ber03}) also used a large number of particles but only a low spatial resolution due to the restriction on low order spherical harmonics. For the excitation of resonant global modes in N-body calculations, one order of magnitude higher particle numbers may be necessary. But this is not our aim here.  The standard usage of Chandrasekhar's formula works with an individual fitting of the constant Coulomb logarithm for each orbit. Fitting the Coulomb logarithm to a single orbit provides in most cases an accurate description of a large part of the orbit. For this task, there is no need to improve the method. But the Coulomb logarithm varies systematically from orbit to orbit as discussed above and the orbital fits drift away for long-living satellites (Lin \\& Tremaine \\cite{lin83}; Jiang \\& Binney \\cite{jia00}; Spinnato et al.  \\cite{spi03}; Pe\\~narrubia et al. \\cite{pen04}). In order to get an improved friction formula, which can be used for a sample of orbits with different parameters, we apply local constraints yielding a systematic variation of the Coulomb logarithm with position (and also velocity). The distance to the galactic centre as the maximum impact parameter for the orbital evolution of the Magellanic Clouds was already proposed by Tremaine (\\cite{tre76}) and a position-dependent Coulomb logarithm was found numerically by Bontekoe \\& van Albada (\\cite{bon87}) and Bertin et al. (\\cite{ber03}). However the properties of a position-dependent $\\ln \\Lambda$ were never investigated systematically. We discuss also the velocity dependence of $\\Lambda$ by excluding slow encounters.  A second effect of the local density gradient of the background distribution is an additional 'inhomogeneous' force. Due to the different symmetry of the density gradient, the net effect of the perpendicular component of 2-body encounters does not vanish like that for the homogeneous part, but leads to an additional force not anti-parallel to the satellite motion (see also Binney \\cite{bin77}). We discuss the new Coulomb logarithm and the inhomogeneous friction term in detail.  We calculate dynamical friction using the explicit deflection of 2-body encounters, because this approach is not based on the diffusion limit in phase space. Therefore, in the case of a massive particle, the derivation also holds for large angle deflections leading to a well-defined behaviour for small impact parameters and slow encounters. Neglecting the effect of the mean tidal field on the encounter yields the well known deflection \\begin{eqnarray} \\vec{\\Delta v_{\\rm M\\|}} & = & \\frac{2m\\vec{V_{\\rm 0}}}{M} \\left[1+\\frac{b^2V_{\\rm 0}^4}{G^2 M^2}\\right]^{-1} \\nonumber\\\\ & = & \\frac{2Gm}{a(V_{\\rm 0})}\\frac{\\vec{V_{\\rm 0}}}{V_{\\rm 0}^2} \\left[1+\\frac{b^2}{a(V_{\\rm 0})^2}\\right]^{-1} \\,, \\label{deltavp}\\\\ \\vec{\\Delta v_{\\rm M\\bot}} & = & \\frac{2m\\vec{b} V_{\\rm 0}^3}{GM^2} \\left[1+\\frac{b^2V_{\\rm 0}^4}{G^2 M^2}\\right]^{-1} \\nonumber\\\\ & = & \\frac{2Gm\\vec{b}}{a(V_{\\rm 0})^2 V_{\\rm 0}} \\left[1+\\frac{b^2}{a(V_{\\rm 0})^2}\\right]^{-1} \\label{deltavo}\\\\ && \\qquad\\mbox{with}\\qquad a(V_{\\rm 0})=\\frac{GM}{V_{\\rm 0}^2} \\label{aV} \\end{eqnarray} (BT, Eqs. (7-10a,b)) for a given impact parameter $\\vec{b}$ and relative velocity $\\vec{V_{\\rm 0}}$. Here $a(V_{\\rm 0})$ corresponds to the impact parameter for a $90^{\\circ}$ deflection. From this acceleration we have to subtract the acceleration along the unperturbed orbit properly in order to avoid artificial results due to differences in the mean field approximation.  We start in Sect. \\ref{coulog} with a detailed discussion of the collision rate of 2-body encounters and the structure of the Coulomb logarithm with all parameter dependences. In Sect. \\ref{dyfric} we compute the inhomogeneous dynamical friction for an isotropic distribution function including a reanalysis of the homogeneous friction force (corresponding to the standard friction) for a consistent comparison of both parts. In Sects. \\ref{applic} and \\ref{results} we apply the new formulae to typical satellite orbits in the Dark Matter halo of the Milky Way and discuss the effect of the improved Coulomb logarithm and of the inhomogeneous friction term. Sect. \\ref{conclu} contains a discussion of the relevance of using a position-dependent Coulomb logarithm and of the inhomogeneous force for other applications (super massive black holes in galactic cores) and other aspects (circularization of the orbits). ",
        "conclusions": "}  We have recomputed the dynamical friction force on small satellite galaxies moving in the DMH of their parent galaxies. We have extended the local approach of Chandrasekhar by adding local constraints for the determination of the friction force. We analysed the effect of the local density gradient of the DMH on the standard (homogeneous) friction force and the properties of the additional inhomogeneous force term.  \\subsection{The Coulomb logarithm}  We have shown that the local scale length $L$ is an appropriate maximum impact parameter in the Coulomb logarithm $\\ln \\Lambda_0$. A numerical test for an isothermal DMH shows a significant improvement of orbits with different eccentricities compared to the standard case of a constant $\\langle\\ln\\Lambda\\rangle$. This fits to the finding of Hashimoto at al. (\\cite{has03}), who also used a position-dependent Coulomb logarithm. The numerical results of high-resolution simulations of Spinnato et al. (\\cite{spi03}) are also in agreement with $L$ as the maximum impact parameter. For a detailed analysis see Just \\& Spurzem (\\cite{jus03}). The improved Coulomb logarithm is also in agreement with the detailed investigation of Bontekoe \\& van Albada (\\cite{bon87}) for a n=3 polytrope with radius $R_{\\rm out}$. They found for the range of central distances, where the density profile is nearly exponential (corresponding to a constant scale length $L\\approx 0.1R_{\\rm out}$), constant values for $\\ln \\Lambda$, which are of the order of unity for a numerical resolution of $\\approx 0.04R_{\\rm out}$.  In general the numerical determination of $\\ln \\Lambda$ strongly depends on the numerical method and resolution. In Cora et al. (\\cite{cor97}) and other work a high level of noise and a wide spread in the numerical determination of the Coulomb logarithm are present. An extreme case is presented in Jiang \\& Binney (\\cite{jia00}), who found $\\ln \\Lambda=8.5$ from orbital fits to the Sagittarius dwarf, which corresponds to the unphysical ratio of $b_{\\rm max}/b_{\\rm min}=5,000$. They used high particle numbers but an expansion into low order spherical harmonics with respect to the halo centre, which strongly suppresses the local gravitational wake, and artificially enhances the excitation of global modes by 'ghost satellites' due to the symmetry. Cartesian meshes are best adapted to a good representation of a local perturbation with well-defined resolution - here global modes may be underestimated by the lack of general symmetry.  The generalization to a velocity-dependent Coulomb logarithm by an additional cut-off for near but slow encounters using $\\ln \\Lambda_2$ gives an ambiguous result. On the one hand we find an  additional improvement of the orbit fits for the first reolutions. But on the other hand in the late phase the orbital decay in the numerical calculations is even stronger than with a constant Coulomb logarithm. The reason may be a deformation of the central halo region due to the relative large mass ratio $q_{\\rm m}$. The linear decay when using  $\\ln \\Lambda_2$ is similar to the findings of Hashimoto et al. (\\cite{has03}). Additionally, the fitting factor $Q_2\\approx 2.5$ shows that by cutting off slow encounters a significant contribution to the friction force is lost. The rescaling by using $Q_2>1$ may change the parameter dependence essentially. Here further investigations are necessary for conclusive results.  \\subsection{Circularization}  One of the most important effects of a position-dependent $\\ln \\Lambda_0$ is the reduction of the circularization of the orbit. {The variation of $\\ln \\Lambda$ along the orbit is a sensitive quantity for the evolution of the orbital shape, since the effect is significant even in the case where the dynamical friction force is changed by about 15\\% along the orbit. The circularization depends on the systematic parameter dependence of the friction force along the orbit, because the energy loss relative to angular momentum loss is strongly enhanced around peri-centre due to the higher satellite velocity (see Eq. \\ref{edot}). The general effect of a reduced friction force around peri-centre is a reduction of the energy loss leading to a lower circularization. The effect of the position-dependent Coulomb logarithm on the circularization of the orbits is in any case important for eccentric orbits. The general finding of a constant or increasing eccentricity in N-body simulations can be understood in this way.  \\subsection{Inhomogeneous dynamical friction}  The inhomogeneous term describes the effect of the asymmetry of the gravitational wake with respect to the orbit. For the calculation of the inhomogeneous force it is essential to use the velocity-dependent $\\ln \\Lambda_1$ or $\\ln \\Lambda_2$ in order to have a robust and consistent cut-off at low impact velocities, where the dominant contribution arises (see Appendix for a comparison with Binney (\\cite{bin77}), who used a lower cut-off velocity). The inhomogeneous force is for typical parameters of the order of 10\\% relative to the homogeneous force. In highly eccentric orbits around apo-centre and for higher satellite masses ($\\sim 10^{10}M_{\\rm \\odot}$) it becomes significantly stronger. The force ratio of inhomogeneous and homogeneous terms scales essentially with $\\sim M/{\\cal M}(r_{\\rm M})$ (the ratio of satellite mass and enclosed mass of the DMH). This is very similar to the ratio of the orbital to the friction time-scale.  The properties of the inhomogeneous term are very different to the standard friction force. For an isotropic velocity distribution it is an even function of the satellite velocity leading to a non-vanishing force at $v_{\\rm M}=0$ and to a non-secular behaviour around apo- and peri-centre. The inhomogeneous force is less inclined to the satellite motion than the density gradient. For vanishing satellite motion it points anti-parallel to the density gradient. The parallel component gives an acceleration during outward motion and a deceleration during inward motion. The net force points outwards of the orbit leading to a less eccentric orbit.  In our test orbits, the effect of the inhomogeneous force is very small leading to an orbital offset of the order of 1~kpc over a few Gyr. Since the magnitude of the inhomogeneous force is so small and their is no secular effect to first order, the corrections to the friction force must be compared to other shortcomings of the dominating homogeneous force. In highly eccentric orbits the time lag of the perturbation leads to a stretching or squashing at peri-centre and a bending at apo-centre of the gravitational wake due to the accelerated unperturbed motion of the satellite. The corresponding corrections may affect the orbital evolution much more than the inhomogeneous force.  \\subsection{SMBHs in galactic centres}  For compact objects an additional velocity dependence of the Coulomb logarithm via $a_{90}$, the typical impact parameter for a $90^0$ deflection, appears in $\\ln\\Lambda_0$. For an application to the orbital decay of super-massive Black holes (SMBH) in cuspy galaxy centres, a significant correction to the decay time-scales is required. A quantitative analysis of the motion of SMBHs in galactic centres is given in Just \\& Spurzem (\\cite{jus03})."
    },
    "0410/astro-ph0410095_arXiv.txt": {
        "abstract": "From millimeter observations of classical T Tauri stars, it is suggested that dust grains in circumstellar disks have grown to millimeter size or larger. However, gas drag on such large grains induces rapid accretion of the dust. We examine the evolution of dust disks composed of millimeter sized grains, and show that rapid accretion of the dust disk causes attenuation of millimeter continuum emission. If a dust disk is composed mainly of grains of $1 \\ {\\rm cm}-1$ m, its millimeter emission goes off within $10^6$ yr. Hence, grains in this size range cannot be a main population of the dust. Considering our results together with grain growth suggested by the millimeter continuum observations, we expect that the millimeter continuum emission of disks comes mainly from grains in a narrow size range of $[1 \\ {\\rm mm}-1 \\ {\\rm cm}]$. This suggests either that growth of millimeter sized grains to centimeter size takes more than $10^6$ yr, or that millimeter sized grains are continuously replenished. In the former case, planet formation is probably difficult, especially in the outer disks. In the latter case, reservoirs of millimeter grains are possibly large ($\\ga 10$ m) bodies, which can reside in the disk more than $10^6$ yr. Constraints on the grain growth time-scale are discussed for the above two cases. ",
        "introduction": "Observations at millimeter wavelengths ($\\lambda=1-3$ mm) have revealed that about 50\\% of classical T Tauri stars (CTTSs) emit detectable millimeter continuum. The detection limit, $10-20$ mJy, of the $1.3$ mm observations by Osterloh \\& Beckwith (1995) corresponds to a minimum disk luminosity of $4 \\pi D^2 \\nu F_{\\nu} \\approx 5 \\times 10^{28} \\ {\\rm erg \\ s^{-1}}$, where we used the distance to the Taurus-Auriga molecular cloud, $D=140$ pc, while the typical value of the disk luminosity is $10^{29} -10^{31} \\ {\\rm erg \\ s^{-1}}$ (Beckwith et al. 1990; Ohashi et al. 1991, 1996). The sizes of the disks that emit millimeter continuum are in a range between $100-500$ AU (Dutrey et al. 1996; Kitamura et al. 2002). Only a small fraction ($\\sim 10$\\%) of weak line T Tauri stars (WTTSs), however, have detectable emission at $1.3$ and $2.7$ mm (Osterloh \\& Beckwith 1995; Dutrey et al. 1996; Duvert et al. 2000). It is not clear whether the typical age of WTTSs is older than that of CTTSs, as both classes have similar age distributions between $10^5$ and $10^7$ yr (Kenyon \\& Hartmann 1995). However, the observations have made clear that the outer part ($\\ga 100$ AU) of the disk that emits millimeter continuum disappears in a timescale between $10^6$ and $10^7$ yr. The standard interpretation assumes that WTTSs are the older counterparts of CTTSs.  A notable feature of the millimeter emission from CTTSs is the evolution of the spectral index from the interstellar value. The frequency dependence of the dust emissivity (or the opacity) $\\kappa_{\\nu}$ is conventionally expressed as a power law, $\\kappa_{\\nu} \\propto \\nu^{\\beta}$, where $\\nu$ is the frequency of radio emission. Beckwith \\& Sargent (1991) argued that the value of the opacity index $\\beta$ evolves from the interstellar value of 2 to $\\la 1$ for CTTS disks (see also Mannings \\& Emerson 1994; Hogerheijde et al. 2003). This change in $\\beta$ is probably due to dust evolution in the disks. However, because different directions of dust evolution (such as in grain size, in chemical composition, and in grain shape) can produce observable changes in the dust opacity, it is not easy to specify the cause of changing $\\beta$ (Beckwith, Henning, \\& Nakagawa 2000). Among the above possible causes, dust growth is an attractive hypothesis in a context of planet formation. The effect of dust growth on the dust opacity was studied by Miyake \\& Nakagawa (1993). Their calculations showed that grain growth to 1 mm would produce a low $\\beta$ value ($\\la 1$), but that it is difficult to explain $\\beta \\la 1$ as long as the maximum size of grains is smaller than $1 \\ \\micron$. In this paper, we stand at the ground that the low $\\beta$ is a sign of dust growth, and we assume that in the disks showing $\\beta \\la 1$ the dust grains have grown to $\\ga 1$ mm.  In summary, radio observations have shown that about half of CTTSs emit detectable millimeter continuum. Images of millimeter continuum show extended disks of $100-500$ AU in size, while their wavelength dependences show the evolution of $\\beta$, suggesting grain growth to $\\ga 1$ mm. Millimeter continuum emissions of CTTSs last for $10^6 - 10^7$ yr.  Dust grains in a gas disk migrate to the central star due to gas drag (see \\S\\ref{sec:rad_mig}). As grains grow larger, their orbital decay time becomes shorter. When the grain size becomes 1 cm, the decay time is only about $10^5$ yr, which is much shorter than the dust disk's life time observed for CTTSs. This dust growth and subsequent dust accretion induce a decline in millimeter emission from the disk. The millimeter emission would disappear in $10^5$ yr, if most of the dust mass had been confined in the largest grains of $\\ga 1$ cm. Thus, we expect that in CTTS disks, there must remain sufficient dust grains less than 1 cm over $10^6$ yr. This expectation is confirmed in \\S\\ref{sec:evo_den} using numerical calculations. At the same time, however, the largest grain size must be larger than about 1 mm, if low $\\beta$ is due to dust growth. Very large bodies (say $\\ga 10$ m) experience little migration, because such bodies almost completely decouple from the gas. However, such large bodies do not contribute to millimeter emission (Miyake \\& Nakagawa 1993). Hence, in this paper we conclude that only a narrow range of the dust size, $[1 \\ {\\rm mm}, 1 \\ {\\rm cm}]$, can fulfill the requirement that the millimeter continuum having $\\beta \\la 1$ lasts over $10^6$ yr.  Our calculations suggest that the millimeter continuum of CTTS disks is mainly emitted from grains of $1 \\ {\\rm mm}- 1 \\ {\\rm cm}$. Thus, the lifetime of observable millimeter continuum, $10^6 -10^7$ yr, constrains the growth time of millimeter sized grains. One possibility is that millimeter sized grains do not grow to centimeter size in $10^6$ yr and they hold the millimeter continuum of $\\beta \\sim 1$. Another possibility is rapid growth of grains in a context of planet formation: grains in CTTS disks have already grown to sizes of $\\ga 10$ m, which is large enough to resist gas drag and to reside in the disk more than $10^6$ yr. We consider these possibilities in \\S\\ref{sec:growth}. ",
        "conclusions": "Motivated by radio observations suggesting grain growth to 1 mm or larger in CTTS disks, we studied the evolution of dust disks whose main components are millimeter sized grains.  1. Gas drag on millimeter sized grains accelerates the dust disk accretion to the star. When grains grow to $\\sim 1$ cm, they cannot reside in the disk for over $10^6$ yr, and as the dust disk accretes, its millimeter continuum emission decreases. The lifetimes of grains are longer for a smaller particle physical density (i.e., fluffy), and in a more massive gas disk.  2. If the observed millimeter continuum of CTTSs comes from millimeter sized grains, a sufficient amount of grains in this size range must exist more than $10^6$yr in their disks. This suggests one of the following possibilities. One is that millimeter sized grains do not grow rapidly, and it takes more than $10^6$ yr to become larger than 1 cm. If the grains' sticking probability is less than $\\sim 0.1$, this condition is satisfied in the outermost region of the disk ($r \\ga 100$ AU). Another possibility is that the grains have already grown to $10$ m or larger and they do not migrate rapidly anymore, and millimeter sized grains are continuously replenished through collisions of the large bodies. If the disk's dust-to-gas ratio is $\\sim 0.1$ and the sticking probability is $C_{\\rm stk} \\sim 0.1$, the observable millimeter continuum holds for more than $10^6$ yr, emitted mainly from grains of $\\sim 1$ mm in the outer disk, and in the inner disk the grains can grow larger than 10 m before falling to the star. The other possibilities are that millimeter sized grains are replenished from a surrounding envelope, or that the small $\\beta$ of the observed millimeter continuum is not due to dust growth. In order to examine these possibilities, we need to perform simulations taking account of both dust migration and dust growth at the same time."
    },
    "0410/astro-ph0410576_arXiv.txt": {
        "abstract": "SPINOR is a new spectro-polarimeter that will serve as a facility instrument for the Dunn Solar Telescope at the National Solar Observatory. This instrument is capable of achromatic operation over a very broad range of wavelengths, from $\\sim$400 up to 1600~nm, allowing for the simultaneous observation of several visible and infrared spectral regions with full Stokes polarimetry. Another key feature of the design is its flexibility to observe virtually any combination of spectral lines, limited only by practical considerations (e.g., the number of detectors available, space on the optical bench, etc). SPINOR is scheduled for commissioning by the end of 2005. In this paper we report on the current status of the project and present actual observations of active regions in the Ca~II infrared triplet and the He~I multiplet at 1083~nm. ",
        "introduction": "\\label{sec:intro}  The new breed of spectro-polarimeters scheduled to begin operation in the next few years make the present a very exciting time for solar physics. Examples of these instruments are SOLIS \\citeP{K98}, Solar-B \\citeP{LES01}, POLIS \\citeP{SBK+03}, Sunrise \\citeP{SCG+03}, and the DLSP \\citeP{SEL+03}. All these instruments, however, are highly specialized and operate under fixed conditions, either with the aim of performing synoptic observations or to optimize for spatial resolution. While these aspects are very important for the future advance of solar physics, we feel that there is also a need for an ``experiment-oriented'' type of instrument. By this we mean an instrument that provides enough flexibility to implement more or less arbitrary optical arrangements at the observer's request. Such an instrument should allow researches to observe any given spectral line (or combinations of lines), either at the disk or off the limb at a diverse range of wavelengths. One should be able to modify parameters like spatial resolution, spectral dispersion, integration times, etc, in order to address a broad range of scientific problems. In some sense, operating an instrument of this kind is like setting up an experiment in a laboratory. Optical components and detector systems are arranged on an optical bench to investigate some specific problem.  For more than a decade, the highly successful Advanced Stokes Polarimeter (ASP, \\citeNP{ELT+92}) has provided these capabilities at {\\it visible} wavelengths. The Spectro-Polarimeter for Infrared and Optical Regions (SPINOR) will replace the ASP as the experiment-oriented spectro-polarimeter on the Dunn Solar Telescope (DST) at the National Solar Observatory (NSO), Sunspot, NM, USA. SPINOR is currently under deveopment by the High Altitude Observatory and the National Solar Observatory. It has been conceived with flexibility as a top priority and intends to fill this gap in the next-generation solar instrumentation. The new instrument will replace and enhance the capabilities of the ASP.  Table~\\ref{table1} below provides some specifications for SPINOR. The most important enhancement over the ASP is the ability to observe simultaneously in the visible {\\it and} the near infrared. Higher spatial resolution than was available over most of the life of the ASP is now possible thanks the the high-order adaptive optics (AO) system \\citeP{R00} developed for the DST. Two different detectors will be available at commissioning time (scheduled for the end of 2005): a PixelVision Pluto 652x488 thinned back illuminated split frame transfer CCD, 60 frames per second readout rate (hereafter Pluto) and a Sarnoff CAM1M100, 1024x512, thinned back illuminated split frame transfer CCD, 100 frames per second readout rate. These will be fully dedicated for observations between 400 and $\\sim$900~nm. The SPINOR optics allows for an even broader range of wavelengths, up to 1.6~$\\mu$m. We plan to take advantage of this in future observations by using an HAO Rockwell infrared camera that is shared with other HAO projects. The modular design of the instrument control will allow for additional detectors to be attached to SPINOR using NSO's Virtual Camera interface.  The present configuration of the instrument, used in the engineering run reported here, does not include the Sarnoff or the Rockwell cameras. For this run we used the ASP control computers and detectors alongside the Pluto camera. With this configuration we were able to demonstrate SPINOR's achromatic spectro-polarimetry up to 1083~nm.  This paper is organized as follows. Section~\\ref{sec:science} discusses the scientific motivations for the development of SPINOR. Section~\\ref{sec:instrument} presents a description of the polarimeter with some details on the new achromatic optics. The polarimetric calibration of the system consisting of the telescope and the polarimeter is explained in section~\\ref{sec:calib}. Some observations taken during an engineering run in June 2004 are presented in section~\\ref{sec:obs}. Finally, section~\\ref{sec:conc} presents the main conclusions of our work and some future perspectives. ",
        "conclusions": "\\label{sec:conc}  This paper presents some initial results from first-light observations with SPINOR, a new facility instrument for the DST. We perceive an urgent demand in the solar community for a new experiment-oriented  (as opposed to specialized) spectro-polarimeter. We expect that SPINOR will fulfill this need and remain at the cutting-edge of solar research at least until the construction of the Advanced Technology Solar Telescope (\\citeNP{KRH+02}). Its broad wavelength coverage will provide a uniquely connected view of photospheric, chromospheric and, to some extent, coronal magnetism. The observations presented in the present paper are intended to emphasize this point and to demonstrate the capabilities of the instrument in the near infrared.  It is important to remark that the final configuration of the instrument will provide significant improvements for infrared observations with respect to the performance reported here. The new achromatic modulator (already deployed) and cameras (larger format and more efficient) will hopefully enhance SPINOR's performance at 1$\\mu$m and extend its range up to 1.6 $\\mu$m.  As with the ASP, researchers from national and foreign institutions will be able to access the new instrument through the usual time allocation competition for the Dunn Telescope, operated by NSO."
    },
    "0410/astro-ph0410606_arXiv.txt": {
        "abstract": "UW CrB (MS1603+2600) is a peculiar short period X-ray binary that exhibits extraordinary optical behaviour. The optical light curve shape of the system changes drastically from night to night, without any changes in overall brightness.  Here we report X-ray observations of UW CrB obtained with {\\it XMM-Newton}. We find evidence for several X-ray bursts confirming a neutron star primary. This considerably strengthens the case that UW CrB is an Accretion Disc Corona (ADC) system located  at a distance of at least 5--7 kpc, (3--5 kpc above the galactic plane). The X-ray and optical monitor (UV+optical) light curves show remarkable shape variation from one observing run to another, which we suggest are due to large scale variations in the accretion disc shape resulting from a warp which periodically obscures the optical and soft X-ray emission. This is also supported by the changes in phase-resolved X-ray spectra. ",
        "introduction": "UW CrB (previously known as MS1603+2600) was discovered by Morris et al. (1990) in the \\textit{Einstein} medium sensitivity X-ray survey.  They identified the optical counterpart with a blue, emission-line only object close to the X-ray position. Subsequent optical photometry revealed a period of 111 minutes. Morris et al. (1990) also noted UW CrB's key distinguishing feature of an optical light curve that changes dramatically from one night to another.  UW CrB is located at a high galactic latitude of 47$^o$, which, together with its low X-ray flux, suggests that either the source is underluminous in X-rays for a low mass X-ray binary (LMXB) or that it is actually located in the galactic halo.  \\begin{figure*} \\includegraphics[scale=0.9]{unfolded.ps} \\caption{The background subtracted unfolded EPIC pn 0.2--10 keV light curves from epoch 1 (top, 2003 Jan 20) and epoch 2 (bottom, 2003, Jan 22). The time resolution is 10 sec. Several X-ray bursts are seen during each observation. } \\end{figure*}  \\begin{figure*} \\includegraphics[scale=0.9]{plot2.ps} \\caption{The phase folded EPIC-pn and OM lightcurves. From top to bottom: 0.2--10.0 keV, 0.2--1.0 keV, 1.0--3.0 keV, 3.0--10.0 keV and OM white light. The left panel is epoch 1 and the right panel epoch 2. The 0.0 phase is arbitrary and taken from the main sharp minimum of epoch 1 OM light curve. } \\end{figure*}    Hakala et al. (1998) published the results of further optical photometry together with the ROSAT observations.  They confirmed the extreme variability seen in the optical light curve pulse shapes. Even if there are dramatic changes in the optical light curve shape, the overall optical flux level does {\\bf not} vary significantly.  ROSAT data showed that the soft X-ray spectrum of UW CrB can be modelled with a single blackbody component, with a temperature of kT$\\sim$0.24 keV. As expected at this galactic latitude, the spectra showed no significant interstellar absorption.  Earlier models for UW CrB have invoked both LMXB and magnetic CV (cataclysmic variable) explanations (Morris et al. 1990, Ergma \\& Vilhu 1993). Hakala et al.  (1998) compared the ROSAT spectra of UW CrB with soft X-ray spectra of various types of interacting binaries. Their conclusion was that the only class where the spectral fits seemed to match those of the source was the (black hole) soft X-ray transients in quiescence. This could also have explained the relatively low \\textit{F$_{X}$/F$_{opt}$} ratio ($\\sim$15) observed.  In addition, Hakala et al.  (1998) report a negative result from their search for circular polarisation.  Mukai et al. (2001) presented ASCA observations of UW CrB. They claim to see a single type I X-ray burst, which would rule out both black hole and white dwarf explanations for the accreting component in the system.  They proposed instead that UW CrB is a short period X-ray dipper, similar to X1916-05. They also note that the system would probably in this case have to reside in the galactic halo, and could have formed in the globular cluster Palomar 14, which is located 11$^o$ from UW CrB and at a distance of 73.8 kpc.  Recently Jonker et al. (2003) report the results from Chandra observations of the source. They conclude that depending on the nature of the single burst event detected by ASCA (Mukai et al. 2001), UW CrB could be either a ``nearby'' quiescent transient or an ADC source at 11--24 kpc distance. They consider the dipper interpretation of Mukai et al. (2001) to be unlikely, mainly on the basis of optical brightness (eg. the system would be intrinsically much brighter in the optical than other, longer period LMXB's that have a larger disc). Quite recently, Muhli et al. (2004) reported that the optical light curves of UW CrB show optical bursts. This has now also been confirmed by Hynes et al. (2004).  In this paper we will present new results of UW CrB taken using {\\it XMM-Newton} and discuss the nature of the source in light of these extensive new data. ",
        "conclusions": "Using our {\\it XMM-Newton} data, we find evidence for evolving vertical disc structure (either warped or an asymmetric flared disc). This is evident from the strong X-ray modulation over the orbital period, that has changed amplitude and shape over the 2 day period separating the two epochs. The nature of X-ray and UV-optical modulation also suggests that UW CrB is most likely an ADC source.  The detection of several X-ray bursts confirms that the compact object in UW CrB is a neutron star. However, as we probably only see scattered X-rays from the ADC, the bursts luminosities cannot be used as a direct tool for distance estimates. Assuming that we would only see about 1 \\% of the total burst flux  then we can estimate that the minimum distance to the source would be $\\sim$ 6 kpc. This would place the source at about 4 kpc above the galactic plane."
    },
    "0410/astro-ph0410430_arXiv.txt": {
        "abstract": "We present kinematics of a sample of $398$ DA white dwarfs from the SPY project (ESO SN Ia Progenitor surveY) and discuss kinematic criteria for a distinction of thin disk, thick disk, and halo populations. This is the largest homogeneous sample of white dwarfs for which 3D space motions have been determined. Radial velocities and spectroscopic distances obtained by the SPY project are combined with our measurements of proper motions to derive 3D space motions. Galactic orbits and further kinematic parameters are computed. Our kinematic criteria for assigning population membership are deduced from a sample of F and G stars taken from the literature for which chemical criteria can be used to distinguish between thin disk, thick disk and halo. Our kinematic population classification scheme is based on the position in the $U$\\/-$V$-velocity diagram, the position in the $J_{\\mathrm{z}}$-eccentricity diagram and the Galactic orbit. We combine this with age estimates and find seven halo and $23$ thick disk white dwarfs. ",
        "introduction": "White dwarfs are presumably valuable tools in studies of old populations such as the halo and the thick disk, since the fraction of old stars among white dwarfs is higher than among main-sequence stars. An open issue is the fraction of white dwarfs in the thick disk and halo populations and their percentage of the total mass of the Galaxy. Studies of their kinematics can help to determine the fraction of the total mass of our Galaxy contained in the form of thick disk and halo white dwarfs which allows the role of white dwarfs in the dark matter problem to be studied.    The common problem of kinematic studies of white dwarfs is the lack of radial velocity measurements. Especially deviating conclusions derived from the white dwarfs of the Oppenheimer et al. (2001) sample demonstrate that different assumptions on the values of $v_{\\rm rad}$ can produce different fractions of halo and thick disk stars and thus have a strong impact on the determination of the white dwarf halo density. Therefore a sample of white dwarfs with known radial velocity measurements is needed in order to obtain the full 3D kinematic information. Pauli et al. (2003) presented a complete 3D kinematical study of 107 white dwarfs. Here we extend this investigation to about 400 stars. ",
        "conclusions": "We have applied the classification scheme developed in Pauli et al. (2003) to kinematically analyze a sample of about $400$ DA white dwarfs from the SPY project. Combining the three kinematic criteria position in the $U$\\/-$V$-diagram, position in the $J_Z$-$e$-diagram and Galactic orbit with age estimates we have found seven halo and $23$ thick disk members.  The velocity dispersions that for the thin disk white dwarfs are well compatible with the ones of Soubiran et al. (2003). The same is the case for the asymmetric drift and the velocity dispersions of the thick disk. For the halo white dwarfs $\\sigma_{\\rm U}$ and $\\sigma_{\\rm V}$ are similar to the values derived by Chiba \\& Beers (2000) for halo stars while our $\\sigma_{\\rm W}$ is much smaller, probably due to low number statistics. Hence the kinematic parameters of the white dwarfs of the three different populations do not differ much from main-sequence samples."
    },
    "0410/astro-ph0410356_arXiv.txt": {
        "abstract": "We compute the emission of gravitational radiation from the merging of a close  white dwarf binary system. This  is done for a  wide range of masses  and compositions  of the  white dwarfs,  ranging  from mergers involving two  He white dwarfs, mergers  in which two  CO white dwarfs coalesce to  mergers in which a  massive ONe white  dwarf is involved. In doing so we follow the  evolution of binary system using a Smoothed Particle Hydrodynamics  code.  Even though the  coalescence process of the white  dwarfs involves  considerable masses, moving  at relatively high  velocities with  a high  degree of  asymmetry we  find  that the signature  of  the merger  is  not very  strong.   In  fact, the  most prominent feature  of the  coalescence is that  in a  relatively small time  scale (of  the order  of the  period of  the last  stable orbit, typically  a  few minutes)  the  sources  stop emitting  gravitational waves.  We also discuss  the possible implications of our calculations for the  detection of the  coalescence within the framework  of future space-borne interferometers like LISA. ",
        "introduction": "Gravitational waves are a direct  consequence of the General Theory of Relativity.  Many  efforts have been done  so far to  detect them, but due to the intrinsic experimental difficulties involved in the process of detection and in the data  analysis no definite result has been yet obtained.  Supernova  core collapse, binary  systems involving compact objects, like double black holes, double neutron stars or double white dwarfs are,  amongst others, promising sources  of gravitational waves --- see  Schutz (1999)  for  a comprehensive  review  on the  subject. Moreover,  with the advent  of the  current generation  of terrestrial gravitational  wave detectors,  like LIGO  (Abramovici et  al.  1992), VIRGO (Acernese et  al.  2004), GEO600 (Willke et  al.  2004), or TAMA (Takahashi  et al.   2004),  and of  space-borne interferometers  like LISA\\footnote{\\tt  http://lisa.jpl.nasa.gov}  (Bender  et  al.   1998, 2000), gravitational  wave astronomy will probably be  soon a tangible reality.  As  already   mentioned,  one  of   the  most  promising   sources  of gravitational waves  are galactic  binary systems containing  at least one  compact  object.   Galactic   binaries,  such  as  neutron  stars binaries,  cataclysmic binaries  or  close white  dwarf binaries,  are guaranteed  sources for LISA  (Mironowsky 1965;  Evans, Iben  \\& Smarr 1987), provided that the  sources are at sufficiently close distances. In fact,  emission from galactic  close white dwarf binary  systems is expected to  be the dominant  contribution to the background  noise in the low frequency region, which ranges from $\\sim 10^{-3}$ up to $\\sim 10^{-2}$~Hz  (Bender et  al.  1998).   Additionally, from  very simple considerations about the initial mass  funtion, it is easy to see that galactic close white dwarf binaries must be quite common (Hils, Bender \\&  Webbink   1990)  and,  consequently,  if  the   amplitude  of  the gravitational waves  is large enough  we should be able  to eventually detect them  during the operation  of LISA.  Moreover, the  merging of two white  dwarfs by emission  of gravitational radiation will  be the final  destiny of  a good  fraction of  this type  of  binary systems. Since during  the merging process  a sizeable amount  of gravitational waves  is  expected  to be  produced  (Guerrero  et  al. 2004)  it  is important  to  characterize  which  would be  the  gravitational  wave emission of such  process and to assess the  feasibility of dectecting them.  The process  of formation of  close white dwarf binaries  involves two mass  transfer episodes  of  the  progenitor stars  when  each of  the components  of  the  binary  system  evolves off  the  main  sequence. Depending on when  during the lives of the  binary components the mass transfer  episodes  occur  the  components  may  have  different  core compositions.  In particular,  we may have He-He systems  with a total mass $M_{\\rm tot}  \\la 0.75\\, M_{\\sun}$, He-CO for  those systems with masss  within the  range $0.75\\,  M_{\\sun}  \\la M_{\\rm  tot} \\la  1.45 M_{\\sun}$,  CO-CO for  masses  larger than  $M_{\\rm  tot} \\sim  1.45\\, M_{\\sun}$  and even He-ONe  or CO-ONe  systems when  one of  the white dwarfs  is a  massive one.   Although the  astrophysical  scenarios in which a merger of two white  dwarfs in a close binary system can occur and their  relative frequencies have been relatively  well studied --- see,  for instance,  Yungelson et  al.   (1994), and  Nelemans et  al. (2001a,  2001b), and  references therein  --- the  process  of merging itself has received little attention until very recently.  Indeed, one of the  probable reasons  for this lack  of theoretical models  is the heavy  computational   demand  involved   in  the  simulation   of  an intrinsically   three-dimensional  phenomenon.    However,   in  sharp contrast, the  coalescence of two  neutron stars has  been extensively studied  --- see,  for instance,  Rosswog et  al.  (2000),  Rosswog \\& Davies (2002),  and Rosswog  \\& Liebend\\\"orfer (2003),  and references therein, for some of the most recent works on this subject.  In a  recent paper (Guerrero et  al.  2004) we  thoroughly studied the merging of white  dwarf binary systems for a wide  range of masses and compositions.   In doing  this  we used  an  up-to-date {\\sl  Smoothed Particle Hydrodynamics}  code. This method was first  proposed by Lucy (1977) and, independently,  by Gingold  \\& Monaghan (1977).   The fact that  the method is  totally Lagrangian  and does  not require  a grid makes   it   specially   suitable   for  studying   an   intrinsically three-dimensional problem like the coalescence of two white dwarfs. In the present paper we will not discuss the simulations presented there. The interested  reader can find them  in mpeg format  at the following URL: {\\tt ftp://ftp.ieec.fcr.es/pub/astrofisica/SPH}. Instead, in this paper we  compute the gravitational  wave signature expected  from the merging of two  white dwarfs as obtained from  those SPH calculations. The paper  is organized as follows.   In \\S 2 we  briefly describe the main ingredients  of our SPH  code, whereas in  \\S 3 we  summarize the physics of  the emission of gravitational  waves.  In \\S  4 we discuss the calibration and consistency checks we have done in order to assess the reliability  of our results, whereas in  \\S 5 and \\S  6 we present our results.  Finally,  in \\S 7 we summarize our  results and draw our conclusions. ",
        "conclusions": "We have computed the emission  of gravitational waves of merging white dwarf binaries,  for a  wide range of  masses and compositions  of the components of the  binary system. For that purpose we  have used a SPH code  which  allowed  us  to  follow the  temporal  evolution  of  the coalescing  white dwarfs.   We  have shown  that  the most  noticeable feature  of  the emitted  signal  is  a  sudden disappearance  of  the gravitational strains.  By contrast  the chirping phase will be easily detectable by  future space-borne interferometers like  LISA. In fact, it can be said that the  most relevant signature of the merger will be the  absence of  any signature  and  the sudden  disappearance of  the source. Since the frequency coverage of LISA will range from $10^{-1}$ to $10^{-4}$~Hz,  the detection of  chirping close white  dwarf binary systems  is  guaranteed  (Farmer  \\&  Phinney 2003;  Nelemans  et  al. 2001a).   Moreover, at  frequencies of  $\\sim 3\\cdot  10^{-3}$~Hz most close white dwarf binary  systems will be spectrally resolved (Cornish \\& Larson  2003). Typically, LISA will  be able to  detect $\\sim 3000$ binaries at  these frequencies (Seto 2002).  Additonally,  it has been recently shown  (Cooray, Farmer \\&  Seto 2004) that the  typical error box  of LISA  for  these  kind of  systems  will be  $\\delta\\Omega\\sim 5$~deg$^2$, and that, given that  many of the detected sources will be eclipsing binaries  with a period  equal to that of  the gravitational waves, optical follow-up campaigns  will allow us to further constrain the location of the sources.  Hence, by combining optical observations and  gravitational  wave  data,  and  taking  into  account  that  the gravitational wave  signal for such mergers suddenly  vanishes --- or, equivalently, that  during the integration  period the signal-to-noise ratio stops  growing --- we  should be able  to gain insight  into the physics of merging and to obtain precise information of the properties of the progenitor  systems. Moreover, it should be  taken into account that for a typical merger ---  namely, the $0.6 + 0.8\\, M_{\\sun}$ case --- the  volume  accesible  to  LISA is  $V_{\\rm  LISA}\\sim  1.2\\times 10^{14}$~pc$^3$.   Since the volume  of the  Galaxy is  $V\\sim 3\\times 10^{11}$~pc$^3$, LISA would be able to detect all the mergers occuring in our Galaxy during its  operation period.  However, the typical rate of  white  dwarf  mergers   is  $r\\sim  8.3  \\times 10^{-3}$~yr$^{-1}$ (Nelemans  2003) and,  hence, although  there is  an uncertainty  of a factor of 5 in the rate of white dwarf mergers, the expected detection rate is consequently small.  \\vspace{0.5 cm}  \\noindent  {\\sl  Acknowledgements.}   This  work  has  been  partially supported   by   the   MCYT  grants   AYA2002--4094--C03--01/02,   and AE-ESP2002-10451-E,  by the  European  Union FEDER  funds  and by  the CIRIT.  We  would also  like to acknowledge  the invaluable  advise of J.M.   Ib\\'a\\~nez   who  largely  contributed   through  his  support, suggestions and comments to improve  the manuscript. We also thank our anonymous referee for valuable comments and criticisms."
    },
    "0410/astro-ph0410160_arXiv.txt": {
        "abstract": "We conducted a daily monitoring campaign on Mrk~421 in 2003 and 2004 with the Whipple 10 m telescope and the large-area instruments aboard {\\em RXTE}, simultaneously covering TeV and X-ray energies. Supporting observations at optical and radio wavelengths were also frequently carried out. Mrk~421 was observed over a wide range of fluxes (with a dynamic range of $\\sim$30 both at TeV and X-ray energies). The source was relatively quiet in 2003 but became unusually active in 2004, with flares reaching peak fluxes of $\\sim$80 mCrab in X-rays and $>$3 Crab at TeV energies! We will describe the multiwavelength campaign and present some preliminary results. We will also discuss the implications of the results on the proposed emission models for TeV blazars. ",
        "introduction": "Over the past decade, one of the most exciting advances in high energy astrophysics has been the detection of sources at TeV energies with ground-based $\\gamma$-ray facilities (see Weekes 2003 for a recent review). Among the sources detected, blazars are arguably the most intriguing. The emission from a blazar is generally thought to be dominated by radiation from a relativistic jet that is directed roughly along the line of sight (Urry \\& Padovani 1995). The spectral energy distribution (SED) of TeV blazars invariably shows two characteristic ``humps'' in the $\\nu F_{\\nu}$ representation, with one located at X-ray energies and the other at TeV energies (Fossati et al. 1998). There is a general correlation between the two SED peaks. TeV blazars are also known to undergo flaring episodes, both at X-ray and TeV energies. The flares have been observed over a wide range of timescales, from months down to minutes. The observed X-ray flaring hierarchy in Mrk 421 seems to imply a scale-invariant physical origin of the flares (Cui 2004).  A popular class of models associates the X-ray emission from a TeV blazar with synchrotron radiation from highly relativistic electrons in the jet and the TeV emission with synchrotron self-Compton (SSC) processes (e.g., Marscher \\& Gear 1985; Maraschi et al. 1992). The SSC models can, therefore, naturally account for the observed X-ray--TeV correlation. Moreover, they have also enjoyed some success in re-producing the observed SEDs. However, the models face challenges, such as the presence of ``orphan TeV flares'' (in 1ES 1959+650, Krawczynski et al. 2004; and in Mrk 421, see Fig.~1 in this work). Alternatively, the jet could be energetically dominated by the magnetic field, and it is the synchrotron radiation from highly relativistic protons that might be responsible for the observed TeV gamma rays (Aharonian 2000; M\\\"ucke et al. 2003). In such a scenario, the X-rays are thought to be mainly due to synchrotron radiation from co-accelerated electrons in the jet. Although the proton-synchrotron models can also describe the observed SEDs and accomodate the X-ray--TeV correlation, they are being challenged by variability timescales below $\\sim$15 min in TeV blazars. Very recently, it has been proposed that the flares in blazars might be associated with magnetic reconnection events in a magnetically dominated jet (Lyutikov 2003), perhaps similar to solar flares in this regard. The model might offer a natural explanation for the hierarchical flaring phenomenon in Mrk 421, in analogy to solar flares. However, it has not been applied to the data in any quantitative manner.  To make further progress on distinguishing emission models proposed for TeV blazars, we believe that simultaneous or contemporaneous data are critically needed over a wide range of fluxes, especially in the crucial X-ray and TeV bands, for quantifying the SED and its variability of a source and for allowing investigations of such important issues as variability timescales, cross-band correlation, spectral variability and hysteresis, and so on. Such data are severely lacking at present, despite intense observational efforts over the years. ",
        "conclusions": "A substantial amount multiwavelength data have been collected on Mrk 421 through our monitoring campaign. The data have allowed us to carry out investigations on variability timescales, cross-band correlation, SED and its flux dependence, and spectral variability and hysteresis across flares. Some preliminary results have been presented in this contribution. They have already begun to shed light on emission mechanisms in TeV blazars. The study of most rapid flares is hindered by the snapshot observing strategy adopted. It is, therefore, highly desirable to take a complementary approach in the future by observing the source continously for sufficiently long time."
    },
    "0410/astro-ph0410210_arXiv.txt": {
        "abstract": "{We present the result of photo-ionizing modelling of the three planetary nebulae (PNe) A~20, A~15 and MeWe~1-3. All three objects are roughly roundish, highly excited and have a high galactic $z$. The PNe displayed low densities in the shell, but relatively dense halos. A low metallicity and a relative high electron temperature were found. Comparisons with radio observations confirmed the obtained properties. The objects are very likely originating from thick disk stellar progenitors. The distances found investigating the PNe shells are somewhat lower than those derived spectroscopically with the central stars in the past. ",
        "introduction": "Although the interacting wind model for the formation of planetary nebulae (PNe) is accepted since a long time (Kwok et al.~\\cite{kwok}), the exact mechanism is not clearly understood up to now. Several numerical simulations have been carried out especially during the last decade (e.g. Frank \\& Mellema~(\\cite{frank2}), Dwarkadas \\& Balick~(\\cite{dwar}) or Sch\\\"onberner \\& Steffen~(\\cite{schon}). The last one carried out an one-dimensional hydrodynamic simulation based on the evolutionary tracks of Bl\\\"ocker~(\\cite{blocker}) and the simulations of the circumstellar shell around Asymptotic Giant Branch (AGB) star (Steffen et al.~\\cite{steffen}). \\\\ The results of their simulation showed that a shock-bounded ionized main PN shell moving supersonically into the ABG-material compressing the inner parts of the matter into a dense but thin shell. The unaffected AGB-material becomes ionized as well and forms a rapidly expanding halo. Due to the drop of luminosity of the central stars of the PNe (CSPN) towards the white dwarf regime the outer part of the PN shell recombines quickly forming an second inner halo. The disadvantage of this simulation is that it just considers one dimension. Thus more complex interactions of the winds like Rayleigh-Taylor instabilities, quickly building a clumpy environment as we see it in HST images, are not taken into account. This causes that the filling factor ($\\rm\\varepsilon$) was set in fact to unity. Comparing densities obtain from forbidden line ratios and the one obtained from radio data or similar methods clearly show a filling factor far below unity (e.g. Mallik \\& Peimbert~\\cite{fill2} and Boffi \\& Stanghellini~\\cite{fill3}). This changes the dramatically the optical thickness for ionizing UV photons and thus a recombining PN-shell is doubtful (Armsdorfer et al. \\cite{birgit}).  \\noindent The three PNe A~20, A~15  (Abell~\\cite{Abell}) and \\mbox{MeWe~1-3} (Melmer \\& Weinberger~\\cite{MeWe}) were selected as they are roughly roundish objects. The central stars of the PNe (CSPNe) are well studied spectroscopically (Rauch et al.~\\cite{rauch1}, McCarthy et al.~\\cite{McC}, Saurer et al.~\\cite{saurer}). This allows us to fix the input of radiation for the nebula. We modelled the physical properties of PNe-shells of three objects and compare the results with the theory.  \\begin{figure*}[ht] \\centerline{\\resizebox{18cm}{!}{\\includegraphics{1607F1A.eps}\\phantom{lkfag}\\includegraphics{1607F1B.eps}}}  \\caption{The observations and the model for several lines of A 20. \\noindent{\\bf{Left:}}\\ Observation with NTT/EMMI (slit N--S);~\\noindent{\\bf{Right:}}\\ Observation with \\mbox{ESO~3.6~m/EFOSC} (slit E--W); especially in the [OIII] lines the radial symmetry of the object is well pronounced. Also the fragmentation in three parts, halo, PN shell and central cavity, is displayed well. The [NII] lines are in an upper limit and thus confirm clearly the low metallicity found here (see Tab. \\ref{chem214}).} \\label{mod214a} \\end{figure*} ",
        "conclusions": "The metallicities of all three PNe investigated are lower by a factor of 10 to 30 compared to the solar abundances, only helium is little bit more present than in the Sun. Soker~(\\cite{soker}) already reports an under-abundance of metals in round PNe.\\\\ Looking at the galactic coordinates and the distances of the planetary nebulae one can easily verify that the height above the galactic plane ($z$) of these nebulae is comparatively large (Fig.~\\ref{dist}).\\\\ The nebulae investigated are located at a $z$ of $320$~pc, $-1127$~pc and $535$~pc for A~20, A~15 and  MeWe~1-3 respectively. To compare these values with other one the PNe catalogued by Acker et al.~(\\cite{ESO}) were taken as a representative sample for all PNe. Half of these planetary nebulae are located below a $z$ of 162~pc. Their average $z$ amounts to 288~pc. The large $z$ of the three objects investigated indicates that they are so called thick disk objects and therefore their progenitor was an underabundant object as well.  \\begin{figure}[ht] \\centerline{\\resizebox{8.8cm}{!}{\\includegraphics{1607F10.eps}}}  \\caption{Here the galactic height $z$ versus the radius from the galactic center is displayed. The radius of the sun was assumed to be 8 kpc. The + shows the PNe catalogued by Acker et al.~(\\cite{ESO}) and the $\\diamondsuit$ the PNe A 20, A 15 and MeWe 1-3.}  \\label{dist} \\end{figure}  The high electron temperature of the nebulae is also linked to the low metallicity (Stasinska~(\\cite{stasi})). As Frank \\& Mellema~(\\cite{frank}) showed one important cooling mechanisms in ionized gas is the radiation emitted by collisionally excited lines. The most important role plays the [NII] lines in low excited nebulae and the [OIII] lines in high excited one. Since no [NII]- and less intensive [OIII] lines, compared to common PNe, are detected a higher temperature is feasible.  As mentioned above a filling  factor ($\\varepsilon$) of 0.1 was adopted. Since radio data are available for A~20 and A~15 (Condon \\& Kaplan~\\cite{con}) it was possible to check this assumption. Adopting a optically thin wind in the radio regime and therefore $\\rm I_{\\nu}\\propto \\nu^{-0.1}$ allows us to compute the radio flux at 5 GHz from the 1.4~GHz data and thus the total ionized mass according Cahn et al.~(\\cite{Dist2}). $$ \\rm M_{ion.}=\\sqrt{2.266\\times 10^{-21}\\cdot D^5\\cdot\\theta^3\\cdot F_{5\\,GHz}} $$ where $D$ is the distance in pc, $\\theta$ the angular radius in arcseconds and $\\rm F_{5\\,GHz}$ the radio flux at 5 GHz in Jansky. We compared this mass with the one of our model, computed by $$ \\rm M=4\\pi\\mu\\int_{0}^{+\\infty}n_{hp}(r)r^{2} dr $$ $\\mu$ is the mean atomic weight and $\\rm n_{hp}(r)$ the heavy particle density as function of the radius. This results in $\\varepsilon$. Although the error of the obtained filling factor is quiet high it is clearly smaller than 0.3. Using the sample and the formula of Mallik \\& Peimbert~(\\cite{fill2}) we obtain a value even below 0.05. The formula of Boffi \\& Stanghellini~(\\cite{fill3}) give a value of 0.096. But especially the sample shows no PN with a radius above 0.2~pc and an filling factor above 0.1. Tests with CLOUDY show that the nebulae are always optically thin for the UV radiation up to $\\varepsilon = 0.5$.\\\\ Further calculations showed that even stars with luminosities down to 160~$L_\\odot$ are able to ionize most of the hydrogen and thus a recombination of the shell seems to be impossible. This is in clear contradiction to the predictions of the 1D hydro simulations (e.g. Sch\\\"onberner \\& Steffen \\cite{schon}).  While the luminosities of the CSPN found from the investigation of the PN shells here are about the same as those found {with the NLTE models of the CSPN spectroscopy}. Our temperatures are - still within the error bars - slightly higher than those from the CSPN spectra. As shown by Lechner \\& Kimeswenger (\\cite{lech}) and Emprechtinger et al. (\\cite{empr}) these parameters are very sensitive to the line ratios of the highly ionized species.  We thus suggest this effect to be real. The distances $D_{CSPN}$ have to be significantly lower than those found by other investigations in the past."
    },
    "0410/astro-ph0410026_arXiv.txt": {
        "abstract": "{  We use numerical N-body simulations in order to test whether the kinematics of stars and brown dwarfs at birth depend on mass. In particular we examine how initial variations in velocity dispersion can affect the spatial distribution of stellar and substellar objects in clusters. We use 'toy' N-body models of a Pleiades-like cluster in which brown dwarfs have their own velocity dispersion $\\sigma_{V_{\\text BD}}$ which is $k$ times larger than the stellar one.  We find that in order to match the broad agreement between the brown dwarf fraction in the field and in the Pleiades, the velocity dispersion of brown dwarfs at birth has to be less than twice the stellar velocity dispersion, i.e. cannot exceed a few km/s in the Pleiades cluster. In order to discern more subtle differences between the kinematics of \\bds and stars at birth, our simulations show that we need to look at clusters that are much less dynamically evolved than the Pleiades. One might especially seek evidence of high velocity brown dwarfs at birth by examining spatial distribution of stars and brown darfs in clusters that are about a crossing timescale old. ",
        "introduction": "So far, most studies of the star formation process have dealt with the formation of stellar systems but with the recent discovery of brown dwarfs in 1995 (Nakajima et al. 1995; Rebolo et al. 1995) new perspectives have opened regarding the formation of condensed objects in molecular clouds. Today more than a hundred brown dwarfs (BDs) in various environments are known. However, their mode of formation is still controversial and the theoretical framework describing the stellar and substellar formation process(es) is not completely satisfactory. Regions of high density ($n(H_{2})\\sim 10^{7}$ cm$^{-3}$) in molecular clouds are needed to form proto-\\bds but then for their mass to remain substellar their reservoir of gas has to be small or the accretion not very efficient.  Two main competing scenarios have been proposed so far to account for the formation of substellar objects. One assumes that brown dwarfs form like solar-mass stars by gravitational collapse of small, dense molecular cloud core and subsequent accretion. The supporting argument is that in the opacity limited regime the Jeans mass can be as low as a few Jupiter masses (Low \\& Lynden-Bell 1976). The alternative view assumes that brown dwarfs are ejected ``stellar embryos'' as proposed by Reipurth \\& Clarke (2001). In this scenario, molecular cloud cores fragment to form unstable protostellar multiple systems which decay dynamically. The lowest mass fragments are ejected from their birth place and deprived of surrounding gas to accrete remain substellar objects.  The brown dwarf properties predicted by these two different formation scenarios may in principle be quite different. In the former case, both stars and brown dwarfs form predominantly as single or binary ($N=2$) systems; in this case there are no obvious reasons why properties such as the binary fraction or kinematics should depend on mass.  In the latter scenario, by contrast, the dominant formation mechanism (again for both stars and brown dwarfs) is in small $N$ ($>2$) clusters, and the gravitational interplay that precedes the break up of the system into stable entities implies a potentially strong mass dependence for resulting properties like the binary fraction and kinematics. In particular it has been suggested (Reipurth and Clarke 2001) that low mass objects (e.g. brown dwarfs) ejected from such clusters would have a higher velocity dispersion than higher mass objects.  Reipurth and Clarke's initial suggestion - that brown dwarfs in star forming regions may have a detectably higher velocity distribution from stars - has {\\it not} been borne out by radial velocity studies (Joergens \\& Guenther 2001). Meanwhile, successive simulations have modified the predictions of small $N$ clusters models. Delgado et al. (2003) and Sterzik and Durisen (2003) have emphasised that, in their simulations, the main difference in velocity dispersion is between single stars and binaries, and that brown dwarfs attain rather larger velocities - with respect to their parent cores - because they are more likely to be ejected as single objects. This dependence of ejection speed on binarity may readily be understood, since one binary is typically formed in each cluster in these simulations: this binary is able to eject the remaining stars from the cluster by sling-shot gravitational encounters, whilst itself remaining close to the center of mass of the natal cluster. In the turbulent fragmentation calculations of Bate, Bonnell and Bromm (2003) and Delgado, Clarke and Bate (2004), by contrast, more than one binary is formed per cluster and so binaries are able to eject each other from the natal cluster. Consequently, in these simulations, the kinematics of the resulting objects do not depend strongly on either mass {\\it or} binarity.  Evidently, the relative kinematics of stars and brown dwarfs and of single stars and binaries can shed some light on the conditions in star forming cores and could ultimately answer the question of whether stars (and brown dwarfs) are formed as isolated single and binary systems, as small $N$ aggregates containing typically one binary or as aggregates containing more than one binary. [Note that this question is not easy to answer by direct observations, since the timescale for the break up of putative small clusters implies that this process occurs in the deeply embedded phase. However, high resolution imaging of the driving sources of Herbig Haro objects by Reipurth (2000) suggests that the multiplicity of stars in deeply embedded regions is indeed high].  Direct observations of the kinematics of young stars and brown dwarfs is unlikely to be fruitful however. The differences in velocity dispersion predicted by theoretical models are small (of the order of a km/s). When one bears in mind that these velocities are measured with respect to star forming cores, which are in themselves in relative motion at $\\sim1$ km/s, it is unsurprising that the study of Joergens and Guenther (2001) - involving small numbers of objects, with velocity resolution of $\\sim0.2$ km/s and rather small dynamic range in mass - did not detect any differences.  In this paper, we propose another approach that could potentially detect any mass dependence of the kinematics of stars (and brown dwarfs) at birth. Here we examine the statistical consequences of such an effect on the spatial distribution of stars and brown dwarfs in clusters. This approach has the advantage that one can work with large samples of stars and brown dwarfs, whose positions and masses are known with high accuracy. On the other hand, we cannot predict how initial variations in velocity dispersion affect the spatial distributions of stars and brown dwarfs in a cluster at a given age without further, N-body, modeling.  This is partly because two body relaxation leads to mass segregation in older clusters, even in the absence of a mass dependent initial velocity dispersion.  Our purpose in this paper, therefore, is to use `toy' N-body models (in which brown dwarfs are introduced with a velocity dispersion that is a variable multiple of the stellar velocity dispersion) in order to establish under what circumstances could one detect a higher velocity dispersion at birth for low mass objects. We stress that this toy models for the kinematics is not supposed to correspond to the outcome of any particular numerical star formation model but is designed to provide a ready parameterization of the problem. We also underline that in no models are any sudden discontinuities in kinematic properties expected at the hydrogen burning mass limit.  We use the Pleiades as the testbed for our calculations. This is because the brown dwarf population of the Pleiades has been the subject of intensive scrutiny in recent years (Moraux et al. 2003, Dobbie et al. 2002, Pinfield et al. 2000, Zapatero-Osorio et al. 1999, Bouvier et al. 1998) so that the present day mass function in this cluster is reasonably well constrained. We shall proceed by first placing an upper limit on the initial velocity dispersion of brown dwarfs in the Pleiades, based on the broad similarity between the normalization of stars to brown dwarfs in the Pleiades and in the field, which limits the number of brown dwarfs that can have left the cluster to date. We shall then explore whether the radial distribution of brown dwarfs in the cluster can place meaningful limits on their initial velocity distribution. ",
        "conclusions": "We have shown that the observed similarity between the brown dwarf to star ratio in the Pleiades, in star forming regions and in the field implies that the velocity dispersion of brown dwarfs at birth cannot exceed $2\\times$ the stellar velocity dispersion in the Pleiades. This imposes an upper limit on the brown dwarf velocity dispersion at birth of a few km/s. Thus either brown dwarfs are not dynamically ejected from their natal cores or else the ejection velocities with respect to their natal cores is low. Such a velocity dispersion limit would rule out the common incidence of very hard encounters (i.e. at $<$ 4 A.U., Delgado et al. 2004) in the natal cores, but would be consistent with the results of current hydrodynamical simulations (Delgado et al. 2003, Bate et al. 2003). We stress that the velocities quoted here are limits on the 3D velocity dispersion for a Maxwellian distribution and so do not necessarily rule out a tail of objects with significantly higher velocities (such, for example, as the (not substellar) object PV Ceph, for which a velocity of $20$ km/s has recently been claimed (Goodman \\& Arce 2003).  We have shown that the radial density profile of brown dwarfs and stars in the Pleiades provides {\\it no} information about the velocity dispersion of brown dwarfs at birth. This is because any brown dwarfs with velocities greater than the escape velocity of the cluster would have long ago escaped the cluster and thus the residual brown dwarf population has a velocity distribution - and hence spatial distribution - that is very similar to the stars.  We therefore conclude that if we seek {\\it positive} evidence for a modestly higher velocity distribution of brown dwarfs at birth, we need to look at clusters that are significantly different from the Pleiades. There are two possibilities here. Firstly, the upper limits we have derived could still imply a significant preferential loss of brown dwarfs from loosely bound regions.  For example, Preibisch et al. (2003) suggested that this could be the cause of the small fraction of \\bds observed in IC348 for which the escape velocity is only $\\sim 0.8$ km/s (Herbig 1998). Kroupa \\& Bouvier (2003) found also that the ejection mechanism could explain a part of the deficit of substellar objects in the Taurus association. Secondly, one might seek evidence of high velocity brown dwarfs at birth by instead examining clusters that are about a crossing timescale old. In this case, the brown dwarfs would be expected to have a more extended distribution than the stars, which could {\\it not}, at such a young age, be confused with the signature of two body relaxation. Even if the brown dwarf velocity dispersion exceeds that of stars by only 50\\%, the spatial distributions of the two populations are quite different at this age (contrast dashed lines in Fig.~\\ref{cumul_cut10} with solid line in Fig.~\\ref{bd_star10}) Thus brown dwarf surveys in clusters of various ages and degrees of richness will afford further opportunities to constrain the velocities of brown dwarfs at birth."
    },
    "0410/hep-ph0410349_arXiv.txt": {
        "abstract": "We obtain the equations of motion for cosmic strings in extensions of the 3+1 FRW model with extra dimensions.  From these we derive a generalisation of the Velocity-dependent One-Scale (VOS) model for cosmic string network evolution which we apply, first, to a higher-dimensional isotropic $D+1$ FRW model and, second, to a $3+1$ FRW model with static flat extra dimensions. In the former case the string network does not achieve a scaling regime because of the diminishing rate of string intersections ($D>3$), but this can be avoided in the latter case by considering compact, small extra dimensions, for which there is a reduced but still appreciable string intercommuting probability. We note that the velocity components lying in the three expanding dimensions are Hubble-damped, whereas those in the static extra dimensions are only very weakly damped. This leads to the pathological possibility, in principle, that string motion in the three infinite dimensions can come to a halt preventing the strings from intersecting, with the result that scaling is not achieved and the strings irreversibly dominate the early universe. We note criteria by which this can be avoided, notably if the spatial structure of the network becomes essentially three-dimensional, as is expected for string networks produced in brane inflation. Applying our model to a brane inflation setting, we find scaling solutions in which the effective 3D string motion does not necessarily stop, but it is slowed down because of the excitations trapped in the extra dimensions. These effects are likely to influence cosmic string network evolution for a long period after formation and we discuss their more general implications. ",
        "introduction": "Introduction} There has been a recent resurgence of interest in cosmic strings both for theoretical and observational reasons. Of particular interest is the generic possibility that cosmic strings can be produced at the end of an inflationary phase in models of brane inflation \\cite{SarTye,PogTWW}. The evolution of the network of the strings created in these models can be very different from the standard field theory case, thus providing a potential observational window on superstring physics \\cite{JoStoTye2,Vilenk,PolchProb}. On the observational side, there are perennial cases of astrophysical phenomena for which cosmic strings have been invoked as an explanation in the absence of some more orthodox mechanism; recent examples are two peculiar gravitational lensing events \\cite{Sazhin,Sazhin1,Schild, Kibblerev} but which need further independent follow-up. However, expected improvements in observational data, particularly from high resolution CMB experiments and gravitational wave detectors, present us with the very real prospect of detecting or constraining cosmic strings over a wide range of predicted energy scales (see, for example, \\cite{PogTWW, DamVil}). Further theoretical motivation has come from a recent phenomenological study of SUSY GUT models \\cite{Jeannerot}, which again found generic cosmic string production (for all cases which solve the monopole problem).  It is the recent work on cosmic strings appearing in higher dimensional theories, such as brane inflation, which primarily motivates our present study.  For spacetime dimension greater than four, strings no longer generically collide, so that loop production will be highly suppressed. Loops radiate away energy from the long string network, so this suppression will result in a much higher density of cosmic strings than in the usual $3+1$ dimensional case. Jones, Stoica and Tye \\cite{JoStoTye2} have estimated this enhancement by using a three-dimensional one-scale model and introducing an intercommuting probability $P<1$ to account for the fact that strings generically miss each other due to the presence of extra dimensions.  They suggest that the enhancement on the energy density of the string network is of order $P^{-2}$, which can be orders of magnitude different than the usual case.  This approach however does not take into account string velocities in the extra dimensions.  In general, cosmic strings are subject to the constraint that the average velocity squared of string segments must be less than $1/2$. Thus the fact that strings are moving in the extra dimensions will slow down their apparent three-dimensional motion. One might naively expect that velocities in the infinite dimensions will be redshifted by the expansion, while velocities in the compact dimensions will not if these dimensions are static. Hence there is the cosmologically dangerous possibility that velocities in the extra dimensions will accumulate and dominate while string motion in the three infinite dimensions will come to a halt. This can occur before SUSY breaking, when the fields corresponding to the string position in the extra dimensions are expected to become massive. However, if 3D string motion stops for long enough in the early universe, then the strings would irreversibly dominate the energy density of the universe making a subsequent hot big bang model impossible.  Thus we need to go beyond the simple analysis presented to date in order to gain a better quantitative understanding of cosmic string evolution in higher dimensional spacetimes, notably taking into account the important role of velocities in the extra dimensions.  String evolution in three spatial dimensions has been studied by various authors. Kibble \\cite{Kibble} described string networks by a single lengthscale, the `correlation length', and showed that it evolves towards a scaling solution in which it stays constant with respect to the horizon size. Bennett \\cite{Bennett,Bennett1} later modified this `one-scale' model with similar conclusions, subject to a condition on the efficiency of small loop production. The existence and stability of this scaling solution was verified by numerical simulations \\cite{BenBouch,AllShel}. These studies also revealed new physics at smaller scales, in particular the accumulation of significant small-scale structure on strings, which results in loop production at much smaller scales than initially thought.  To try to incorporate small-scale structure in analytic models, a number of different approaches have been attempted. These include a `kink-counting' model \\cite{AllCald,Austin}, a functional approach \\cite{Embacher}, a `three-scale' model \\cite{AusCopKib} and a `wiggly' model \\cite{thesis}. Including small-scale structure in analytical models comes with the cost of introducing several extra parameters, which need to be fixed by simulations.  However, the large-scale properties of string networks can be quantitatively described by the Velocity-dependent One-Scale (VOS) model \\cite{vos0,vos,vosk}, which does not suffer from this problem. By introducing a variable rms string velocity, the VOS model extends its validity from the friction dominated regime at early times, through the matter-radiation transition to $\\Lambda$-domination at late times , thus describing the complete cosmological history of string networks. Though it does not directly model small-scale structure, it provides a `thermodynamic' large-scale description of cosmic string evolution, which agrees remarkably well with high resolution numerical simulations.  The purpose of this paper is to extend the VOS model to spacetimes of higher dimension. Although strong motivation is provided by brane inflation, where the extra dimensions are small and stabilised, we intend to keep the discussion as general as possible to include time-varying extra dimensions, as for example in a $(D+1)$-dimensional FRW universe.  The paper is organised as follows. In section \\ref{dynamics} we discuss cosmic string dynamics in a FRW spacetime with isotropic and flat (but possibly expanding) extra dimensions. Starting from the Nambu-Goto action in this spacetime we derive the equations of motion as well as an expression for the energy of a cosmic string network. In section \\ref{evolution} we derive the averaged equations describing the evolution of cosmic strings in that $(D+1)$-dimensional spacetime. After we briefly review the $3+1$ VOS model, we extend it to $D+1$ dimensions and comment on some qualitative features of solutions. In section \\ref{braneinflation} we consider possible application of this extra-dimensional VOS (EDVOS) model to the case of brane inflation and discuss the dependence of the results on various parameters as for example the intercommuting probability of strings. We conclude in section \\ref{conc}. Finally, there is an Appendix where we derive approximate formulae for the momentum parameters of the $(D+1)$-dimensional VOS model. ",
        "conclusions": ""
    },
    "0410/nucl-ex0410025_arXiv.txt": {
        "abstract": "Detailed Monte-Carlo simulations were used to evaluate the performance of a liquid neon scintillation detector for dark matter and low-energy solar neutrino interactions. A maximum-likelihood event vertex fitter including PMT time information was developed, which significantly improves position resolution over spatial-only algorithms, and substantially decreases the required detector size and achievable analysis energy threshold. The ultimate sensitivity to WIMP dark matter and the $pp$ flux uncertainty are evaluated as a function of detector size. The dependence on the neon scintillation and PMT properties are evaluated. A 300 cm radius detector would allow a $\\sim$13 keV threshold,  a $pp$ flux uncertainty of $\\sim$1$\\%$, and limits on the spin-independent WIMP-nucleon cross-section of $\\sim10^{-46}$ cm$^2$ for a 100 GeV WIMP, using commercially available PMTs.  Detector response calibration and background requirements for a precision $pp$ measurement are defined. Internal radioactivity requirements for uranium, thorium, and krypton are specified, and it is shown that the PMT data could be used for an {\\it{in-situ}} calibration of the troublesome ${}^{85}$Kr. A set of measurements of neon scintillation properties and PMT characteristics are outlined which will be needed in order to evaluate feasibility and fully optimize the design of a neon-based detector. ",
        "introduction": "This paper presents details on the design of a dual-purpose liquid neon based detector for dark matter and low-energy solar neutrino interactions. The origin of dark matter in our universe and study of low-energy solar neutrinos are two of the foremost topics in particle astrophysics. Current observational evidence strongly suggests that a large fraction of the total matter in our universe is non-luminous, non-baryonic matter outside of the current standard model of particle physics.  A currently favored hypothesis is that the dark matter consists of WIMPs (Weakly Interacting Massive Particles). A natural candidate for this particle is the lightest superpartner which arises from supersymmetry, the neutralino.  Current direct dark matter searches attempt to detect WIMP interactions through elastic scattering on target nuclei, and this is the basis for WIMP detection in this study.  The current generation of solar neutrino experiments have shown that mixing occurs between the neutrino flavors.  One of the goals for next-generation experiments is to measure with high-precision the low-energy $pp$ neutrinos.  These neutrinos account for greater than 90\\% of all solar neutrinos, and are well-constrained by the measured solar luminosity.  A precision measurement of the flux of $pp$ neutrinos would lead to improved knowledge of neutrino mixing in the solar sector ($\\theta_{12}$),  tests of the neutrino oscillation prediction in the low-energy region where vacuum mixing dominates, and precision tests of solar evolution models.  WIMP particles ($\\chi$) and solar neutrinos ($\\nu$) can be observed in liquid neon through their elastic scattering from nuclei and electrons, respectively: \\begin{eqnarray} \\chi + {\\rm{Ne}} &\\longrightarrow& \\chi^{'} + {\\rm{Ne}}^{'} \\label{wimp_scatter} \\\\ \\nu + e^{-} & \\longrightarrow & \\nu^{'} + e^{-'} \\label{nu_scatter}. \\end{eqnarray} Monte-Carlo simulations were performed to evaluate the performance of a neon scintillation detector sensitive to the recoiling nuclei (${\\rm{Ne}}^{'}$) and electrons ($e^{-'}$) from equations~\\ref{wimp_scatter} and \\ref{nu_scatter}. The concept of a liquid neon scintillation detector for solar neutrinos was first proposed by McKinsey and Doyle~\\cite{clean_paper}. The experimental design of a next-generation WIMP dark matter experiment should allow a sensitivity increase of several orders of magnitude over present dark matter experiments to probe interesting regions of supersymmetry parameter space.  Ultra-low backgrounds and a large fiducial target are required to achieve this goal.  A large spherical volume of neon, along with precise event position reconstruction allows for a large fiducial mass and a large mass to surface area ratio, with the advantage that external source backgrounds are physically distant from the fiducial target and can be dealt with through position reconstruction.  Simulations were performed assuming nominal characteristics of scintillation light in liquid neon, and detector performance was estimated  for commercially available photomultiplier tubes (PMTs) with ultra-low background glass.  The details of the simulation inputs and algorithms are presented in section~2. The detector response to scintillation events is characterized in section~3, where reconstruction algorithms and response functions useful for evaluating predicted observables for dark matter and solar neutrino interactions are presented. The sensitivity to solar neutrinos is discussed in section~4.  Section~5 describes the reduction of external and internal backgrounds through reconstruction and analysis cuts, and the ultimate background requirements.  Optimizing the detector for dark matter and solar neutrino sensitivity is discussed in section~6. Calibration requirements for the full-scale experiment are outlined in section~7.  The requirements from photon absorption and sensitivity to neon scintillation and detector parameters are presented in section~8, along with a list of which measurements are needed to fully optimize the detector design.  Conclusions from this work, and potential advantages of this approach are presented in section~9. ",
        "conclusions": "In summary, we have demonstrated the possibility for a simultaneous dark matter and $pp$ neutrino experiment using liquid neon, assuming nominal scintillation characteristics and background contamination levels with detailed Monte-Carlo calculations.  This approach to direct dark matter detection has the advantage that WIMP event discrimination is based only on PMT signals, allowing for a simple detector design which can in principle be scaled to a very large target mass. This detector design also has the potential advantage that external sources of background are now well-separated from the inner target material, and can be dealt with through position reconstruction.  Assuming light attenuation and internal source radioactive backgrounds can be controlled, the large target mass achievable with neon may lead to the best sensitivity of any of the direct dark matter searches. We have shown that form factor suppression for high-A target nuclei leads to less of an advantage over lower-A target nuclei than that expected by the cross-section coherence alone.  A specific list of issues which must be addressed in order to determine the ultimate detector sensitivity has been provided.  An experimental program is underway to address these issues."
    },
    "0410/nucl-ex0410031_arXiv.txt": {
        "abstract": "A detailed examination of current data on \\xs17 (as opposed to an examination of \\S17 only) excludes quoting \\S17 with sufficiently small uncertainty. In contrast to suggestions that \\S17 is now known with the accuracy of \\xpm 3\\%, the exact value of \\S17 is dependent on the choice of the data and the choice of theory used for extrapolation. In addition recent high precision results (including the Seattle data) on \\xs17 which are in good agreement, still differ on the measured slopes, as do theoretical models that predict different d-wave contribution, precluding an accurate extrapolation to zero energy of the consistent data. Using a common extrapolation of only the consistent high precision data, suggests a value of \\S17 = 21.2 \\xpm \\ 0.5 eV-b, but a value equal to or smaller than 19.0 eV-b can not be excluded due to the uncertainty in the extrapolation, leading to an additional error of $^{+0.0}_{-3.0}$ eV-b. This (unacceptable) situation must be cleared by future experiments. . \\vspace{1pc} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410204_arXiv.txt": {
        "abstract": "Two families of models are currently considered to describe an accretion flow onto black holes and production of the observed X-ray radiation: (1) a standard cold accretion disk with a hot corona above it and (2) an outer truncated accretion disk with a hot semispherical inner flow. We compute spectra in the scenario with a hot inner flow surrounded by a truncated accretion disk covered by a hot corona and test the results on a sample of high redshift ($z > 4$) quasars observed with {\\it Chandra}. We find that in order to reproduce the ratio of optical to X-ray fluxes (the $\\alpha_{\\rm ox}$ parameter), the optical depth of the Comptonizing plasma has to be rather low ($\\tau = 0.02 - 0.15$ in the corona above the disk, and $\\tau = 0.15 - 0.70$ in the hot inner flow). This, together with the observed X-ray photon indices, implies either a high temperature in a thermal plasma ($kT_{\\rm e} = 90 - 500$ keV), or a nonthermal electron distribution in the plasma. We put an upper limit on the disk truncation radius, $r_{\\rm tr} \\leq 40 R_{\\rm S}$. The modeled accretion rate is high, $\\dot{m} > 0.2\\, \\dot{M}_{\\rm Edd}$, which may suggest that high-$z$ radio quiet quasars are analogs of X-ray binaries in their high or very high state. ",
        "introduction": "Modern X-ray observatories can detect the most distant objects in the universe and provide high quality data for quasars at redshifts as high as $\\sim 6$ (Becker et at.\\ 2001; Brandt et al.\\ 2002b; Bechtold et al.\\ 2003, hereafter B03), i.e. for objects born in the very early universe. By comparing high- and low-redshift sources, we may be able to determine whether these objects evolve and if there is any particular pattern that they follow during the evolution. This provides bases for studying the formation and evolution of the universe.  We focus our studies on the radio quiet quasars (RQQs), which comprise $\\sim85\\%-90\\%$ of the quasar population (e.g., Stern et al.\\ 2000). The fact that RQQs are more common, and thus more typical, among the quasars makes it important to understand the origin of their broad-band spectra. In these systems a contribution from possible additional components related to the radio emission should be negligible. On the other hand, RQQs are X-ray quieter than radio loud quasars (RLQs) in terms of higher $\\alpha_{\\rm ox}$ parameter --- the spectral index of a reference power law connecting the optical and X-ray rest frame fluxes, for which we use the following definition: $\\alpha_{\\rm ox} = - \\left ( \\log{f_{\\rm x}} - \\log{f_{\\rm o}} \\right ) / \\left ( \\log{\\nu_{\\rm x}} - \\log{\\nu_{\\rm o}} \\right )$, where $f_{\\rm o}$ and $f_{\\rm x}$ are rest frame optical and X-ray fluxes at $\\lambda = c / \\nu_{\\rm o} = 2500$ \\AA\\ and $E = h \\nu_{\\rm x} = 2$ keV (e.g., Zamorani et al.\\ 1981; Worrall et al.\\ 1987; Laor et al.\\ 1997) --- so the X-ray data of RQQs are of worse quality than those of RLQs.  Understanding observed spectral features in terms of theoretical model parameters can provide important insights into physical processes taking place in quasars, especially about the origins of the hard X-ray radiation.  The ``standard'' model so far successful in explaining optical/UV/X-ray spectra of active galactic nuclei (AGNs) (and those of X-ray binaries, hereafter XRBs) contains as a crucial element an optically thick accretion disk around a central black hole. Gravitational energy is dissipated and thermalized in the accretion disk (Shakura \\& Sunyaev 1973). It is emitted in the form of thermal radiation, giving a characteristic bump in the optical/UV band (``big blue bump'') in AGN spectra (Shields 1978, Czerny \\& Elvis 1987, Siemiginowska et al.\\ 1995, Koratkar \\& Blaes 1999) or soft X-rays in the case of XRBs (see, e.g., a review by McClintock \\& Remillard 2003). In the case of high-$z$ RQQs, the central black hole is expected to be supermassive, with a mass of up to $\\sim 10^{10}\\, M_{\\odot}$ (Rees 1984). Such a high mass is suggested by the high luminosities of high-$z$ RQQs (of the order of $\\sim10^{48}$ ergs~s$^{-1}$), which should not exceed the Eddington luminosity of $L_{\\rm Edd} \\approx 1.3 \\times 10^{46} \\left ( M / 10^8\\, M_{\\odot} \\right )$~erg~s$^{-1}$.  Some of the disk photons gain energy in the inverse Compton process by scattering off energetic electrons in a hot, optically thin plasma. This mechanism is generally accepted to explain hard power--law--like spectral X-ray tails extending beyond 100 keV (see Mushotzky, Done \\& Pounds 1993 for a review; Done 2001).  The spectra commonly show cut-offs beyond $\\sim 100$ keV in both Seyfert galaxies and XRB (in their hard/low states), suggesting a thermal energy distribution of Comptonizing electrons (Gierli\\'nski et al.\\ 1997, Zdziarski et al.\\ 1998).  Observed X-ray spectra of XRBs (soft/high states) may extend to energies beyond $500$ keV (e.g., Gierli\\'{n}ski et al.\\ 1999) with no cut-off, implying a non-thermal energy distribution or, possibly, a hybrid (i.e., a mixture of thermal and non-thermal; Coppi 1999) electron distribution.  Some of the Comptonized X-ray photons are likely to illuminate the accretion disk, where they can be reprocessed and mostly thermalized, thus increasing the disk flux.  Since the disk flux provides the seed photons for the inverse Compton processes, a feedback loop between the cold disk and the hot plasma is created (Haardt \\& Maraschi 1991, 1993).  Some of the illuminating X-ray photons will be reflected by Compton scattering off cold electrons and emerge back from the disk, with a characteristic spectrum (Lightman \\& White 1988; George \\& Fabian 1991).  The spectrum is composed of (1) a broad continuum (``reflection hump'') peaking at 20--30 keV, whose shape is determined by a competition between photo-absorption in the soft X-ray band ($E<10$ keV) and Klein-Nishina effects on the Compton scattering cross-section in the hard X-ray band, (2) the Fe K$_{\\alpha}$ fluorescence/recombination line at 6.4--6.97 keV (depending on the Fe ionization state), and (3) the Fe-K$_{\\alpha}$ edge at $\\sim7.1$--10 keV.  The fluorescence features can be further influenced by relativistic effects and/or the ionization state of the reflecting medium (Fabian et al.\\ 1989; {\\.Z}ycki \\& Czerny 1994, Nayakshin 2000).  An unresolved issue is that of the geometry of the accretion flow, i.e., a geometrical configuration of the two phases of accreting plasma (cold, optically thick; and hot, optically thin). There are two general scenarios considered in the modeling: (1) the plane-parallel geometry, in which a hot corona is located atop the accretion disk, which is assumed to extend down to the innermost stable orbit ($r_{\\rm in} = 3\\, R_{\\rm S}$ for a Schwarzschild black hole, where Schwarzschild radius $R_{\\rm S} \\equiv 2GM/c^2$), and (2) the spherical geometry, which assumes that the cold disk is truncated at a certain radius $r_{\\rm tr} > r_{\\rm in}$ and surrounds a hot inner plasma flow within $r < r_{\\rm tr}$ (Shapiro et al.\\ 1976; Poutanen et al.\\ 1997; Esin, McClintock \\& Narayan 1997; R\\'{o}\\.{z}a\\'{n}ska \\& Czerny 2000). In the first geometry, the hot corona sandwiching the disk may be continuous (e.g., Haardt \\& Maraschi 1991) or ``patchy'' i.e., composed of a number of active regions driven by, e.g., magnetic activity (Galeev, Rosner \\& Vaiana 1979).  A variant of the spherical geometry (case 2 above) is the spherical cloud model of Collin-Souffrin et al.\\ (1996), in which the cold plasma forms small clouds accreting roughly spherically onto the central X-ray source.  In principle, one can distinguish between the two basic geometries by studying the X-ray spectrum in detail (see Done 2001 for a review). One argument is based on the amplitude of the reflection component. The amplitude is defined as \\be \\label{eq:refl} R \\equiv \\Omega/2\\pi, \\ee where $\\Omega$ is a solid angle subtended by the reflector as seen from the source of X-rays. For plane-parallel geometry, $R$ should yield values close to unity (the disk covers $\\approx 2\\pi$ from the point of view of the X-ray source), whereas for spherical geometry $R<1$ would be expected. Furthermore, the amount of broadening of the reflection features gives an estimate of the inner radius of the accretion disk (e.g., Done et al.\\ 2000).  For example, the Fe K$\\alpha$ line would be narrower if the disk were truncated far away from $r_{\\rm in}$.  Another argument in determining the geometry may be based on the X-ray photon index, $\\Gamma$. A continuous, plane-parallel corona produces rather soft X-ray spectrum, $\\Gamma>2$, even when the entire gravitational energy is dissipated in the corona (Haardt \\& Maraschi 1991, 1993).  This is because the underlying cold disk provides a lower limit to the cooling flux. The spectra in the spherical geometry are harder, $\\Gamma < 2$, because the disk flux entering the central sphere is strongly reduced, compared to that in the plane-parallel geometry (e.g., Poutanen et al.\\ 1997).  However, these diagnostics do not determine the geometry in a unique way, since it is possible to modify the continuous plane-parallel corona scenario in such a way that the predicted spectra are hard, with $\\Gamma<2$, and the amplitude of reflection is reduced, $R<1$. This may be realized in a number of scenarios: (1) the hot plasma may form a patchy corona (Stern et al.\\ 1995); (2) the cold disk may be covered by a hot ionized skin, decreasing the effectiveness of the reprocessing/thermalization (Nayakshin, Kazanas \\& Kallman 2000); or (3) the hot plasma may be outflowing at relativistic speeds, so that relativistic beaming makes the emission pattern significantly anisotropic (Beloborodov 1999a). In particular, both the spherical geometry scenario and models 2 and 3 above, reproduce the correlation observed between the X-ray spectral index and amplitude of reflection of both Seyfert galaxies and XRB (Zdziarski et al.\\ 1999).  An additional model diagnostic, which helps to break the degeneracy, is the ratio of observed optical and X-ray fluxes, the X-ray loudness, $\\alpha_{\\rm ox}$. This parameter is of more use for AGNs than for XRBs, since the disk emission can be more easily observed in AGNs than in XRBs. This is especially true for high-$z$ quasars, where the big blue bump may be redshifted into the observable optical band. On the other hand, the quality of X-ray data from these objects is usually not sufficient to accurately determine the amplitude of the reprocessed component (see Siebert et al.\\ 1996, Lawson \\& Turner 1997, Cappi et al.\\ 1997, Reeves \\& Turner 2000, Mineo et al.\\ 2000, and Page et al.\\ 2003 for attempts to fit reflected component and/or iron K$_{\\alpha}$ line to the data of quasars).\\\\  The ``complex'' models described above were applied to Seyfert galaxies and low redshift RQQs with high S/N (signal-to-noise ratio) data (recently, e.g., by Chiang 2002; Chiang \\& Blaes 2003; Janiuk, Czerny \\& Madejski 2001).  So far the quasars' spectral energy distribution, especially that of high redshift quasars, has been modeled by the thermal emission from an accretion disk (Sun \\& Malkan 1989, Bechtold et al.\\ 1994a, 1994b, Tripp et al.\\ 1994), with the addition of a Comptonizing medium explaining the X-rays (Band \\& Malkan 1988, Fiore et al.\\ 1995). There have been no systematic studies of the parameter space of these models appropriate for high$-z$ quasars, while the number of the observations has been increasing.  In this paper we consider a geometry with a hot inner flow surrounded by a truncated accretion disk covered with a hot corona. We apply the model to the high redshift RQQs. Similar geometry was applied by Chiang (2002) and Chiang \\& Blaes (2003) to Seyfert galaxies. Their Compton scattering medium was located within the central sphere, which was surrounded by an accretion disk. They extended the model of Zdziarski et al.\\ (1999) by introducing viscous dissipation in the cold disk and modeled X-ray data simultaneously with the optical/UV data. In our model we include two Comptonizing media, the sphere and the corona. We explore the model parameter space and identify the model parameters suitable to generate the observed spectra of RQQs. The observational constraints include three quantities: the X-ray photon index, $\\Gamma$, the X-ray loudness, $\\alpha_{\\rm ox}$, and the optical/UV luminosity, $l_{\\rm UV} \\equiv \\log \\left ( \\nu L_{\\nu} \\right )$ at 2500 \\AA\\, in the rest frame. Optical/UV luminosities of high-$z$ RQQs are of the order of $l_{\\rm UV} \\gtrsim 46.0$--47.0 (B03 and references therein). B03 and Vignali et al.\\ (2003a, hereafter V03; 2003b) derived the mean $\\sim 2$--30 keV rest frame X-ray photon index to be in the range $\\sim 1.5$--2. The X-ray loudness was found to be $\\sim 1.7$--1.8 (B03; V03; Vignali et al.\\ 2003b).  The structure of the paper is as follows: In Section \\ref{sec:model}, we describe the geometry of accretion flow and describe the model. In Section \\ref{sec:results}, we present the results of computations, and discuss the constraints imposed on the parameter space by the observations. In Section \\ref{sec:app}, we apply the spherical geometry accretion model to the high-$z$ RQQ data of B03. Finally, Sections \\ref{sec:disc} and \\ref{sec:concl} contain discussion and concluding remarks, respectively.  In a subsequent paper (Sobolewska et al.\\ 2004, hereafter Paper II), we study the accretion disk model covered by a non-uniform corona composed from hot clouds (a patchy corona).\\\\ ",
        "conclusions": "\\label{sec:concl}  In this paper we explored a parameter space of the semi-spherical accretion models applicable to high-$z$ RQQs. Our results show the following:  \\noindent 1. The observed optical/UV luminosity suggests that high redshift quasars host supermassive black holes with masses of the order of $10^{10} M_{\\odot}$. The modeling results in accretion rates higher than $\\sim 0.2$ times the Eddington rate.  \\noindent 2. The high accretion rates may suggest similarity between high-$z$ RQQs and XRBs in the high or very high state. If high-$z$ RQQs correspond to high or very high states of X-ray binaries, the high values of plasma temperature (up to 500 keV) and low optical depths might suggest instead that Comptonization takes place in a non-thermal or hybrid plasma.  \\noindent 3. The observed X-ray loudness puts strong constraints on the modeled plasma temperature and optical depth. {\\it Chandra} or {\\it XMM-Newton} observations with better S/N are needed to give tighter constraints on the X-ray slope. Future mission (such as {\\it Constellation-X}) can provide information about the high-energy cut-off and the reflected component that is necessary to uniquely determine geometry of the accretion flow.  \\noindent 4. Precise observations constraining the optical/UV spectral index for individual sources would help to reduce the number of model parameters and perform more accurate fits. However, the intrinsic reddening of the quasar spectrum is the main observational uncertainty in this spectral region.  \\noindent 5. With the present data, a number of theoretical descriptions are possible. In Paper II, we focus on another interesting possibility, in which the accretion disk extends to the last stable orbit and the hot Comptonizing corona is formed by clouds above the disk (the ``patchy'' corona geometry)."
    },
    "0410/astro-ph0410174_arXiv.txt": {
        "abstract": "We present {\\it BVRIK}~images and spectroscopic observations of the $z=0.17$\\ host galaxy of the compact steep-spectrum radio source MRC~B1221$-$423. This is a young ($\\sim10^5$~yr) radio source with double lobes lying well within the visible galaxy. The host galaxy is undergoing tidal interaction with a nearby companion, with shells, tidal tails, and knotty star-forming regions all visible.  We analyse the images of the galaxy and its companion pixel-by-pixel, first using colour-magnitude diagrams, and then fitting stellar population models to the spectral energy distributions of each pixel. We also present medium-resolution spectroscopy of the system.  The pixels separate cleanly in colour-magnitude diagrams, with pixels of different colours occupying distinct regions of the host galaxy and its companion.  Fitting stellar population models to these colours, we have estimated the age of each population. We find three distinct groups of ages: an old population ($\\tau \\sim 15$~Gyr) in the outskirts of the host galaxy; an intermediate-age population ($\\tau \\sim 300$~Myr) around the nucleus and tidal tail, and a young population ($\\tau \\simlt 10$~Myr) in the nucleus and blue ``knots''.  The spectrum of the nucleus shows numerous strong emission lines, including \\oi\\ $\\lambda$6300, \\oii\\ $\\lambda$3727, \\sii\\ $\\lambda\\lambda$6716, 6731, H$\\alpha$, and \\nii\\ $\\lambda\\lambda$6548, 6583, characteristic of a LINER spectrum. The companion galaxy shows much narrower emission lines with very different line ratios, characteristic of a starburst galaxy.  We have evidence for three distinct episodes of star formation in B1221$-$423. The correlation of age with position suggests the two most recent episodes were triggered by tidal interactions with the companion galaxy. The evidence points to the AGN in the centre of B1221$-$423 having been ``caught in the act'' of ignition.  However, none of the components we have identified is as young as the radio source, implying that the delay between the interaction and the triggering of the AGN is at least $3 \\times 10^8$~years. ",
        "introduction": "\\label{sec:intro}  The onset of AGN activity in galaxies appears to be closely related to starbursts, possibly triggered by mergers, which lead to large increases in the amount of material being fed to the central black hole, thereby triggering the radio emission.  Many radio galaxies show signatures of tidal interactions, such as tails, bridges, shells, and double nuclei \\citep[see][for a review]{bh92}.  HST observations of powerful radio galaxies at $z \\sim 1$\\ show disturbed, knotty rest-frame UV emission in the host galaxies \\citep*{lbr95,blr96}.  \\begin{figure*} \\centerline{\\psfig{figure=figure1.ps,width=\\textwidth,clip=t}} \\caption{{\\it BVRIK} images of MRC~B1221$-$423. Knots of star-formation and a tidal tail joining the galaxy to the companion are clearly visible. The size of each image is 22~arcsec; the separation of the galaxy and the companion is 10~arcsec ($1\\,\\mathrm{arcsec} = 2.88\\,\\mathrm{kpc}$\\ assuming $H_0 = 71\\;{\\rm km\\,s^{-1}\\,Mpc^{-1}}$\\ and $\\Omega_\\mathrm{M}=0.27$). The orientation of the slit for the spectroscopic observations (PA=$-$18\\fdg5) is shown as the line joining the galaxy to the companion on the $B$\\ image. The labels on the $V$\\ band image indicate the various regions described in the text.} \\label{fig:images} \\end{figure*}  Compact steep-spectrum (CSS) sources are bright but compact radio galaxies, with sizes 1--20~kpc. Unlike the majority of powerful radio galaxies, the radio source is wholly contained within the envelope of the host galaxy. They have a steep spectral index, $\\alpha < -0.5$\\ ($S_\\nu \\propto \\nu^\\alpha$), and a radio spectrum peaking below 500~MHz.  CSS sources are powerful ($P_{1.4} \\ga 10^{25}\\;{\\rm W\\;Hz}^{-1}$), and make up about 20\\% of bright flux-limited low-frequency radio catalogues like the 3CR catalogue \\citep{ben62} or the Molonglo Reference Catalogue \\citep{lml+81}.  They typically have redshifts between 0.1--2 \\citep[see][for a review]{ode98}. The two competing models for CSS sources are {\\em (a)} that they are small in size because the jets are ``frustrated'' by interaction with dense surrounding gas; and {\\em (b)} that they are young sources which will eventually evolve into classical large double-lobed radio sources \\citep{ode98}.  Most evidence suggests that CSS sources are young \\citep[e.g.][]{ffs+90}, and will fade as they grow into large-scale radio sources. Significant luminosity evolution is required to explain the fact that the age of CSS sources is only 1\\% that of the large radio sources, yet 20\\% or more of bright centimetre-wavelength--selected samples are CSS sources \\citep[e.g.][]{sch74b,ka97,rtpw96}. The best evidence for the youth of CSS sources is the detection of hotspot advance speeds of 0.1 to 0.3$c$\\ in compact symmetric objects \\citep[see][and references therein]{con02}.  MRC B1221$-$423 is one of the nearest CSS sources, located in a host galaxy with a redshift $z=0.1706$\\ \\citep{scrw93}. The radio source has a steep spectrum ($\\alpha = -0.85$) and double lobes with a separation of only 1\\farcs5, which lie well within the envelope of the visible galaxy \\citep*{shp03}.  The structure is suggestive of the lobes seen in larger double sources; there is no evidence of a core. \\citet{shp03} estimated the age of the radio source to be $10^5$~yr, using the tight correlation between the total kinetic power of the jet and the narrow-line luminosity \\citep{rs91}. At the distance of the galaxy, 1~arcsec corresponds to a distance of 2.876~kpc, assuming $H_0 = 71\\;{\\rm km\\,s^{-1}\\,Mpc^{-1}}$\\ and $\\Omega_\\mathrm{M}=0.27$.  The host galaxy is highly disturbed, showing signs of tidal interaction with at least one close companion. An archival image taken using the Anglo-Australian Telescope in 1\\farcs4 seeing shows faint concentric arcs to the north-east \\citep{shp03}, evidence of past tidal interaction.  We have obtained multi-colour images of MRC~B1221$-$423 in order to study the change in properties of the stellar population across the galaxy.  The use of spatially-resolved colours to study differences in stellar populations across a galaxy has been shown to be an effective way to study age distributions, localised extinction and star formation history in galaxies. The technique was pioneered by \\citet{bot86}, who used a spatially resolved $(B-R)$ vs. $B$\\ diagram to examine the star formation history of NGC$\\,$4449. More recently, \\citet{aef+99} used pixel colour-magnitude diagrams to investigate the evolutionary history of galaxies in the \\textit{Hubble Deep Field}, while \\citet{kzc+00}, \\citet{dlh+03} and \\citet{kfp+03} used such diagrams to study various nearby galaxies.  In this paper, we present and analyse five-colour images and optical spectroscopic observations of MRC~B1221$-$423.  In \\S\\ref{sec:obs} we describe the observations and their reduction.  In \\S\\ref{sec:Modell-pixel-colo}, we analyse the images, first by means of colour-magnitude diagrams, and then by fitting spectral energy distributions to each pixel. In \\S~\\ref{sec:Surface-photometry} we examine the surface-brightness profile of the galaxy.  In \\S\\ref{sec:Spectroscopy-results} we discuss the spectroscopic results, and then discuss the implications for the history of the galaxy in \\S\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  We have shown that the host galaxy of MRC~B1221$-$423 and its interacting companion can be resolved into different stellar populations. We have achieved this in two different ways, using colour-magnitude diagrams, and by using models for the spectral energy distribution of individual pixels. Both these methods separate the pixels into distinct regions with different colours and/or ages. We interpret this as indicating that the galaxy and its companion have undergone several episodes of star formation, possibly associated with the tidal interaction between the two.  The ages we deduce for the stellar population fall into three distinct groups. An old population, with age $t \\sim 15$~Gyr, occupies the outskirts of the host galaxy. An intermediate age population, with $t \\sim 300$~Myr, is found near the nucleus and around the tidal tail joining the two galaxies, while the youngest population, with ages less than 10~Myr, is concentrated in the nucleus and the blue regions to the west, as well as in the companion galaxy.  The location of these populations is not identical to the separation suggested by the regions in the colour-magnitude diagram (Figure~\\ref{fig:regions}). This may be due to the effects of dust, which our SED modelling suggests is concentrated in the nucleus of the host galaxy.  The relationship of the different stellar populations with the underlying powerful radio source is much less clear. The age of the source, estimated from the power of the radio jets and total stored energy in the synchrotron plasma of the lobes \\citep{rs91}, is only $10^5$~yr \\citep{shp03}.  This is much younger than the youngest population we detected, with an age of about 10~Myr (Section~\\ref{sec:Fitt-spectr-energy}), and suggests there must be a substantial time delay between the most recent burst of star formation and the triggering of the radio source.  A possible scenario to explain the three component populations is as follows: the tidal interaction began 300~Myr ago, which triggered star formation in both the host galaxy and the companion. We see the results of this episode of star formation as the intermediate-age population.  Gas is driven down into the nucleus on timescales of $\\sim 10^8$~y, where it both cools to form stars, which we see as the 10~Myr-old population, and triggers activity in the central supermassive black hole. Delays of this order are predicted in theoretical calculations, \\citep[e.g.][]{lpr88}, which give time delays of a few$\\times 10^8$~yr between the tidal interaction and the gas reaching the centre of the galaxy.  This would be classed as a ``minor merger'', involving interaction of the galaxy with a dwarf satellite. Such interactions have been shown to be effective in feeding material into the nuclear regions \\citep{hm95}. This scenario for B1221$-$423 would fit with the idea that low redshift radio galaxies are triggered by material falling into a pre-existing supermassive black hole, in contrast with very high redshift radio galaxies and quasars, which are often associated with major mergers and the simultaneous formation of a massive bulge and the central black hole.  \\label{sec:Acknowledgments}  We thank Oliver Prouton for obtaining the optical images of B1221$-$423, and the anonymous referee for helpful comments. RWH and EMS acknowledge support for this project from the Australian Research Council and the University of Sydney Sesqui grant scheme."
    },
    "0410/astro-ph0410497_arXiv.txt": {
        "abstract": "{ \\em GALEX} has identified large numbers of extra galactic objects that are too faint to be observed by {\\em FUSE}.  These star forming systems are potential sources of Lyman continuum (Lyc) radiation and may contribute to the metagalactic background radiation that ionizes most of the baryons in the universe.  By measuring the level of Lyc radiation leaking from these galaxies we can quantify the contribution of stars, relative to quasars, to the ionizing background at low redshift ($0.025 \\la z \\la 0.4$) and provide insight into the reionization processes at high redshift.  A meter class instrument with a moderate spectral dispersion, wide-field, multi-object spectrograph can detect Lyc leak from a statistically significant number of these galaxies, if the average escape fraction $ f_{esc} \\ga$ 1\\%.  Such an instrument will have enormous discovery potential because of the large number of extra-galactic objects that exist within its grasp.  LyContEx  will be scientifically productive, synergistic with existing surveys, and can be developed into a far ultraviolet mission in the near future. ",
        "introduction": "\\vspace*{-.1in}  For all its considerable success {\\em FUSE} has three limitations.  The spectral resolution is too low, the spectral resolution is too high and the \\'{e}tendue ($A_{eff}\\Omega$, grasp, information content)  is limited to a single object.  The challenge for future far ultraviolet missions is to identify compelling new science opportunities enabled by instrumentation that overcomes these limitations.  Advances in high resolution spectroscopy ($R \\ge$ 20000) will require a very large telescope  \\citep{Moos:2004}, a new approach to spectrograph design, and a multi-decadal timescale for development.  A meter class, high \\'{e}tendue, moderate dispersion far ultraviolet spectrograph will have a shorter development timescale and an enormous potential for discovery because of the large number of extra galactic sources within its grasp too faint to be observed by {\\em FUSE}.  One new science opportunity enabled by such an instrument will be the search for the elusive Lyc radiation leaking from low redshift star-forming galaxies.   \\citet{McCandliss:2004} have detailed a new moderate dispersion, wide field, far ultraviolet spectroscopic survey instrument incorporating an array of micro-shutters \\citep{Moseley:2000} designed to reach significantly lower sensitivities than {\\em FUSE} and increase the information gathering power.  Observations of objects with flux levels below the {\\em FUSE} background equivalent flux (BEF $\\sim 5 \\times 10^{-15} \\:{\\rm ergs\\:cm^{-2} s^{-1}\\:\\AA^{-1}}$),  face an exposure time ``Wall'' requiring increasingly deep exposures.  A LyContEx mission with  $R \\sim$ 3000  will yield a BEF $\\sim 2 \\times 10^{-17} \\:{\\rm ergs\\:cm^{-2}\\:s^{-1}\\:\\AA^{-1}}$.  \\citet{Xu:2004} finds  $>$ 850 gal $(\\deg)^{-2}$ brighter than  23 mag$_{AB}$ = $3 \\times 10^{-17} \\:{\\rm ergs\\:cm^{-2}\\:s^{-1}\\:\\AA^{-1}}$ at 1540 \\AA\\ from {\\em GALEX} number counts (Figure 1).  I show here a statistically significant number of these galaxies will have detectable levels of Lyc if the $ f_{esc} \\ga$ 1\\%.    \\begin{figure*} \\vspace*{.2in} \\psfig{file=mccandliss_fig1.ps,angle=90,width=4.5in} \\caption[]{Differential number counts from {\\em GALEX} observations. The cross hatched areas have been integrated to determine the total number of galaxies brighter than the limiting fluxes on the top of the figure. The integrated number of galaxies per $(\\deg)^{2}$ are printed under the cross-hatchings. } \\label{numcts} \\vspace*{-.2in} \\end{figure*} ",
        "conclusions": ""
    },
    "0410/astro-ph0410668_arXiv.txt": {
        "abstract": "Supernovae (SNe) were long suspected as possible progenitors of gamma-ray bursts (GRBs).  The arguments relied on circumstantial evidence.  Several recent GRBs, notably GRB~030329, have provided direct, spectroscopic evidence that SNe and GRBs are related.  The SNe associated with GRBs are all of Type Ic, implying a compact progenitor, which has implications for GRB models.  Other peculiar Type Ic SNe may help to expand understanding of the mechanisms involved. ",
        "introduction": "The mechanisms that produce gamma-ray bursts (GRBs) have long been among the mysteries of modern astrophysics.  There have been a wide variety of models (see M{\\'e}sz{\\'a}ros 2002 for a review of the theories of GRBs).  The discovery of optical afterglows opened a new window on the field (see, e.g., van Paradijs, Kouveliotou, \\& Wijers 2000).  Early identification of the optical afterglows led to the measurement of redshifts for several GRBs (e.g.,~GRB~970508: Metzger et~al.~1997), providing definitive evidence for their cosmological origin.  Observations at other wavelengths, especially radio, have revealed many more details about the bursts (e.g., Berger et al. 2000; Frail et al. 2003).  This only applies to the long-duration bursts at this point.  Models that use supernovae (SNe) to explain GRBs have been part of the discussion from the start (e.g., Colgate 1968; Woosley 1993; MacFadyen \\& Woosley 1999).  Strong evidence for the GRB-SN association was first provided by GRB~980425; the optical afterglow was not found, but an unusual Type Ic SN was seen in the error box of the GRB (Galama et al. 1998; Patat \\& Piemonte 1998).  Many other circumstantial arguments were made to connect GRBs with SNe.  Direct evidence, where an afterglow is seen to transform into an SN, was first seen with GRB~030329 (Stanek et~al. 2003; Matheson et~al. 2003; Hjorth et~al. 2003; Kawabata et~al. 2003; Kosugi et~al. 2004).  Subsequently, two other cases have been identified: GRB~021211 (Della Valle et~al. 2003) and GRB~031203 (Malesani et~al. 2004).  In this paper, I will review the evidence for the GRB-SN association and describe the nature of the SNe that are associated with GRBs. ",
        "conclusions": "While circumstantial arguments had linked SNe with GRBs, spectroscopic data was needed for irrefutable proof.  This crucial piece of evidence arrived in the form of GRB~030329, wherein a direct transformation from GRB afterglow to peculiar Type Ic SNe was thoroughly observed.  Later that same year, GRB~031203 provided a further link.  These two, along with SN~1998bw, showed similar spectra.  That they were Type Ic implied that the progenitor was a stripped-envelope star, lending credence to jet-based models for GRBs.  Although these GRB-SNe are rare, low-redshift examples of peculiar Type Ic SNe are not, providing a laboratory to understand stellar explosions in extreme circumstances."
    },
    "0410/astro-ph0410424_arXiv.txt": {
        "abstract": "We present a deep VLA search for HI emission from the low-luminosity Virgo Cluster elliptical galaxy NGC 4476, which contains 1.1$\\times$10$^8$$M_{\\odot}$ of molecular gas in an undisturbed disk in regular rotation.  No HI was detected.  The rms noise in the final image corresponds to a 3$\\sigma$ column density sensitivity of 1.2$\\times$10$^{20}$ cm$^{-2}$ at the position of NGC 4476, averaged over the 4~kpc beam.  The total HI mass is less than 1.5$\\times$10$^{7}$ $M_{\\odot}$. If we compare our HI upper limit to the H$_2$ content, we find that NGC 4476 is extremely deficient in HI compared to other galaxies detected in these two species.  The H$_2$/HI mass ratio for NGC 4476 is~ $>$ 7, whereas typical H$_2$/HI ratios for elliptical galaxies detected in both HI and H$_2$ are $<$~2.  Based on this extreme HI deficiency and the intra-cluster medium (ICM) density at the projected distance from M87 we argue that either NGC 4476 has undergone ram-pressure stripping while traveling through the Virgo cluster core or its average molecular gas density is larger and its interstellar UV field is smaller than in typical spiral galaxies. NGC 4476 is located 12\\arcmin~ in projection from M87, which causes extreme continuum confusion problems. We also discuss in detail the techniques used for continuum subtraction.  The spectral dynamic range of our final image is 50,000  to 1. ",
        "introduction": "In recent years it has become clear that the majority of elliptical galaxies contain small amounts of cold interstellar gas and dust (eg. Knapp \\etal~1989; Colbert \\etal~2001; Sadler \\etal~2000).  Fifty to seventy percent of ellipticals show HI emission in amounts M(HI)/L$_B$$\\geq$10$^{-3}$ (Huchtmeier, Sage, $\\&$ Henkel 1995). The origin of the neutral atomic gas is still being debated; possibilities include accretion of gas-rich companions, fallback of gas after the major merger event that created the elliptical, secondary infall, and (especially for the smaller systems) stellar mass loss (e.g.\\ van Gorkom \\& Schiminovich 1997; Oosterloo et al.\\ 1999).  Work done by Quillen \\etal~(1992), Wiklind \\etal~(1997), and Young (2002) has shown that molecular gas in elliptical galaxies can be found in flat disks in regular rotation about the central nucleus.  HI studies in low-luminosity elliptical galaxies show that their HI systems fundamentally resemble those of spiral galaxies (Sadler \\etal~2000).  The primary difference between the HI disks in low-luminosity ellipticals and those in spiral galaxies is that the surface density of the HI disk is a factor of 5 to 10 lower in ellipticals (Oosterloo \\etal~1999).  Typical $M_{H_2}$/M$_{HI}$ for spiral and elliptical galaxies are 2 (Kenny $\\&$ Young 1989, Thronson \\etal~1989) and 0.8 (Lees \\etal~1990,  Wiklind \\etal~1995), respectively.  It is well established that cluster spirals are deficient in HI (Solanes \\etal~2001).  HI imaging of a few clusters has shown that the HI disks have smaller diameters in galaxies closer to the centers of clusters (Warmels 1988, Cayatte \\etal~1990; Bravo-Alfaro \\etal~2000). Studies of the Virgo Cluster show extreme HI deficiencies for galaxies within 6$^{\\circ}$ of the cluster center.  Inside this region, six spiral and lenticular (S0) galaxies have $M_{H_2}$/M$_{HI}$ ratios of 10 and greater (see Table 1).  The fate of the HI in these galaxies has been the subject of  much speculation.  Two main processes are thought to affect the gas content of galaxies in clusters: (i) ICM-ISM interactions (Gunn $\\&$ Gott 1972; Nulsen 1982; Valluri $\\&$ Jog 1990; Schulz $\\&$ Struck 2001; Vollmer \\etal~2001), and (ii) gravitational interactions, including galaxy-galaxy encounters and cluster-galaxy encounters (e.g.\\ Mihos 2004; Moore, Lake, \\& Katz 1998; Henriksen \\& Byrd 1996).  ICM-ISM interactions would affect the gas in a galaxy but not the existing stellar populations, whereas gravitational interactions would affect both gas and stars. Since the gas disks in low-luminosity ellipticals resemble those of spirals, it is very likely that the cluster environment will have the same effects on HI disks in both spirals and ellipticals.  In this paper we investigate this type of alteration of the ISM in the low-luminosity elliptical NGC 4476.  Spirals which are deficient in HI but not in H$_2$ may have enhanced H$_2$/HI mass ratios.  The largest H$_2$/HI mass ratios in Virgo spirals are listed in Table \\ref{tbl-1}, and so far no elliptical galaxies are known to have mass ratios as large as these.  Of the ellipticals which are detected in CO, Wiklind \\etal~(1995), Lees \\etal~(1991), and Georgakakis \\etal~(2001) find only five with $M_{H_2}$/M$_{HI}$ of at least 3, and these ratios are all based on HI upper limits except for NGC 2623 (see Table 2).  The remaining six entries in Table 2 are H$_2$ non-detections so the H$_2$/HI ratios calculated for these galaxies are upper limits.  The lenticulars which are studied in detail by Welch and Sage (2003) have H$_2$/HI in the range 0.01 to 1.5 which may be related to the fact that their sample excludes Virgo cluster members.  In this paper we present a search for HI emission from the low-luminosity Virgo elliptical NGC 4476.  The main challenge in this observation was to achieve good spectral dynamic range, since M87 is a strong radio source ($\\sim$220 Jy at 1.4 GHz).  In Sec. 3 we describe the method used to achieve optimal bandpass calibration and continuum subtraction.  No HI was detected in the entire 40\\arcmin~ by 40\\arcmin~ by 647 km s$^{-1}$ data cube.   In Sec. 4 we discuss the implications of this null result, and provide evidence that the HI gas has been stripped by the ICM surrounding M87 and/or the physical conditions of the ISM in NGC 4476 are very different than observed for normal stripped spirals. In Sec. 5 we discuss the effects of ram pressure on the star formation in NGC 4476.    In Sec. 6 we present our conclusions. ",
        "conclusions": "In this paper we present high dynamic range HI observations of the low-luminosity elliptical galaxy NGC 4476.  In  summary:\\\\\\\\ $\\bullet$ The spectral dynamic range of the final image (estimated as the peak continuum intensity in the field divided by the rms noise level in spectral channels) is $\\sim$~50,000.\\\\  \\\\ $\\bullet$ Our observations yield an upper limit to the HI mass of the galaxy of approximately  $<$~10$^{7}$~$\\mathrm{M_\\Sun}$. \\\\ \\\\ $\\bullet$ Assuming a minimum radius equal to that of the CO disk radius, we obtain an HI surface density of $\\leq$ 1.5~$\\mathrm{M_\\Sun}$~pc$^{-2}$. \\\\ \\\\ $\\bullet$ The ratio of the molecular gas mass to atomic gas mass in the galaxy is $>$ 7, which puts NGC 4476 among the most seriously HI deficient (compared to H$_2$) cluster  galaxies. \\\\\\\\ $\\bullet$ An analysis of these new HI data suggest that the HI in NGC 4476 may have been removed due to ram-pressure stripping as the galaxy traveled on a highly radial orbit through the Virgo Cluster core. It is puzzling that there is no HI left in the inner disk, unless the physical conditions in the molecular ISM of NGC 4476 are very different (about a factor of 10 higher average density, coupled with lower UV fields)  from those of typical Virgo spirals."
    },
    "0410/astro-ph0410338_arXiv.txt": {
        "abstract": "We present a double distribution function of dark matter halos, with respect to both object mass and local over- (or under-) density. This analytical tool provides a statistical treatment of the properties of matter {\\em surrounding} collapsed objects, and can be used to study environmental effects on hierarchical structure formation. The size of the ``local environment'' of a collapsed object is defined to depend on the mass of the object. The Press-Schechter mass function is recovered by integration of our double distribution over the density contrast. We also present a detailed treatment of the evolution of overdensities {\\em and underdensities} in $\\Omega_{\\rm m}+\\Omega_\\Lambda=1$ and $\\Omega_{\\rm m}=1$ universes according to the spherical evolution model. We explicitly distinguish between true and linearly extrapolated overdensities and provide conversion relations between the two quantities. ",
        "introduction": "Cosmological distributions have long been used with great success as an analytical tool complementary to numerical simulations. They have been used to constrain the cosmological parameters; interpret results of cosmological simulations; study regions of the parameter space which cannot be approached by simulations due to prohibitive computational cost; exploring the effects of various physical processes in an efficient if approximate way. The analytical tool used most widely in cosmology is the mass function of dark matter halos (distribution of the number density of halos with respect to halo mass). Analytical descriptions of dark matter halos are usually based on the Press-Schechter formalism (\\cite{ps},\\cite{b91}) and its extensions (e.g.,\\cite{lc93}, \\cite{ph90}, \\cite{j95}, \\cite{bow91}). The Press-Schechter mass function has is in good agreement with the results of N-body simulations (e.g. \\cite{wef}, \\cite{lc94},  \\cite{ecf}). More sophisticated approaches taking into account deviations from spherical symmetry (e.g., \\cite{ls98}, \\cite{smt01},\\cite{st02}) have improved this agreement even further.  To derive the Press-Schechter mass function, regions in space are smoothed on successively smaller scales. The mass of a collapsed object is then taken to be the largest smoothing mass scale for which the average linear overdensity exceeds some threshold. In this way, matter in the universe is distributed among collapsed structures of different masses, which all share the same value of average overdensity (the threshold value). Information about the local environment of the collapsed objects (whether they live in underdensities or overdensities) is thus erased.  For this reason, and despite its wide applicability, the expression for the mass function cannot be used to address environment-related questions: Does the mass function of structures in superclusters differ from the mass function inside  voids? Are structures of some particular mass more likely to reside inside underdense or overdense regions in space? How does such a preference evolve with redshift, and  how sensitively does it depend on the cosmological parameters? How does the state of the material surrounding and accreted by a collapsed object depend on the mass of the object and the cosmic epoch?  To address such questions, we seek a {\\em double distribution} of the number density of structures with respect to mass but also to local overdensity (or underdensity). In order to extract information about the surroundings of collapsed structures, we use the same random walk formalism  which rigorously yields the Press-Schechter mass function (\\cite{b91}, \\cite{lc93}). Integration of this distribution over density contrast should return the Press-Schechter mass function so that the successes of the Press-Schechter formalism be retained.  Two complications arise in the effort to expand the Press-Schechter mass function to incorporate a description of the local overdensity. First, the concept of the ``local environment'' is somewhat vague and needs to be defined in a more rigorous way. The size of the ``local environment'' cannot be the same for all structures. If this was the case, very small structures would represent only a tiny fraction of the ``environment'', while very large structures could even exceed the size of the ``environment'', which would be an unphysical situation. This problem is not exclusive to analytical tools, but also needs to be addressed when analyzing the results of numerical simulations. Second, the Press-Schechter treatment of the density field uses linear theory, and ways of converting this information to a more physical non-linear result need to be determined.  We address the first problem by introducing a {\\em clustering scale parameter}, $\\beta$, which allows us to define the size of the ``environment'' of each structure as a function of its mass. We address the second concern by calculating conversion relations between the linear-theory overdensities (or underdensities) and those predicted by the spherical-evolution model.  Environment-related questions in cosmological structure formation have also beed addressed using analytical models for the clustering properties of dark halos which evaluate quantities such as the cross-correlation function between dark halos and matter, and the biasing factor (e.g. \\cite{mw96}, \\cite{kns97}, \\cite{s98}, \\cite{tp98}, \\cite{lk99}, \\cite{st99}, \\cite{sj00}, also see review by \\cite{cs02} and references therein). These analyses are based on  the same random walk formalism which we use here to derive our double distribution (also see \\cite{j98} for fitting formulae from N-body simulation results). However, the information content of the double distribution, which treats the ``environment'' in a mass-dependent fashion, is complementary to that of correlation functions, which describe the clustering properties of the dark halo population at some {\\em fixed spatial scale}. The double distribution is ideally fitted for population studies of cosmological objects. If the properties of a single object can be parametrized as a function of its mass and its environment, then the double distribution can be used to predict the statistical properties of such objects, as well as their evolution with time, for any cosmological model.  Our paper is organized as follows. In \\S \\ref{form} we briefly review the ``random walk'' formalism used in the derivation of the Press-Schechter mass function by \\cite{b91} and \\cite{lc93}. We then use the same formalism to derive the double distribution for dark matter halos. We also discuss converting between linear and spherical-evolution density contrasts, and we present some interesting derivative quantities of the double distribution. In \\S \\ref{res} we explore the information content of our double-distribution by plotting the distribution itself as well as its derivative quantities for concordance ($\\Omega_{\\rm m}+\\Omega_\\Lambda =1$) and Einstein-deSitter universes. Finally, we conclude and discuss our findings in \\S \\ref{disc}. ",
        "conclusions": "\\label{disc}  We have presented an extension of the Press-Schechter mass function, in the form of double distribution of structures with respect to mass and local overdensity. We have done so by introducing a clustering scale parameter $\\beta > 1$, which we use to associate with each collapsed object of mass $m$ a larger environment of mass $\\beta m$.  The scale parameter $\\beta$ can be expressed as a fucntion of the number of virial radii included in the local environment of each structure. We found that for reasonable values $\\beta \\sim 2$, the shape of the distribution does not depend sensitively on this parameter. Integration over linearly extrapolated overdensity returns the original Press-Schechter mass function, independently of the value of $\\beta$.  We present the double distribution in terms of the true, physical, nonlinear density contrast $\\delta$. However, in calculating the distribution it is useful to identify regions using instead the overdensity obtained via linear analysis, $\\ed$, extrapolated to the present epoch. A useful fitting function is given for the $\\delta-\\ed$ conversion.   The double distribution is useful because it allows us to have an explicit analytical if approximate description of the environment in which collapsed objects of all masses reside. Using the tools we have developed, it can be readily calculated for any flat cosmology, and evaluated at any epoch. Consequently, it offers new insight into the growth of structure as well as the present distribution of collapsed objects.  We have evaluated the double distribution and some of its integral moments for both a concordance cosmology and an Einstein-de Sitter universe. Some key results are that at any redshift, the double distribution is dominated by a peak which shifts in mass but is always at a relatively low value of $|\\delta|$. For each mass, there is a most probable $\\delta$, which increases with structure mass. Moreover, at the present epoch in the concordance universe, the most probable environment is a modest {\\em under}density, for all objects below about $10^{14} M_\\odot$; thus, underdensities are preferentially populated by low-mass objects. Finally, the fraction of mass in underdensities increases with time, and in the concordance cosmology the present underdense mass fraction in objects of $M>10^{12} {\\,\\rm M_\\odot}$ is about 40\\%. These trends can be understood in terms of hierarchical clustering in which overdense regions are the site of vigorous merging that clears out low-mass objects, which then find their last refuge in voids.  These results are consistent with other analyses in the literature which use Press-Schechter-like formalism to probe the correlations between structures. For example, \\cite{mw96} extend the Press-Schechter random-walk picture in a way very similar to ours, but identify a local environment in terms of a fixed, mass-independent radius.  Our results are thus complementary to these, because we adopt a mass-dependent environment based on the virial radius.   Finally, it is well-known that the Press-Schechter mass function provides an excellent characterization of the results of numerical $N$-body simulations (e.g., \\cite{lc94}). It would be of great interest to compare the analytic double distribution we have presented with numerical results. We plan to address this issue in future work."
    },
    "0410/astro-ph0410612_arXiv.txt": {
        "abstract": "This work investigates the use of a well-known empirical correlation between the velocity dispersion, metallicity, and luminosity in H$\\beta$ of nearby H$\\rm{II}$ galaxies to measure the distances to H$\\rm{II}$-like starburst galaxies at high redshifts. This correlation is applied to a sample of 15 starburst galaxies with redshifts between $z=2.17$ and $z=3.39$ to constrain $\\Omega_m$, using data available from the literature.  A best-fit value of $\\Omega_m=0.21^{+0.30}_{-0.12}$ in a $\\Lambda$-dominated universe and of $\\Omega_m=0.11^{+0.37}_{-0.19}$ in an open universe is obtained.  A detailed analysis of systematic errors, their causes, and their effects on the values derived for the distance moduli and $\\Omega_m$ is carried out.  A discussion of how future work will improve constraints on $\\Omega_m$ by reducing the errors is also presented. ",
        "introduction": "Precision cosmology, or accurately constraining the parameters describing the universe, has recently become an active field of research due to the precision of available data sets.  Stringent contraints have recently been placed on cosmological parameters from measurements of the microwave background \\citep{B03}, type Ia supernovae \\citep{R00}, and galaxy surveys \\citep{H03,Ba03}.  Although these sources of data are sufficient for generating consistent values for the mass density ($\\Omega_m$), vacuum energy density ($\\Omega_\\Lambda$), the dark energy equation of state parameter ($w$), and the value of spatial curvature in the universe ($\\Omega_k$), these values must be checked via as many independent methods as possible for consistency, accuracy, and avoiding systematic biases.  Furthermore, without the data from supernovae, there would be weak evidence at best for stating that $w \\approx -1$, thus it is important to seek another, independent observation supporting the existence of dark energy.  The cosmological parameter with the greatest number of observable cross-checks is $\\Omega_m$.  It has been derived using many techniques, including the Sunyaev-Zel'dovich effect \\citep{G01}, weak gravitational lensing \\citep{H01}, X-ray luminosities \\citep{B01}, large scale clustering \\citep{SBCG03}, peculiar velocities of galaxy pairs \\citep{F03}, and supernovae data \\citep{K03}.  These methods yield results ranging from $\\Omega_m = 0.13$ to $\\Omega_m = 0.35$, and are all consistent with one another at the 2-$\\sigma$ level.  However, they all face difficulties when attempting to differentiate between cosmological models, as they are only weakly dependent on $\\Omega_\\Lambda$, $\\Omega_k$, and $w$.  If a reliable standard candle were found at high redshifts, cosmological models could be discriminated between by precise and accurate observations, as the distance modulus becomes sensitive to $\\Omega_\\Lambda$, $\\Omega_k$, and $w$ at higher redshifts.  It is known that local H$\\rm{II}$ galaxies and giant H$\\rm{II}$ regions in local galaxies are physically similar systems \\citep{M87}.  This work extrapolates a link between nearby H$\\rm{II}$ galaxies and H$\\rm{II}$-like starburst galaxies at high redshifts to use such objects as standard candles.  This is accomplished through the application of the known correlation between the luminosity in the $ H\\beta $ line ($L_{H\\beta}$), the velocity dispersion ($\\sigma$), and metallicity ($O/H$) of nearby H$\\rm{II}$ galaxies \\citep[discovered by][hereafter MTM]{MTM88} to the H$\\rm{II}$-like starburst galaxies found at high redshifts.  This correlation, when applied to starburst galaxies at $z > 2$, allows for discrimination between different values of $\\Omega_m$ \\citep[as first suggested by][hereafter MTT]{MTT00}, and can discern which cosmological model is most favored by the data.  H$\\rm{II}$ galaxies (and H$\\rm{II}$ regions of galaxies) are characterized by a large star-forming region surrounded by singly ionized hydrogen.  The presence of O- and B-type stars in an H$\\rm{II}$ region causes strong Balmer emission lines in $H\\alpha$ and $H\\beta$.  The size of a giant H$\\rm{II}$ region was shown to be correlated with its emission line widths by \\citet{M78}.  This correlation was improved upon by \\citet{TM81}, who showed that $L_{H\\beta}$ of giant H$\\rm{II}$ regions is strongly correlated with their $\\sigma$.  This basic correlation, its extension to H$\\rm{II}$ galaxies, and its usefulness as a distance indicator have been explored in the past \\citep[MTM;][]{MTM87,M87}.  The empirical correlation for H$\\rm{II}$ galaxies (MTM) relates their $L_{H\\beta}$, $\\sigma$, and $O/H$.  The relationship is \\begin{equation} \\log{L_{H\\beta}}=\\log{M_z} + 29.60 \\mathrm{,} \\qquad M_z \\equiv \\frac{\\sigma^5}{O/H} \\mathrm{,} \\end{equation} where the constant 29.60 is determined by a zero-point calibration of nearby giant H$\\rm{II}$ regions (MTT) and from a choice of the Hubble parameter, $H_0 = 71 \\, km \\, s^{-1} \\, Mpc^{-1}$ \\citep{H02,F01}.  The 1-$\\sigma$ rms scatter about this correlation is $\\pm 0.33 \\, \\, \\rm{dex}$ on $\\log{L_{H\\beta}}$ from the local sample of H$\\rm{II}$ galaxies found in MTM.  Starburst galaxies observed at high redshifts exhibit the same strong Balmer emission lines and intense star formation properties \\citep{P01,E03} as nearby H$\\rm{II}$ galaxies.  This work follows the suggestion of MTT that equation 1 applies to the H$\\rm{II}$-like starburst galaxies found at high redshifts, and provides evidence to validate that assumption.  The remainder of this paper discusses the constraints that can be placed on $\\Omega_m$ and the restrictions that can be placed on the choice of cosmology using starburst galaxies.  Section 2 details how the data set was selected and analyzed.  Section 3 states the results obtained from the selected data.  Section 4 provides a discussion of the random and systematic errors associated with this method, including a detailed discourse on the assumption of universality between local H$\\rm{II}$ galaxies and high redshift starburst galaxies.  Finally, section 5 presents the conclusions drawn from this work, and points towards useful directions for future work on this topic. ",
        "conclusions": "This paper has demonstrated that using H$\\rm{II}$-like starburst galaxies at high redshift as a standard candle is a promising and well-motivated avenue to explore for precision cosmology.  A future survey of high redshift starburst galaxies with measurements of $z$, $\\sigma$, $A_{H\\beta}$, $F_{H\\beta}$, $O/H$, and $EW$ will reduce both random and systematic errors dramatically.  Since the inherent scatter of the method is large, a large sample size is required to obtain meaningful constraints.  This paper contains a sample size of only 15 galaxies, but future surveys should be able to obtain hundreds of starburst galaxies that survive the selection cuts.  For a sample of 500 galaxies, this will improve constraints on $\\Omega_m$ to a restriction of $\\pm 0.03$ due to random errors.  Additionally, all of the systematics specific to this sample due to incomplete data will disappear.%  If the assumption of universality between local H$\\rm{II}$ galaxies and high redshift starburst galaxies is correct, this method of measuring $\\Omega_m$ is capable of providing very tight constraints, independent of any constraints arising from other sources, including CMB and SNIa data.  % Additionally, if galaxies are obtained at a variety of redshifts between $2 \\leq z \\leq 4$, different cosmological models (including vacuum-energy dominated models with different values of $w$) can be tested for consistency with the future data set.  If $\\Omega_m \\leq 0.3$, the differences in $DM$ at various redshifts become quite pronounced, and meaningful results as to the composition of the non-matter components of the universe can be obtained as well.  % Future work on this topic has the potential to provide strong independent evidence either supporting or contradicting the concordance cosmological model of $\\Omega_m + \\Omega_\\Lambda = 1$, $w=-1$, in addition to providing a very stringent constraint on the $\\Omega_m + \\Omega_\\Lambda$ parameter space."
    },
    "0410/astro-ph0410048_arXiv.txt": {
        "abstract": "In a previous paper, we described a new method for including detailed information about substructure in semi-analytic models of dark matter halo formation based on merger trees. In this paper, we present the basic predictions of our full model of halo formation. We first describe the overall properties of substructure in galaxy, group or cluster haloes at the present day. We then discuss the evolution of substructure, and the effect of the mass accretion history of an individual halo on the mass function and orbital grouping of its subhalo population. We show, in particular, that the shape of the subhalo mass function is strongly correlated with the formation epoch of the halo. In a third paper in this series we will compare the results of our semi-analytic method with the results of self-consistent numerical simulations of halo formation. ",
        "introduction": "\\label{sec:1}  Following the analysis of the first-year data from the Wilkinson Microwave Anisotropy Probe (WMAP), there is now very strong evidence that the matter content of the Universe is in large part non-baryonic (Spergel et al.\\ 2003). Observations of large scale structure (e.g.\\ Percival et al.\\ 2001; Maller et al.\\ 2003; Tegmark et al.\\ 2004), measurements of weak lensing (e.g.\\ Hoekstra et al.\\ 2002; Van Waerbeke et al.\\ 2002; Jarvis et al.\\ 2003; Hamana et al.\\ 2003; Brown et al.\\ 2003; Rhodes et al.\\ 2004), and modelling of the Lyman-$\\alpha$ forest (e.g.\\ Croft et al.\\ 2002; Demia{\\' n}ski \\& Doroshkevich 2003; Kim et al.\\ 2004) indicate a power spectrum of density fluctuations in this component continuing to subgalactic scales, and consistent with `cold' dark matter (CDM) models. The implications of the CDM power spectrum for structure formation are well established. Dark matter haloes, the dense regions that surround galaxies, groups and clusters, form from the bottom up, through the merging of progressively larger structures. This process of hierarchical merging has been studied extensively for two decades, and the overall properties of galaxy or cluster haloes formed in this way are now fairly well determined.  To learn more about dark matter, and to search for features in the power spectrum that could reveal new phases in the evolution of the very early universe, we must push the theory of structure formation to smaller scales. Since the present-day evolution of structure is already nonlinear on the scale of galaxy clusters, it is challenging to make even the simplest predictions on subgalactic scales. The usual approach to the problem is to combine the results of simulations on different scales, using numerical simulations of large (100\\,Mpc$^{3}$ or more) volumes to determine the formation rates and properties of the overall population of dark matter haloes at low resolution (e.g.\\ Kauffmann et al.\\ 1999; Bullock et al.\\ 2001; Gottl{\\\" o}ber, Klypin, \\& Kravtsov 2001; Yoshida, Sheth, \\& Diaferio 2001; Evrard et al.\\ 2002; Mathis et al.\\ 2002;  Wambsganss, Bode \\& Ostriker 2003; Lin, Jing, \\& Lin 2003; Gottl{\\\" o}ber et al.\\ 2003; Zhao et al.\\ 2003a, 2003b; Hatton et al 2003; Reed et al.\\ 2003; Klypin et al.\\ 2003; Percival et al.\\ 2003; Tasitsiomi et al.\\ 2004; Yahagi, Nagashima, \\& Yoshii 2004; Gao et al.\\ 2004b), more detailed re-simulations of objects selected from these volumes to determine the structure of individual haloes for reasonably large samples (e.g.\\ Klypin et al.\\ 2001; Jing \\& Suto 2002; Fukushige \\& Makino 2001, 2003; Ascasibar et al.\\ 2003; Hayashi et al.\\ 2003; Fukushige, Kawai \\& Makino 2004; Hoeft, M{\\\" u}cket, \\& Gottl{\\\" o}ber 2004; Navarro et al.\\ 2004; Tasitsiomi et al.\\ 2004), and finally very high-resolution simulations of individual haloes focusing on their structure and substructure, either on cluster scales (e.g.\\ Tormen 1997; Tormen, Diaferio \\& Syer 1998; Ghigna et al.\\ 1998; Col{\\'{\\i}}n et al.\\ 1999; Klypin 1999a; Okamoto \\& Habe 1999; Ghigna et al.\\ 2000; Springel et al.\\ 2001; Governato, Ghigna \\& Moore 2001; De Lucia et al.\\ 2004; Gill, Knebe, \\& Gibson 2004a; Gill et al.\\ 2004b; Gao et al.\\ 2004a), on galaxy scales (e.g.\\ Klypin et al.\\ 1999b; Moore et al.\\ 1999a, 1999b; Stoehr et al.\\ 2002; Power et al.\\ 2003), or both (e.g.\\ Diemand et al.\\ 2004c; Gao et al.\\ 2004b; Weller, Ostriker \\& Bode 2004; Reed et al.\\ 2004).  This approach has provided a detailed picture of the structure and dynamics of dark-matter-dominated systems over a wide range of scales, from tens of kiloparsecs to hundreds of Megaparsecs. There is a hard limit to the dynamic range that can be achieved using a multi-scale approach, however. Structure formation mixes information on many different scales as haloes form. To model the formation of a halo accurately, one needs to include the effects of very long-wavelength fluctuations together with the smaller fluctuations that produce substructure. Thus the minimum scale that can be included in any self-consistent simulation of a present-day halo is constrained by the requirement that the total volume studied still be in the linear regime (where the effect of larger fluctuations can be accounted for) at the present day, and by the finite numerical resolution available computationally. For the highest-resolution simulations possible currently, this leads to a minimum mass scale for resolved substructure of around 10$^{-4}$--10$^{-5}$ of the mass of the main halo considered. To study halo substructure below this mass limit requires analytic or semi-analytic extensions to the numerical results. It is precisely this sort of small-scale information, however, that is required in many current applications including galaxy dynamics, strong lensing, direct or indirect dark matter detection or tests of dark matter physics in general.  In earlier work (Taylor \\& Babul 2001, TB01 hereafter), we developed a model for dynamical evolution of satellites orbiting in the potential of larger system. This model includes dynamical friction, tidal mass loss and tidal disruption as its main components. It calculates satellite evolution over a many short timesteps, rather like a restricted $N$-body simulation, but uses only global properties of the satellite to determine its evolution, thus reducing the computational expense considerably. More recently (Taylor \\& Babul 2004a, paper I hereafter), we applied this model of satellite evolution to the merging subcomponents involved in the hierarchical formation of galaxy, group or cluster haloes. As one of the main results of paper I, we established a non-parametric way to include realistic amounts of substructure in semi-analytic merger trees without increasing the computational cost substantially. This `pruning' method, together with established merger tree methods and the analytic model of satellite dynamics from TB01, constitute a new semi-analytic model for the formation of individual dark matter haloes.  In this paper, we discuss the basic predictions of this model. In section \\ref{sec:2}, we first review the model, which was described fully in paper I. In section \\ref{sec:3}, we then discuss the distribution functions for the mass, peak  circular velocity, and spatial location of subhaloes that it predicts. In section \\ref{sec:4}, we show that these results are relatively insensitive to specific model parameters, uncertainties or assumptions. The model predicts strong correlations among the different properties of subhaloes, and between the age of subhaloes and their other properties, which are discussed in section \\ref{sec:5}. In section \\ref{sec:6}, we show how the systematic dependence of subhalo properties on age may allow the use of substructure as an indicator of the dynamical history of an individual system. Finally in section \\ref{sec:7}, we consider kinematic grouping of substructure due to the merger sequence, and the effects of encounters between subhaloes. We discuss our results in section \\ref{sec:8}.  In a subsequent paper (Taylor \\& Babul 2004b, paper III hereafter) we will compare the predictions of this semi-analytic model with the results of self-consistent numerical simulations of halo formation. This comparison is  particularly interesting, since the only free parameters in the semi-analytic  model were fixed in paper I, either by matching restricted simulations of  individual subhaloes (to fix the parameters of the dynamical model), or by  assuming self-similarity in the merging process (to fix the one parameter in the pruning method). Thus we have no remaining parametric freedom when comparing our results to self-consistent simulations, making the comparison a meaningful one.  Finally, we note that as in paper I, here and in paper III we will generally consider results for the former `standard' CDM (SCDM) cosmology with $h = 0.5$ and $\\sigma_8 = 0.7$, because the simulations we compare to directly in paper III assumed this cosmology. As illustrated in section \\ref{subsec:4.2}, the properties of substructure depend only weakly on cosmology, and thus we expect our main results to be equally valid for $\\Lambda$CDM or similar cosmologies. We will consider the cosmological dependence of our results in more detail in future work. ",
        "conclusions": "\\label{sec:8}  In this paper, we have discussed the basic properties of halo substructure, as predicted by the semi-analytic model of halo formation described in paper I. This model includes several distinct components: \\begin{itemize} \\item Halo merger histories are generated randomly, using the merger-tree algorithm of Somerville and Kolatt (1999).  \\item The higher order branchings in these trees are `pruned', using the method described in paper I, to determine whether each branch merging with the main trunk contributes a single subhalo or a group of associated subhaloes to the main system. This produces a single list of subhaloes merging with the main system at various redshifts.  \\item Each subhalo from this final list is the placed on a random orbit starting at the virial radius of the main system, and evolved using the analytic model of satellite dynamics described in TB01, experiencing orbital decay due to dynamical friction, and heating and stripping due to tidal forces.  \\item Haloes which were associated with a given parent before pruning fall in together with the parent on similar orbits, as part of a kinematic group.  \\item The properties of the main system change over time, its mass growing according to the merger tree and its concentration changing according to the relations in ENS01.  \\item No baryonic component is included in the models presented here, although one can easily be added, given a prescription for gas cooling and star formation. \\end{itemize}  Overall accuracy and parametric dependence are important considerations for any semi-analytic model. From paper I, we expect our dynamical model of satellite evolution to predict the properties of individual subhaloes to an accuracy of roughly 20 percent until they have lost a large fraction ($\\simgt$ 80--90 percent) of their original mass. This component of the full semi-analytic model has three principal free parameters: the (halo) Coulomb logarithm, $\\Lambda_{\\rm h}$, which determines the strength of dynamical friction, the heating coefficient $\\epsilon_{\\rm h}$, which determines the mass loss rate, and the disruption criterion $f_{\\rm crit}$, which determines how long stripped systems continue to survive as bound objects. The values of these parameters were fixed by comparison with the simulations of Velazquez and White (1999) and H03, as described in TB01 and in paper I.  Examining the dependence of our results on the values of these parameters, (in section \\ref{sec:4}), we conclude that $\\Lambda_{\\rm h}$ is probably well-enough determined for most applications. The heating parameter $\\epsilon_{\\rm h}$ may introduce larger uncertainties, as it was determined from a limited number of simulations with $10^3$--$10^4$ particles each. As a result it may be that we {\\it over}estimate the importance of heating somewhat, and subhalo masses are $\\sim$ 20--30 percent larger than we predict. The disruption criterion is also slightly uncertain, although it only affects the survival of highly stripped subhaloes. Here again, we are likely to overestimate the disruption rate, as we are basing our results on simulations of finite resolution.  The rest of the semi-analytic model has no other free parameters, although there are a number of other minor uncertainties, assumptions and approximations in the model. The main assumptions concern the form of the density profile for the main halo, the dependence of halo concentration on mass and redshift, and the distribution of orbital parameters when haloes first merge. The average halo density profile is probably well-enough determined from simulations for most of the applications considered here, as is the initial distribution of satellite orbits. The halo concentrations used here (from ENS01) may be somewhat too small for low-mass haloes (cf.\\ Bullock et al.\\ 2001); in this case we are probably underestimating subhalo masses by 10--15 percent.  The model also makes a number of simplifications or approximations, treating the main halo as spherically symmetric, for instance, and ignoring some of the higher-order terms in subhalo evolution, such as collisions between systems or `harassment'. The anisotropy of subhalo orbits has been studied by several authors in recent simulations (e.g.\\ Knebe et al.\\ 2004; Gill et al.\\ 2004b; Aubert, Pichon \\& Colombi 2004; Benson 2004) and is generally fairly small. We do not expect it to affect our results substantially. Similarly, while the density profiles of simulated CDM haloes are flattened or triaxial to a certain degree, the corresponding potentials are only slightly non-spherical, so we do not expect that flattening will change our results on the spatial distribution of satellites substantially. The higher-order terms in subhalo evolution probably increase mass loss and disruption rates slightly, but the estimates of section \\ref{sec:4} suggest this is at most a 10--20 percent effect.  In summary, given the calibration performed in TB01 and paper I, our semi-analytic model has no free parameters, although it includes various assumptions that could be modified, and various simplifications that could be treated in more detail. We expect it to be accurate to roughly 20 percent, the remaining inaccuracy begin due to several competing effects, including possible overestimates of heating and disruption rates, or underestimates of subhalo concentration and mass loss due to encounters. Given this estimated accuracy, it is interesting to compare our results with self-consistent simulations of halo formation.  Overall, the semi-analytic model predicts simple, universal distributions of subhalo mass, circular velocity, and position within the main halo, similar to those seen in simulations. We also predict strong correlations between a subhalo's merger epoch or formation epoch and its properties, and as a result, between the dynamical history of a halo and the average properties of its substructure. Here again, these are qualitatively similar to trends seen in recent studies of substructure in sets of simulated haloes.  In paper III, we will make a more quantitative comparison between numerical and semi-analytic predictions. We have seen preliminary indications of the results of this comparison in section \\ref{sec:5}: subhaloes in simulations generally experience more mass loss and are disrupted faster than in the semi-analytic model. While this disagreement could be due to a systematic error in the semi-analytic model, it could also indicate that the simulations are still subject to important resolution effects. We will discuss this issue in more detail in paper III."
    },
    "0410/astro-ph0410562_arXiv.txt": {
        "abstract": "Three different maps of the quiet Sun, observed with the Advanced Stokes Polarimeter (ASP) and the Diffraction-Limited Stokes Polarimeter (DLSP), show evidence of strong ($\\simeq$1700~G) and weak ($<$500~G) fields coexisting within the resolution element both at network and internetwork locations. The angular resolution of the observations is of 1\\arcsec (ASP) and 0.\\arcsec6 (DLSP), respectively. Even at the higher DLSP resolution, a significant fraction of the network magnetic patches harbor a mixture of strong and weak fields. Internetwork elements that exhibit kG fields when analyzed with a single-component atmosphere are also shown to harbor considerable amounts of weak fields. Only those patches for which a single-component analysis yields weak fields do not show this mixture of field strengths. Finally, there is a larger fractional area of weak fields in the convective upflows than in the downflows. ",
        "introduction": "\\label{sec:intro}  Our understanding of quiet Sun magnetic fields is evolving at a very rapid pace. The classical picture of the quiet Sun is based on a sprinkling of kilo-Gauss (kG) fluxtube-like structures of sub-arcsecond sizes (unresolved in the observations) forming a magnetic network at the boundaries of the supergranular cells. The internetwork (cell interiors) would be almost devoid of flux, exhibiting only sparse weak turbulent flux concentrations. As new observations became available with improved sensitivity and spatial resolution, the internetwork region started to gain importance. Strong kG fields were also found in many internetwork locations and the amount of magnetic flux and energy detected has increased as the instrumentation improved (\\citeNP{KDE+94}; \\citeNP{L95}; \\citeNP{GDKS96}; \\citeNP{LR99}; \\citeNP{SAL00}; \\citeNP{L02}; \\citeNP{SNSA02}; \\citeNP{KCS+03}; \\citeNP{DCKSA03}; \\citeNP{DCSAK03}; \\citeNP{SNMPL04}; \\citeNP{LSN04}).  At the same time, the concept of organized magnetic structures (e.g., fluxtubes) as the building blocks for quiet Sun fields seems to be losing some ground, at least outside of the network. Recent numerical simulations (\\citeNP{C99}; \\citeNP{EC01}; \\citeNP{SAEC03}; \\citeNP{S03}; \\citeNP{SN02}) reveal a more disorganized, almost chaotic, scenario with the field being dragged around by turbulent convective motions. Another important element in this picture is the discrepancy in the distribution of internetwork fields as seen in visible and infrared observations. Infrared data (\\citeNP{L95}; \\citeNP{KCS+03}) show a predominance of sub-kG fields, with a distribution that peaks around 350~G. On the other hand, authors working with visible observations (e.g., \\citeNP{SNSA02}) obtain that most of these fields are of kG strength. \\citeN{SNSA03} proposed that this discrepancy is a natural consequence of unresolved small-scale inhomogeneities of the field. If one has a mixture of weak and strong fields coexisting in the resolution element, visible and infrared observations tend to emphasize different parts of the distribution.  In a recent paper, \\citeN{SN04b} (see also \\citeNP{SN04a}) showed that the visible \\ion{Fe}{1} lines at 6302 \\AA \\, exhibit some sensitivity to unresolved field strength inhomogeneities. In the present work we make use of this property to seek evidence that mixed strengths indeed occur in the quiet Sun. As we discuss below, we are able to detect a rather large number of such mixed strengths. Our results support the ``disorganized'' picture of the quiet Sun fields discussed above, not only in the internetwork but also to some extent in the magnetic network. ",
        "conclusions": "\\label{sec:conc}  This paper reports on observational evidence for mixed field strengths in the quiet Sun ($\\sim$300 and $\\sim$1700~G) on spatial scales smaller than 0.\\arcsec6. Mixed strengths are found in network and internetwork magnetic elements. The results for the internetwork are not entirely unexpected. \\citeN{SNSA03} proposed that such mixture is the most natural explanation for the discrepancy between visible and infrared observations. Our work strongly supports their conjecture and starts to bridge the gap between the two types of observations. The presence of weak fields inside network patches is more surprising and had not been anticipated before.  The reliability of our results is backed by several facts. Our 1C analysis reproduces what had been obtained in the past from visible spectro-polarimetric observations. When we extend the analysis to 2C we find that the merit function $\\chi^2$ is systematically smaller, even though we did not add free parameters in the model. Therefore, our 2C analysis is at least as reliable as comparable 1C studies which have been used thus far. In order to ensure a conservative criterion for the occurrence of 2C, our procedure was implemented with a ``preference'' for 1C solutions when possible. The spatial distributions of the filling factors $\\alpha_w$ and $\\alpha_s$ are smooth, exhibiting spatial coherence even though each pixel has been inverted separately. These distributions, as well as the other results obtained from our study, are consistent in the three maps analyzed (which have been obtained from two different instruments). Several numerical tests (\\S\\ref{sec:reliability}) show that our method is able to distinguish between the filling factors of weak and strong fields in the presence of unresolved PDFs and/or line asymmetries. Finally, our results are sensible from a physical point of view. For example, we find mixed strengths in internetwork locations showing strong 1C fields but not in those showing weak 1C fields. Moreover, the relative fractional area occupied by weak and strong fields is different for upflows and downflows, with the downflows having a larger filling factor of kG fields (in agreement with existing simulations).  It is important to keep in mind that the mixed strengths that we have detected may be just the tip of the iceberg. It is very likely that we only see the most conspicuous ones. We have chosen to parameterize the sub-pixel field distribution by two discrete values at 300 and 1700~G. However, we know from both observations and simulations that quiet Sun fields obey a continuous probability distribution, with the weaker fields covering a larger fraction of the resolution element. Unfortunately, according to \\citeN{SN04b}, we can only infer (at most) two points of such distribution using visible observations. Simultaneous visible and infrared observations will offer a much more detailed picture. It might then be possible to infer more than two points of the distribution."
    },
    "0410/astro-ph0410081_arXiv.txt": {
        "abstract": "October 2004 marks the 400th anniversary of the sighting of SN 1604, now marked by the presence of an expanding nebulosity known as Kepler's supernova remnant. Of the small number of remnants of historical supernovae, Kepler's remnant remains the most enigmatic. The supernova type, and hence the type of star that exploded, is still a matter of debate, and even the distance to the remnant is uncertain by more than a factor of two.  As new and improved multiwavength observations become available, and as the time baseline of observations gets longer, Kepler's supernova remnant is slowly revealing its secrets.  I review recent and current observations of Kepler's supernova remnant and what they indicate about this intriguing object. ",
        "introduction": "Four hundred years ago, when Johannes Kepler and others observed the ``new star'' of 1604, those observing the event had no concept of what it was that they were observing.  Today we know that supernovae are exploding stars and that they even come in different varieties.  The type Ia SNe, which have become so important for cosmology because they are standard candles, arise from the incineration of white dwarf stars.  The other main class, core collapse SNe, come from more massive stars and produce the sub-classes called type Ib, type Ic, and type II.  Of the six historical SNe in our Galaxy that have occurred over the last millennium (including Cas A for which the SN itself apparently escaped detection), only Kepler's SN has remained uncertain as to the type of the star that exploded.  Today of course, what we observe is the expanding young supernova remnant (SNR) that resulted from the explosion.  By studying this SNR across the electromagnetic spectrum, modern astronomers are still trying to discern a clear picture of the precursor star of this event. In doing so, they are hindered by the lack of an accurate distance to the object, which makes the derivation of even basic properties like the diameter or mean expansion velocity uncertain.  In this paper, I will review the current observational status of this enigmatic object and point out some of the apparent inconsistencies in existing interpretations. ",
        "conclusions": ""
    },
    "0410/astro-ph0410225_arXiv.txt": {
        "abstract": "The Square Kilometre Array is conceived as a telescope which will both test fundamental physical laws and transform our current picture of the Universe. However, the scientific challenges outlined in this book are today's problems---will they still be the outstanding problems that will confront astronomers in the period 2020 to 2050 and beyond, when the SKA will be in its most productive years? If history is any example, the excitement of the SKA will not be in the old questions which are answered, but the new questions that will be raised by the new types of observations it will permit.  The SKA is a tool for as-yet-unborn users and there is an onus on its designers to allow for the exploration of the unknown. We outline a philosophy for the design and operation of the SKA that can lead the radio astronomers in the $21^{\\mathrm{st}}$ century to add to the many discoveries of new phenomena made by radio astronomers in the $20^{\\mathrm{th}}$ century. ",
        "introduction": "\\label{sec:eou.intro}  \\begin{quote} \\emph{``Now my own suspicion is that the Universe is not only queerer than we suppose, but queerer than we \\emph{CAN} suppose'': J.~B.~S.~Haldane} \\end{quote}  Most of the phenomena we observe today, using telescopes to observe across the electromagnetic spectrum, were unknown a few decades ago and, to an amazing extent, were discovered by radio astronomers using increasingly powerful instruments and either looking for something else or just following their curiosity. Examples include non-thermal radiation, radio galaxies, quasars, the cosmic microwave background, cosmic evolution, pulsars, gravitational lensing, cosmic masers, molecular clouds, dark matter, and extrasolar planetary systems. These discoveries have changed astronomy in fundamental ways.  Some discoveries resulted from increased sensitivity, others from better spatial or temporal resolution, still others by observing in a new wavelength band or even testing misguided theory.  Most involved recognizing a new phenomenon and being able to distinguish it from a spurious instrumental response.  This scenario is, of course, not restricted to radio astronomy.  Perhaps the most spectacular example from astronomy in other wavebands was the discovery of $\\gamma$-ray bursts by a military satellite---currently a major field of contemporary astrophysics.  It is fashionable to imagine that all research follows some classical model of the scientific method---formulation of a model or hypothesis followed by experimental confirmation. Observations not based on testable theoretical predictions are sometimes called ``butterfly collecting'' or appeals to ``serendipity'' rather than ``real science.''  Time allocation committees, referees of grant applications, and reviewers of instrument proposals tend to focus on specific questions that will be answered.  Yet, astronomy is not an experimental science. We can only observe our Universe and its content with ``eyes'' as wide-open as possible.  We cannot make little changes or experiments to see what happens, except perhaps for some areas of planetary research. So how do we plan for discovery? Despite the apparent capriciousness of our aim, history tells us that a basic requirement is to carry out systematic work with at least an order of magnitude improvement over what has been achieved before in one or more observing capabilities (sensitivity; spatial, temporal, or spectral coverage; spatial, temporal, or spectral resolution). An observing instrument which can offer major advances in several dimensions of parameter space is more likely to make transformational discoveries---history also shows that much greater sensitivity along with flexibility of operation is a wise path to follow. The sensitivity of the Arecibo telescope and the imaging capabilities of array telescopes are excellent paradigms.  Merely providing access to new areas of parameter space with new technology is not a sure-fire recipe for making ground-breaking discoveries, however.  There are other, human, factors to take into account which are just as important for ensuring the SKA's success as a discovery instrument. We are all familiar with what is now the traditional method of using a large common-user telescope involving: i)~a proposal to tackle a single small problem; ii)~review by time allocation committees; iii)~the award of a few hours or maybe even days of observing time; iv)~the analysis of the data via a standard suite of software, and, then if all goes well, v)~a publication following filtering by a referee. We dub this ``the standard model'' of observational astronomy and it is perhaps inevitable that the SKA will allocate much of its operations to this analytic mode.  When, however, the phenomenon or problem is less well-defined, there can be a rich mix of possible ``solutions,'' only some of which may have been explored by theorists and for which the standard model provides a poor response. It is, therefore, vital to develop a philosophy of design, operations, and data archival which allows individuals, small groups, and larger communities freedom to innovate and encourages users to explore completely new ways of collecting, reducing and analyzing data---in other words \\emph{to allow for discovery as well as explanation}. ",
        "conclusions": "\\label{sec:eou.conclude}  The impact of the SKA will depend not so much on the cleverness of the astronomers in defining the important astrophysical problems of today, but in the cleverness and ambition of its designers to obtain better sensitivity, higher dynamic range, larger field-of-view, and perhaps some other parameters not yet contemplated.  The SKA planners and designers must therefore have a vision for a system that is not only much more powerful than previous radio telescopes but at the same time is highly flexible, easy-to-use, and has an operating philosophy which positively encourages and allows the astronomers of tomorrow to look at the sky and to examine their  data in new and creative ways. It must also allow the radio engineers of tomorrow the space to conceive and design the innovative upgrades which increased signal processing and computing power will allow.  Just as important will be to attract and nourish talented and careful observers who can become so familiar with particular types of data as to be able to spot the important new clues which will lead to fresh discoveries. The SKA will need to solve their problems, not our problems, which they will have already solved or shown to be naive or irrelevant. Including ``Exploration of the Unknown'' as one of the prime goals of a highly-flexible SKA is founded firmly on the profound contributions which radio astronomy has already made to our knowledge of the Universe."
    },
    "0410/astro-ph0410539_arXiv.txt": {
        "abstract": "The High  Mass  X-ray Binaries (HMXRB's) SMC X-1 and  4U1700-37 have been observed with FUSE to study the effect of the X-ray source on the stellar wind of the primary. In both systems phase dependent changes in the wind lines have been observed, indicating the creation of a X-ray ionization zone in the stellar wind. The high X-ray luminosity of SMC X-1 ionizes much of the wind and leaves a Stromgren zone. This disrupts the resonance-line acceleration of the wind in portions of the orbit, quencing the wind and disrupting the mass flow. A similar but less dramatic effect was found for the first time in 4U1700-37.  This so-called Hatchett-McCray (HM) effect had been predicted for 4U1700-37, but was not previously detected. ",
        "introduction": "SMC X-1 is a luminous accretion-powered high-mass X-ray binary (HMXRB) with a B0 I primary and a rapid spinning pulsar (0.71 sec) in a 3.89-day period (Wojdowski et al. 1998). FUSE observed SMC X-1 for almost an entire 3.89 day orbital period in 2003 July. A total of 71 exposures (148 ksec) were obtained over 3.4 days, covering phase 0.20, through X-ray source conjunction ($\\phi=0.5$), and primary eclipse ($\\phi=0.93 - 0.07$). The exposures were divided into nine phase bins, each with $t_{expo}=14-20$ ksec and S/N$\\sim15$ per 0.05 \\AA.  The O\\,VI 1032 stellar wind line profile was corrected for Milky Way and SMC interstellar H$_2$ and O\\,VI 1032 absorption by simultaneously modelling the ISM lines in all nine spectra.  The resulting ISM model was subtracted from the observed O\\,VI 1032 line at each phase (see Figure 1a for examples).  The equivalent width  of the O\\,VI line was measured after the ISM correction.  Figure 1b shows the strong variation of the O\\,VI equivalent width with orbital phase. The orbital modulation of O\\,VI in the stellar wind, the HM effect, is a result of ionization of the wind by the X-ray source.  The large X-ray luminosity of SMC X-1 ($L_x =2\\times 10^{38}$ erg s$^{-1}$) ionizes much of the wind and leaves a large Stromgren zone.  This disrupts the resonance-line acceleration of the wind in portions of the orbit, quencing the wind and disrupting the mass flow. The O\\,VI wind absorption is strongly asymmetric around the orbit. The line is at maximum strength during the eclipse of the pulsar $\\phi=$ 0.0. The O\\,VI line virtually disappears near $\\phi \\sim 0.4$. The O\\,VI 1032 column density during eclipse is $\\sim 7\\times 10^{17}$ cm$^{-2}$. The O\\,VI terminal velocity ($\\sim -$700 km/s) drops to near zero at $\\phi=$ 0.3-0.4.  \\begin{figure}[!ht] \\plotfiddle{iping1_fig1.ps}{1.0in}{0}{42}{42}{-240}{-240} \\plotfiddle{iping1_fig2.ps}{1.0in}{0}{45}{45}{-50}{-170} \\caption{({\\it left}) O\\,VI 1032 stellar wind profiles are shown for three orbital phases (0.0, 0.25, and 0.58) of SMC X-1.  The observed profiles are in the lower panel, where the location of Milky Way and SMC interstellar H$_2$ and O\\,VI features are identified.  The upper panel shows the same profiles after subtracting the ISM lines. ({\\it right}) The equivalent width of the O\\,VI 1032 stellar wind line is shown as a function of orbital phase.  The data points are repeated half a cycle before and after $0<\\phi<1$ to illustrate the trend.} \\end{figure} ",
        "conclusions": ""
    },
    "0410/astro-ph0410649_arXiv.txt": {
        "abstract": "{ Using Greenwich sunspot data for 120 years it was recently observed that activity regions on the Sun's surface tend to lie along smoothly changing longitude strips 180\\degr  apart from each other. However, numerical experiments with random input data show that most, if not all, of the observed longitude discrimination can be looked upon as an artifact of the analysis method. ",
        "introduction": "The time distribution of sunspot latitudes is well known and the ubiquitous butterfly diagram is becoming a kind of trademark in solar research. The same is not true for longitudes. It is still not known how persistent in time are different statistical features of longitude distributions for various activity indicators (see for instance Bai~\\cite{Bai03} and references therein). The present letter is inspired by the recent contribution by Berdyugina \\& Usoskin~\\cite{Berd03} (hereafter BU). By using certain data processing techniques they built smooth curves which run along central longitudes  of sunspot activity centres. They observed that distribution of activity centre phases (if computed against mean flow) has a double peak distribution at a high level of confidence (see Fig.5 in BU). They concluded that there is a century scale persistent feature in the sunspot longitude distribution: the local neigbouring maxima of activity tend to lie 180\\degr apart. They also observed several ``flip-flop'' events (see Jetsu et al.~\\cite{Jetsu93}) and concluded that there is a strong analogy between the Sun and rapidly rotating magnetically active late--type stars (for an overview of active stars see Tuominen et al.~\\cite{Tuominen2002}). We can also formulate their result in metaphoric terms: the Sun has two faces, and as a kind of peculiar Janus it sometimes switches them back and forth.  In this paper we analyse {\\em randomly generated} distributions to check the possibility that the observed effect in BU is a result of dependencies and correlations hidden in the statistical method itself. And indeed, we found that it is possible to obtain quite similar double-peak distributions for random data. Consequently the statistically significant part (or even all)  of the effect described in BU can be safely ascribed to the artifacts of data processing. It is still reasonable to assume that the mean magnetic field on the Sun (at least in the regions of spot formation) is axisymmetric, when the mean is taken over sufficently long time.  What follows is a short technical account of the analysis done.  \\subsection{Simulations for active longitude distributions} The strongest argument in BU for the century scale persistence of active longitudes is given in their Fig.~5 where the two strong and symmetric distribution maxima indicate that sunspot groups tend to concentrate around two centres, 180\\degr\\ ($0.5$ in phase) apart.  To check their result and its methodological underpinnings we did some very simple numerical simulations\\footnote{Full source codes of relevant software can be downloaded from {\\it http://www.aai.ee/$\\sim$pelt/soft.htm}.}. Instead of using actual observational data we generated artifical random data sets and tried to process the obtained distributions along the same lines as in the original paper. There was no need to simulate all data processing steps performed in BU. It was enough to start from the point where for each Carrington rotation the phases for activity maxima were computed. For simplicity we assume that for every rotation there are at least two activity maxima, each representing a sunspot group concentration.  Each individual simulation run proceeded then as follows. First, for each of the $N=1720$ Carrington rotations with starting times $t_i, i=1,\\dots,N$ two {\\em fully random and statistically independent} phases $\\phi_1(t_i),\\phi_2(t_i), i=1,\\dots,N$ were generated. The evenly distributed values \\hbox{$0<=\\phi_1(t_i)<1$ and $0<=\\phi_2(t_i)<1$} were then used as simulated phases for principal and secondary maxima of sunspot density distribution along longitude for particular rotation $t_i$.  For both sequences of phases $\\phi_1$ and $\\phi_2$ the half year means were calculated and from the mean values, by linear interpolation, the two contiuous mean curves $M_1(t)$ and $M_2(t)$ were built. To get the analogue of the Fig.5 of BU the distribution of the differences of the all generated phases and one of the continuous curve ($M_1(t)$) was computed. In Fig.~\\ref{fg0000fig} a short fragment of generated phases together with two running means and in Fig.\\ref{sm0000fig} the corresponding distribution for all differences are plotted. \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{fg0000.eps}} \\caption{Fragment of the computer generated phases $\\phi_1(t_i)$ (diamonds) and $\\phi_2(t_i)$ (points). Thick line is a half year average of the principal maxima ($M_1(t)$), thin line is a half year average of the secondary maxima ($M_2(t)$). The mean curves can cross each other.} \\label{fg0000fig} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{sm0000.eps}} \\caption{Distribution of the differences between phases $\\phi_1(t_i), \\phi_2(t_i)$ and the half year mean $M^1(t)$ for one particular run (first from the series of 100).} \\label{sm0000fig} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{sc0000.eps}} \\caption{Mean distribution of the $100$ separate runs. The distribution is flat in the centre.} \\label{sc0000fig} \\end{figure} To eliminate the effect of the statistical fluctuations we generated 100 independent distributions and in Fig.~\\ref{sc0000fig} the mean distribution of these runs are plotted. As we see the means are uncorrelated and the distribution is approximately flat in its central part.  \\subsection{Phase adjustments} In BU, the original phases for active longitude centers were not used straightforwardly, but were modified. The authors say ``We plot the recovered phases of sunspot clusters versus time and find that regions migrate in phase as rigid structures. When a region reaches $\\phi=1$, it appears again near $\\phi=0$. In such cases we add an integer to the phase and unfold continuous migration of the regions''. Unfortunately the desription of the actual procedure of the described processing step remains somewhat obscure. From Usoskin (\\cite{Usoskin03}) we understood that what was done was something similar to the following: \\begin{itemize} \\item There was a preset constant, say $\\delta \\phi$. \\item When particular maximum occured at a phase which exceeded $1-\\delta \\phi$ and when one of the maximum phases for the next rotation was less than $\\delta \\phi$, then such situations were looked upon as a candidate for a phase jump. \\item If the prospective phase jump improved ``continuity'' of the assumed migration paths it was accepted. \\end{itemize} It is important to note that the jumps were not constrained by the type of maxima (main or secondary). To ensure continuity of the migration paths the authors of BU allowed swaps between the principal and secondary maxima (``flip-flops'').  To avoid any subjective judgement we applied a similar but {\\em fully automatic} procedure in our simulation program. Because the actual value for $\\delta \\phi$ (if it ever existed in numerical form) was not communicated to as, we tried multiple values of it. In a ``crossing situation'' we constructed all possible jumping schemes (no jump at all; largest phase jumps; if both can jump, then smaller phase jumps) and selected in each situation a scheme with a smallest change in phases. Thus, if before crossing point, the phases for two maxima were $\\phi_1$ and $\\phi_2$ and after crossing $\\phi_3$ and $\\phi_4$, then the selection criterion was the minimization of the $|\\phi_1-\\phi_3|+|\\phi_2-\\phi_4|$. We believe that the actual procedure of the original authors was carried out similarily, even if it was not implemeted as a rigorous algorithm but was done by manual adjusting.  We computed analogously to Fig.~\\ref{sc0000fig} distribution diagrams for several values of $\\delta \\phi$ in the range $0.05-0.25$. All diagrams had a characteristic bimodal form with two distribution maxima (persistent longitudes!). For presentation purposes we give here an example with one particular parameter value $\\delta \\phi = 1/18$. This is the most natural value for the case when the histogram method of activity smoothing is used and the number of histogram bins is $18$ (see BU). In Fig.~\\ref{fg1100fig} we see a short (approximately one solar cycle) fragment of the two mean curves with upward drifting phases. \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{fg1100.eps}} \\caption{Two rows of random phases with phase adjustments applied. The upward trend is due to the accumulation of the added full cycles.} \\label{fg1100fig} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{sm1100.eps}} \\caption{Distribution of the differences between phases and the mean $M_1(t)$  after adjustment. Crossing parameter $\\delta \\phi = 1/18$. The bimodal structure of the distribution demonstrates that century scale persistence of the active longitudes appear even in randomly distributed phases.} \\label{sm1100fig} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{sc1100.eps}} \\caption{Mean distribution of the differences for $100$ runs. Crossing parameter $\\delta \\phi = 1/18$. The bimodality in Fig.~\\ref{sm1100fig} is not an exceptional fluctuation.} \\label{sc1100fig} \\end{figure} In Fig.~\\ref{sm1100fig} the one concrete distribution (effectively the first from series of 100) is given. The Fig.~\\ref{sc1100fig} depicts the mean distribution of all runs. As we see, the distributions are of two-peaked nature.  We started from random and uncorrelated pairs of phases. How it is then possible that there are now certain preferred positions for the random phases? It is not hard to see that the selection of phase jumping points tends to introduce extra statistical dependencies between neigbouring maxima in different rotations. Some of the extra large differences $\\phi(t_i)-\\phi(t_{i+1})$ can be excluded by doing jumps for appropriate phases. As a result overall mean of differences between sequential phases tends to be less. When these extra correlations and allowed swaps between principal and secondary minima are combined the double-peaked distribution results.  \\subsection{Restrictions from continuity of distributions} For further discussion we need to introduce the notion of circular distance $D(\\phi_1,\\phi_2)$. It is defined as $\\min(|\\phi_1-\\phi_2|,1-|\\phi_1-\\phi_2|)$ and it takes into account the circular nature of the phases (the point $1$ is the starting point $0$ for the next turn).  Our simulation model has been unrealistic up to this point. It is possible that both randomly generated phases $\\phi_1(t_i)$ and $\\phi_2(t_i)$ are equal or very near to each other. For real data this can not be so. Depending on the smoothing and maximum finding method, there is a certain minimal circular distance $\\Delta \\phi$ between two phases of maxima. To simulate this kind of more realistic data we generated random phases as was done before. However for each phase pair we computed the circular distance $D(\\phi_1(t_i),\\phi_2(t_i))$ and pairs whose phases were too close ($D(\\phi_1(t_i),\\phi_2(t_i))<=\\Delta \\phi$) were excluded from the resulting data set. In this way we allow certain statistical dependencies between two phases but phase pairs remain totally independent (from rotation to rotation).  A short fragment of data which is checked against minimal $\\Delta \\phi$ and with adjusted phases is given in Fig.~\\ref{fg1110fig} (crossing parameter $\\delta \\phi = 1/18$, distance parameter $\\Delta \\phi = 0.3$). The tendency of phases to cluster strongly around mean curves is well demonstrated in Fig.~\\ref{sm1110fig}. The plot of the mean distribution for 100 separate runs (Fig.~\\ref{sc1110fig}) shows that we are not dealing here with random fluctuations. It is possible to make the ``persistency'' effect even stronger if we select certain ``optimal'' input parameters $\\delta \\phi$ and $\\Delta \\phi$. \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{fg1110.eps}} \\caption{Fragment of random phases for constrained data. The series of maxima were generated by checking against minimum allowed distance between main and secondary maximum. Crossing parameter $\\delta \\phi = 1/18$ and distance parameter $\\Delta \\phi = 0.30$. The adjustments to follow phase drift were also applied. } \\label{fg1110fig} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{sm1110.eps}} \\caption{ Phase differences from mean $M_1(t)$ for constrained data. } \\label{sm1110fig} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{sc1110.eps}} \\caption{Mean distribution of the phases for distributions like in Fig.~\\ref{sm1110fig}. } \\label{sc1110fig} \\end{figure} \\subsection{Local correlations} If we return now to the original setup of the generation mechanism for the random data sets, we recall that the distributions for all Carrington rotations were assumed to be statistically independent. This is not of course the case for real distributions. It is well known that there are somewhat persistent activity regions on the Sun. If we try to model this kind of local correlation, the separation of the two maxima becomes even more pronounced.  \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{fg1111.eps}} \\caption{Two rows of random phases and corresponding mean curves for locally correlated data. Crossing parameter $\\delta \\phi = 1/18$, distance parameter $\\Delta \\phi = 0.3$. In this kind of simulation even ``flip-flop'' events can be seen. } \\label{fg1111fig} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{sm1111.eps}} \\caption{ Phase differences from mean $M_1(t)$ for locally correlated data.} \\label{sm1111fig} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{3.3 cm}{\\includegraphics{sc1111.eps}} \\caption{Mean distribution of the phase differences for distributions like in Fig.~\\ref{sm1111fig}. } \\label{sc1111fig} \\end{figure}  Look at the Figures \\ref{fg1111fig} through \\ref{sc1111fig}  which were obtained using locally correlated data sets. The two-peaked nature is now very pronounced and corresponding plots tend to be quite similar to the plots in the original paper. By fine tuning of the parameters it is possible to get a perfect match. But this is not our goal here.  The local correlations were introduced by using the following generation scheme. For every new phase to be generated, one from the four possible actions was selected (with equal probability): \\begin{itemize} \\item generate totally new phase, \\item repeat previous phase, \\item add to previous phase a small amount $\\delta$, \\item subtract from previous phase a small amount $\\delta$. \\end{itemize} As a result we got a sequence of globally uncorrelated (due to the first option) but locally correlated random walk (due to the other options) fragments. We are not arguing here that this is a very good model for a solar activity, our aim is just to show that distributions like these in Fig.~5 of BU can be obtained from locally correlated random sequences and that the correlatedness can strengthen distribution peaks of preferred longitudes. ",
        "conclusions": "The analysis above shows that phase distributions of various parameters (maxima in our case) tend to be very sensitive to hidden dependencies and correlations. It is not very easy to see from the first glance that if we look at random primary and secondary maxima, then a phase adjusting to follow ``rigid patterns'' introduces a significant hidden statistical dependencies which show up as double peaked distribution of phase differences.  When we combine all three types of dependencies and correlations, from phase tracking, from constrained phase differences and from real short time correlations, then we easily get significantly bimodal diagrams like Fig.5 in BU. The flexibility due to the possibility to swap principal and secondary maxima (``flip-flop'') and adjust phases (``rigid patterns'') allows one to amplify every local short range correlation into magnificent century scale persistent phenomena.  It is important to note that our simulation analysis does not rule out that in principle there can be certain persistent phenomena in sunspot longitude distributions. We just demonstrated that current evidence is not sufficient."
    },
    "0410/astro-ph0410363_arXiv.txt": {
        "abstract": "The interstellar thick disks of galaxies contain not only gas, but significant quantities of dust.  Most of our knowledge of extraplanar dust in disk galaxies comes from direct broadband optical imaging of these systems, wherein the dust is identified due to the irregular extinction it produces against the thick disk and bulge stars.  This observational technique is sensitive to only the most dense material, and we argue much of the material identified in this way traces a cold phase of the interstellar thick disks in galaxies.  The presence of a cold, dense phase likely implies the interstellar pressures in the thick disks of spiral galaxies can be quite high.  This dense phase of the interstellar medium may also fueling thick disk star formation, and \\halpha\\ observations are now revealing \\HII\\ regions around newly-formed OB stars associations in several galaxies.  We argue that the large quantities of dust and the morphologies of the structures traced by the dust imply that much of the extraplanar material in disk galaxies must have been expelled from the underlying thin disk. ",
        "introduction": "Understanding the structure of the gaseous component of disk galaxies is an important step for unraveling their on-going evolution. In particular, extraplanar gas in galaxies is a potentially-important probe of the effects of kinetic and radiative feedback from massive stars to the gas in galaxies, as well as of the role that the accretion of metal-poor gas plays in the evolution of modern galaxies. Most studies of extraplanar matter have focused on the gas content of the interstellar thick disks and halos of galaxies.  However, the processes that transport matter from the thin disks of galaxies into the thick disks and extended halos will act on both gas and interstellar dust grains.  The presence of extraplanar dust in galaxies can strongly affect the thermal balance of the gas, and it will definitely affect an observer's view of a galaxy through its impact on the transfer of radiation.  The presence of extraplanar dust in the canonical edge-on galaxy NGC 891 was noted many years ago (Sandage 1961; Dettmar 1990; Keppel et al.  1991).  However, an analysis of the implications of the extensive web of dusty extraplanar material in this galaxy did not occur until recently (Howk \\& Savage 1997, 2000).  In the last few years, extraplanar dust has been recognized and studied in many edge-on systems (Sofue, Wakamatsu, \\& Malin 1994; Howk \\& Savage 1999; Alton et al. 2000; Rossa \\& Dettmar 2003; Thompson, Howk, \\& Savage 2004), and we now know that the presence of extraplanar dust and gas are indeed coupled.  In this contribution, I will concentrate on the physical interpretation of the observed extraplanar dust and its implications for the nature of the interstellar thick disks of spiral galaxies.  In particular, I will summarize some of the evidence that material identified with extraplanar dust may represent a dense, cold medium in the thick disks of galaxies. ",
        "conclusions": ""
    },
    "0410/astro-ph0410680_arXiv.txt": {
        "abstract": "\\baselineskip 11pt Dark energy constraints have forced viable alternatives that differ substantially from a cosmological constant $\\Lambda$ to have an equation of state $w$ that {\\it evolves} across the phantom divide set by $\\Lambda$. Naively, crossing this divide makes the dark energy gravitationally unstable, a problem that is typically finessed by unphysically ignoring the perturbations. While this procedure does not affect constraints near the favored cosmological constant model it can artificially enhance the confidence with which alternate models are rejected.  Similar to the general problem of stability for $w< 0$, the solution lies in the internal degrees of freedom in the dark energy sector.  We explicitly show how to construct a two scalar field model that crosses the phantom divide and mimics the single field behavior on either side to substantially better than 1\\% in all observables.  It is representative of models where the internal degrees of freedom  keep the dark energy smooth out to the horizon scale independently of the equation of state. ",
        "introduction": "A self-consistent model for the dark energy requires not only a parameterization of the evolution of its equation of state in the background $w=p/\\rho$ but also a physical model for its spatial fluctuations to guarantee gravitational stability. Cosmological constraints on a constant equation of state have continued to close in upon $w=-1$ (e.g. \\cite{Seletal04}). Since models with $w<-1$ are called phantom dark energy models (e.g. \\cite{Cal02}), we call this the ``phantom divide\".  Viable alternate models for a strongly evolving $w$ therefore must cross the divide at intermediate redshift so that the effects on either side cancel. Simple generalizations of a single scalar field model that cross the divide  in fact cause severe gravitational instabilities in the dark energy sector (e.g. \\cite{Vik04}).  The usual approach in the literature for dealing with such cases is to artificially turn off the dark energy perturbations.    Doing so violates energy-momentum conservation whenever $w \\ne -1$. The justification for dropping these perturbations is that observations already place $w$ close to $-1$ and so the induced error is presumably small in some physical realization of a crossing model. While true for the currently allowed $1\\sigma$ deviations of $\\Delta w \\sim 0.1$,  the confidence level at which larger deviations can be rejected can be affected.   Furthermore, a strong time evolution allows $w$ to differ substantially from $-1$ during some epochs and still be consistent with the distance data. It is therefore important to show explicitly that models exist where the dark energy remains smooth as it crosses the phantom divide as implied by the usual procedure.  The need for a self-consistent treatment of the evolution of the dark energy is most apparent for the cosmic microwave background. Here the ISW effect is sensitive to the decay of the gravitational potential and, for example, the difference between the smooth and clustered regimes of the dark energy for a constant $w\\approx -2/3$ is roughly $50\\%$.     While the impact for a canonical scalar field which is smooth out to the horizon scale is lower, it is well known \\cite{CalDavSte98} that it changes CMB predictions significantly for larger $w$ and hence the confidence level of constraints on highly deviant $w$ (e.g. \\cite{WelLew03}).  In this Brief Report, we explicitly construct a two scalar field model of the dark energy that is gravitationally stable across the phantom divide and matches the canonical single scalar field predictions on either side to much better than $1\\%$ in all observables. Taken literally, such a model of course compounds the coincidence problem of dark energy but as we shall show it is a proxy for a potentially wider class of dark energy models whose internal degrees of freedom keep it smooth out to the horizon scale. ",
        "conclusions": "\\label{sec:discussion}  In Fig.~\\ref{fig:time} we show examples of the performance of the model on the dark energy observables of comoving distance $D(z)$, Hubble parameter $H(z)$ and linear growth rate $G(z)$ in the smooth dark energy regime (see Eqn.~{\\ref{eqn:smooth}) for strong variations in the equation of state $w_a = \\pm 0.5$.   To better show the relevant performance we choose $w_0$ such that $w_n \\approx -1$ and plot the observables relative to a nearly degenerate fiducial $\\Lambda$ model of $\\Omega_{\\rm DE}  = 1-\\Omega_m = 0.74$ and $h=0.735$ which is a good fit to the current CMB data.    In all cases we adjust the $\\Omega_{\\rm DE}$ of the target model to match the comoving angular diameter distance to recombination and hence the CMB peak results. Note that for $w_a>0$, despite a mismatch in $w$ at very high redshift, all observables as a function of redshift are modeled to the $0.1\\%$ level.  The performance for $w_a<0$ is somewhat worse but never exceeds a few times this level.  Thus the two field model will remain an adequate parameterization until the statistical and systematic errors in the dark energy observable measurements reaches the sub $1\\%$ level.  Mismatches at this point merely reflect the unavoidable fact that higher order derivatives in $w(a)$ produce observable effects and our target constant $w_a$ model is itself inadequate.  Even in this regime the 2 field model can be useful since the parameter $\\epsilon_n$ can be used to marginalize or probe the second derivative.  The two field model allows one to also calculate the CMB anisotropy. For a case with no crossing of the divide, e.g. $w_a=0.5$, $w_0=-1$, the single field model with the exact $w(a)$ matches the 2 field model with the approximate $w(a)$ to $\\sim 0.1\\%$, in particular at the low multipoles of the ISW effect (see Fig.~\\ref{fig:clmodel}).  In addition to the adequate matching of $w(a)$, this indistinguishability is a consequence of setting all of the sound speeds to  $c_e=1$. Aside from $w(a)$ the remaining degrees of freedom in a dark energy model involve the transition scale between the smooth and clustered regimes. We also show the predictions for a model that does cross the phantom divide $w_a=0.5$, $w_0= -1.15$.  As this model has an equation of state at the normalization point of $w_n = -1.025$, it predicts even smaller deviations from the fiducial $w=-1$ model with no instabilities in the evolution of the gravitational potential.  The model constructed here permits a self-consistent likelihood analysis of dark energy observables involving both models that cross the phantom divide and those that differ strongly from $\\Lambda$ in the pivot $w_n$. Although predictions for the former class models differ little from the currently favored smooth cosmological constant case, the confidence level at which the latter can be excluded can be affected by a consistent model of dark energy clustering.  A detailed study of the effect on current cosmological constraints is beyond the scope of this Brief Report.  If future observations require an evolutionary crossing of the phantom divide, it will be a good indication that the dark energy contains hidden internal degrees of freedom in its physical structure.  \\smallskip \\noindent{\\it Acknowledgments:}  We thank C. Armendariz-Picon, S. Carroll, A. Lewis and D. Pogosian for useful discussions. WH was supported by the DOE and the Packard Foundation. This work was begun in the Aspen Center for Physics and completed at the KICP under NSF PHY-0114422.  \\noindent{\\it Note added in proof:} After completion of this work, we became aware of related work by Guo et al [astro-ph/0410654]"
    },
    "0410/astro-ph0410155_arXiv.txt": {
        "abstract": "In a recently completed survey of the stellar populations properties of LINERS and LINER/HII Transition Objects (TOs), we have identified a numerous class of galactic nuclei which stand out because of their conspicuous $10^{8-9}$ yr populations, traced by high order Balmer absorption lines and other stellar indices. These objects were called ``Young-TOs'', since they all have TO-like emission line ratios. In this paper we extend this previous work, which concentrated on the nuclear properties, by investigating the radial variations of spectral properties in Low Luminosity Active Galactic Nuclei (LLAGN). Our analysis is based on high signal to noise long-slit spectra in the 3500--5500 \\AA\\ interval for a sample of 47 galaxies. The data probe distances of typically up to 850 pc from the nucleus with a resolution of $\\sim 100$ pc ($\\sim 1^{\\prime\\prime}$) and $S/N \\sim 30$. Stellar population gradients are mapped by the radial profiles of absorption line equivalent widths and continuum colours along the slit. These variations are further analyzed by means of a decomposition of each spectrum in terms of template galaxies representative of very young ($\\le 10^7$ yr), intermediate age ($10^{8-9}$ yr) and old ($10^{10}$ yr) stellar populations.  This study reveals that Young-TOs also differ from Old-TOs and Old-LINERs in terms of the spatial distributions of their stellar populations and dust. Specifically, our main findings are: (1) Significant stellar population gradients are found almost exclusively in Young-TOs. (2) The intermediate age population of Young-TOs, although heavily concentrated in the nucleus, reaches distances of up to a few hundred pc from the nucleus. Nevertheless, the Half Width at Half Maximum of its brightness profile is more typically 100 pc or less. (3) Objects with predominantly old stellar populations present spatially homogeneous spectra, be they LINERs or TOs. (4) Young-TOs have much more dust in their central regions than other LLAGN. (5) The B-band luminosities of the central $\\la 1$ Gyr population in Young-TOs are within an order of magnitude of $M_B = -15$, implying masses of order $\\sim 10^7$--$10^8$ M$_\\odot$. This population was 10--100 times more luminous in its formation epoch, at which time young massive stars would have completely outshone any active nucleus, unless the AGN too was brighter in the past. ",
        "introduction": "\\label{sec:Introduction}   Low luminosity active galactic nuclei (LLAGN) are the most common form of activity in the nearby universe. Their proximity allows us to sample their properties on linear scales which are not accessible for farther AGN populations like Seyferts and quasars. At optical--UV wavelengths, however, this advantage is compensated by the difficulty in isolating the light from these intrinsically weak nuclei out of a dominant stellar ``background''.  In a series of papers, we have been working our own contribution to this field.  In Paper I (Cid Fernandes \\etal 2004a) we analyzed ground based nuclear optical spectra of a sample of 51 LINERS and TOs, while Paper II (Gonz\\'alez Delgado \\etal 2004) complements this data set with archive HST/STIS spectra of 28 nearby LLAGN. These samples cover nearly half of the LLAGN in the survey of Ho, Filippenko \\& Sargent 1997 (hereafter HFS97). Our focus throughout this series is on the stellar populations of LLAGN, with the ultimate goal of establishing the role of stellar processes in the physics of these objects.  The main result of Papers I and II is that we have uncovered a very strong relation between the nuclear stellar population and the gas excitation, as measured by [OI]/H$\\alpha$, the most important diagnostic line ratio in LLAGN. The relation is in the sense that virtually all systems containing strong populations of $\\sim 1$ Gyr or less have [OI]/H$\\alpha \\le 0.25$, while nuclei dominated by older stars span the full range in [OI]/H$\\alpha$ (up to nearly 1).  In other words, virtually all systems with relatively young stellar populations have TO-like emission line spectra, whereas older nuclei can have either TO or LINER-like line ratios. This finding lead us to introduce a combined stellar population and emission line classification into 4 types: Young-TOs, Old-TOs, Old-LINERs and Young-LINERs. This latter class is extremely rare.  This relation between line-ratios and stellar population is analogous to the one found in Seyfert 2s, where nuclei with strong circumnuclear starbursts tend to have relatively small values of [OIII]/H$\\beta$ and HeII/H$\\beta$, while systems dominated by old stars can reach larger values of these line ratios (Cid Fernandes \\etal 2001 and references therein). It is thus tempting to interpret Young-TOs as low-luminosity analogs of starburst + Seyfert 2 composites, where the relatively low excitation is explained by the starburst contribution to H$\\beta$, which dilutes [OIII]/H$\\beta$ and HeII/H$\\beta$. Intriguingly, however, Young-TOs are substantially older than the starbursts around Seyfert 2 nuclei, many of which are just a few Myr old, as deduced by the detection of O and WR stars. While Papers I and II revealed a surprisingly large number of systems containing $10^8$--$10^9$ yr populations, massive young stars of the type often found in Seyfert 2s seem to be rare in LLAGN. The analogy between Young-TOs and starburst + Seyfert 2 composites thus rests upon the hypothesis of the existence of a population of massive stars which remains essentially undetected at optical wavelengths. Clearly, further work is necessary to clarify the precise nature of the connection between stellar and gaseous properties in LLAGN.  One type of study which has been carried out for Seyferts is the mapping of stellar populations based on spatially resolved spectroscopy. Variations of absorption line equivalent widths ($W_\\lambda$) and colours ($C_\\lambda$) as a function of distance from the nucleus were mapped by means of long slit spectroscopy by Cid Fernandes, Storchi-Bergmann \\& Schmitt (1998), Boisson \\etal (2000), Gon\\'zalez Delgado, Heckman \\& Leitherer (2001), Joguet \\etal (2001). These variations can be transformed into stellar population profiles, as in the study by Raimann \\etal (2003), who found that star-formation in starburst + Seyfert 2s composites, although concentrated in the central regions, is not confined to the nucleus, but spread over the inner $\\sim 1$ kpc. Spatial gradients in spectral indices are also useful to detect the presence of a central continuum source, which dilutes the nuclear $W_\\lambda$'s with respect to off-nuclear positions. Both a compact nuclear starburst and an AGN featureless continuum can produce this effect, but in Seyfert 2s the papers above showed that significant dilution only occurs when a starburst is present in the innermost extraction.  Spatially resolved spectroscopy of LLAGN has so far been limited to relatively few studies (eg, Cid Fernandes \\etal 1998). While these previous works advanced our comprehension of individual sources, the small number of objects, differences in spectral coverage, data quality and method of analysis prevents us from drawing general conclusions about the radial distribution of stellar populations in LINERs and TOs.  In this third paper we take advantage of our recently completed spectroscopic survey to extend this type of study to a large sample of LLAGN. Variations of spectral properties with distance from the nucleus are mapped with the general goal of investigating the relation between spatial gradients, emission line and nuclear stellar population properties. In particular, we aim at evaluating the spatial distribution of intermediate age populations, a distinguishing feature of Young-TOs.  In \\S\\ref{sec:Data} we describe the data set and present examples of our spatially resolved spectra. In \\S\\ref{sec:GradentsInStellarProperties} we investigate the spatial variations of a set of spectral properties and quantify these gradients by means of suitable empirical indices.  These gradients are further analyzed in \\S\\ref{sec:Analysis} with the goal of producing estimates of the sizes, luminosities, masses and extinction of the intermediate stellar population in the central regions of Young-TOs. These estimates provide useful hints on the past and future history of these sources.  Finally, \\S\\ref{sec:Conclusions} summarizes our results. ",
        "conclusions": "\\label{sec:Conclusions}  In this third paper in our series dedicated to the stellar populations of LLAGN, we have investigated the radial variations of stellar populations properties in a sample of 42 LINERs and TOs plus 5 non-active galaxies.  The analysis was based on high quality 3500--5500 \\AA\\ long-slit spectra covering angular regions of at least $\\sim 10^{\\prime\\prime}$ in diameter with a resolution of $\\sim 1^{\\prime\\prime}$ (corresponding to $\\sim 100$ pc).  The main result of Papers I and II was the identification of a class of objects which stand apart from other LLAGN in having a strong $10^{8-9}$ yr population. In terms of emission lines nearly all of these nuclei have weak [OI]/H$\\alpha$, hence their denomination as ``Young-TOs''.  Here we have shown that Young-TOs are also distinct from other LLAGN in terms of the way stellar populations and dust are spatially distributed.  This general conclusion was reached through two distinct and complementary ways.  First, radial profiles of absorption line equivalent widths, continuum colours and the total flux along the slit were used to trace the spatial distribution of stellar populations. The results of this empirical analysis can be summarized as follows.    \\begin{enumerate}  \\item We find that the $W_\\lambda$ profiles are of essentially two types: flat and diluted.  Flat profiles, which indicate spatially uniform stellar populations, are more common, accounting for $\\sim 60\\%$ of the sample. They occur exclusively in galaxies dominated by an old, bulge-like stellar population, regardless of the LINER/TO emission line classification.  \\item Diluted profiles, on the other hand, are produced by a central ``young'' population (CYP) dominated by stars of $10^8$--$10^9$ yr, whose relatively blue continuum dilutes the $W_\\lambda$'s of metal lines with respect to their off-nuclear values.  \\item Although concentrated in the nucleus, these CYPs are spatially extended, reaching distances of up to 400 pc from the nucleus.   \\item The relation between diluted profiles and nuclear stellar population is clearly expressed by the $\\sim$ one-to-one relation between the radial dilution index $\\delta_\\lambda$ and the nuclear $W_\\lambda$ for the CaII K line: Virtually {\\it all} sources with $\\delta_K > 10\\%$ have $W_K^{\\rm nuc} < 15$ \\AA\\ and vice-versa.  This range of $W_K^{\\rm nuc}$ corresponds exactly to our definition of ``Young'' stellar population, meaning populations of 1 Gyr or less.  \\item Since these stars are found almost exclusively in objects with [OI]/H$\\alpha \\le 0.25$ (Papers I and II), it follows that stellar population gradients are typical of Young-TOs.  The fact that these stars are located in their central regions, and not spread over the whole galaxy, reinforces the suggestion that they are somehow connected to the ionization of the nuclear gas.   \\ls Second, a more detailed analysis of stellar population gradients was achieved by means of a decomposition of each spectra in terms of templates representative of very young ($\\le 10^7$ yr), intermediate age ($10^{8-9}$ yr) and old ($10^{10}$ yr) stellar populations. This analysis shows that: \\ls   \\item The CYPs in Young-TOs account for at least $\\sim 1/3$ of the total flux at 4020 \\AA. We confirm the finding of Papers I and II that these populations are dominated by $10^8$--$10^9$ yr stars. Young starbursts, even when present, make a small contribution to the optical light.  \\item Yet another property which distinguishes Young-TOs from other members of the LLAGN family is dust content.  Young-TOs are $\\sim 3$ times more extincted than Old-LINERs and Old-TOs. This finding is confirmed using the HFS97 measurements of the H$\\alpha$/H$\\beta$ ratio.  \\item Dust tends to be concentrated towards the nucleus, although asymmetric extinction profiles are also common.  \\item The radial flux distribution of CYPs have HWHM radii of $\\sim 100$ pc or less. While their core is at best partly resolved in our data, their outer regions are clearly resolved.  \\item The 4020 \\AA\\ luminosities of the CYPs are within an order of magnitude of $10^{4.3}$ L$_\\odot\\,$\\AA$^{-1}$, implying B-band absolute magnitudes of $\\sim -15$ and masses of order $\\sim 10^7$--$10^8$ M$_\\odot$. This population was 10--100 times more luminous in their formation epoch, at which time young massive stars would have completely outshone the bulge light.  The active nucleus would also be swamped by these young starbursts, unless it too was brighter in the past.  \\end{enumerate}   This investigation has unveiled several interesting connections between stellar population, emission line properties, spatial distribution and extinction, paving the road to a better understanding of the physics of low luminosity AGN. Future papers in this series will explore these and other connections in further detail."
    },
    "0410/astro-ph0410633_arXiv.txt": {
        "abstract": "We investigate the properties of the (complex) coefficients obtained in a spherical harmonic representation of temperature maps of the cosmic microwave background (CMB). We study the effect of the coefficient phase only, as well as the combined effects of phase and amplitude. The method used to check for anomalies is to construct a ``random walk'' trajectory in the complex plane where the step length and direction are given by the amplitude and phase (respectively) of the harmonic coefficient. If the fluctuations comprise a homogeneous and isotropic Gaussian random field on the sky, the path so obtained should be a classical ``Rayleigh flight'' with very well known statistical properties. We illustrate the use of this random-walk representation by using the net walk length as a test statistic, and apply the method to the coefficients obtained from a Wilkinson Microwave Anisotropy Probe (WMAP) preliminary sky temperature map. ",
        "introduction": "The process of gravitational instability initiated by small primordial density perturbations is a vital ingredient of cosmological models that attempt to explain how galaxies and large-scale structure formed in the Universe. In the standard cosmological models, a period of accelerated expansion (``inflation'') generated density fluctuations with simple statistical properties through quantum processes (Starobinsky 1979, 1980, 1982; Guth 1981; Guth \\& Pi 1982; Albrecht \\& Steinhardt 1982; Linde 1982). In this scenario the primordial density field is assumed to form a statistically homogenous and isotropic Gaussian Random Field (GRF). Over years of observational scrutiny this paradigm has strengthened its hold in the minds of cosmologists and has survived many tests, culminating in those furnished by the Wilkinson Microwave Anisotropy Probe (WMAP; Bennett et al. 2003; Hinshaw et al 2003).  Since the release of the first (preliminary) WMAP data set it has been subjected to a number of detailed independent analyses that have revealed some surprising features. Eriksen et al. (2004) have pointed out the existence of a North-South asymmetry suggesting that the cosmic microwave background (CMB) revealed by the WMAP data is not statistically homogeneous over the celestial sphere. This is consistent with the results of Coles et al. (2004) who found evidence for phase correlations in the WMAP data; see also Hajian \\& Souradeep (2003) and Hajian, Souradeep \\& Cornish (2004). The low--order multipoles of the CMB also display some peculiarities (de Oliveira-Costa et al. 2004a; Efstathiou 2004). Vielva et al. (2004) found significant non--Gaussian behaviour in a wavelet analysis of the same data, as did Chiang et al. (2004), Larson \\& Wandelt (2004) and Park (2004). Other analyses of the statistical properties of the WMAP have yielded results consistent with the standard form of fluctuation statistics (Komatsu et al. 2003; Colley \\& Gott 2003). These unusual properties may well be generated by residual foreground contamination (Banday et al. 2003; Naselsky et al. 2003; de Oliveira-Costa et al. 2004; Dineen \\& Coles 2004) or other systematic effects, but may also provide the first hints of physics beyond the standard cosmological model.  In order to tap the rich source of information provided by future CMB maps it is important to devise as many independent statistical methods as possible to detect, isolate and diagnose the various possible causes of departures from standard statistics. One particularly fruitful approach is to look at the behaviour of the complex coefficients that arise in a spherical harmonic analysis of CMB maps. Chiang et al. (2004), Chiang, Naselsky \\& Coles (2004), and Coles et al. (2004) have focussed on the phases of these coefficients on the grounds that a property of a statistically homogenous and isotropic GRF is that these phases are random. Phases can also be use to test for the presence of primordial magnetic fields (Chen et al. 2004; Naselsky et al. 2004) or evidence of non-trivial topology (Dineen, Rocha \\& Coles 2004). In this paper we suggest an extension of this idea which involves viewing the entire set of harmonics (both amplitude and phase) in the complex plane.  We describe the basic harmonic description in the next Section. In Section 3 we show how to represent the spherical harmonics as random-walks in the complex plane and give some simple analytic results. We apply the idea to the preliminary WMAP 1-year data in Section 4 and briefly discuss the results in Section 5. ",
        "conclusions": "We have presented a novel representation of the behaviour of spherical harmonic coefficients obtained in the representation of of CMB temperature maps. Indeed, the main result of this paper is the construction of two-dimensional walks as discussed in Section 3. In order to illustrate the method in the simplest way we performed a simple statistical test on the net displacement of the random walks for different $l$ modes. The results of this analysis do not provide conclusive evidence for departures from a homogeneous and isotropic GRF. For example, two modes are discrepant at 99.9 per cent confidence, compared with the 0.6 expected by chance. In a series of independent random skies with an expectation of 0.6 modes, Poisson statistics give a probability of 88 per cent that less than two modes would be discrepant. The detection of two modes is therefore not an indication of a strong departure from chance.  It is worth noting that the method does not pick out any significant departures at low $l$, which is where the non-parametric tests of Coles et al. (2004) were applied with positive results. This test works better with large numbers of independent $m$ modes. In the simple form presented here it seems less powerful than the phase-mapping technique of Chiang et al. (2003) and Chiang et al. (2004). Whatever forms of non-Gaussianity and/or non-stationarity there are in these data, this very simple test is not sensitive to them.  As we explained above, however, the net displacement of the random walk is a simple but rather crude indication of the properties of the $\\{{\\bf a_{l,m}}\\}$, as it does not take into account the ordering of the individual steps.  The possible non-Gaussian behaviour of the set $\\{{\\bf a_{l,m}}\\}$ is encoded not so much in the net displacement but in the {\\em shape} of the random walk. To put this another way, there are many possible paths with the same net displacement, and these will have different shapes depending on the correlations between step size and direction. Long runs of directed steps or regular features in the observed structure could be manifestations of phase correlation (Coles et al. 2004).  For example, the presence of a primordial magnetic field contribution to the CMB fluctuations would result in a correlation between ${\\bf a_{l-1,m}}$ and ${\\bf a_{l+1,m}}$ at fixed $m$ (Chen et al. 2004). Examples of how similar effects might be produced by foregrounds, galactic cuts and so on are discussed by Naselsky et al. (2005). These phase correlations would manifest themselves in higher-order statistics of the random walk as a function of $l$. It will also be useful, say, to combine random walks as functions of $m$ for different $l$: correlations between the net displacements of such walks would indicate the presence of phase correlations too; see Coles et al. (2004). The graphical representation of the set $\\{{\\bf a_{l,m}}\\}$ in the form illustrated by Figure 1 provides an elegant way of visualizing the behaviour of the harmonic modes and identifying any oddities. These could be quantified using a variety of statistical shape measures: moment of inertia (Rudnick, Beldjenna \\& Gaspari 1987), fractal dimension, first-passage statistics, shape statistics (e.g. Kuhn \\& Uson 1982), or any of the methods use to quantify the shape of minimal spanning trees (Barrow, Bhavsar \\& Sonoda 1985). It would be interesting to try these more sophisticated measures on subsequent releases of the WMAP data."
    },
    "0410/astro-ph0410319_arXiv.txt": {
        "abstract": "Using 7 years of MACHO survey data, we present a new determination of the optical depth to microlensing towards the Galactic bulge.  We select the sample of 62 microlensing events (60 unique) on clump giant sources and perform a detailed efficiency analysis. We use only the clump giant sources because these are bright bulge stars and are not as strongly affected by blending as other events.  Using a subsample of 42 clump events concentrated in an area of 4.5 deg$^2$ with 739000 clump giant stars, we find $\\tau =  2.17^{+0.47}_{-0.38} \\ten{-6}$ at $(l,b) = (1 \\hbox{$.\\!\\!^\\circ$} 50, -2 \\hbox{$.\\!\\!^\\circ$} 68)$, somewhat smaller than found in most previous MACHO studies, but in excellent agreement with recent theoretical predictions.  We also present the optical depth in each of the 19 fields in which we detected events, and find limits on optical depth for fields with no events. The errors in optical depth in individual fields are dominated by Poisson noise. We measure optical depth gradients of $(1.06 \\pm 0.71) \\times 10^{-6} {\\rm deg}^{-1}$ and $(0.29 \\pm 0.43) \\times 10^{-6} {\\rm deg}^{-1}$ in the galactic latitude $b$ and longitude $l$ directions, respectively. Finally, we discuss the possibility of anomalous duration distribution of events in the field 104 centered on $(l,b) = (3 \\hbox{$.\\!\\!^\\circ$} 11, -3 \\hbox{$.\\!\\!^\\circ$} 01)$ as well as investigate spatial clustering of events in all fields. ",
        "introduction": "\\lab{sec:intro} The structure and composition of our Galaxy is one of the outstanding problems in contemporary astrophysics. Microlensing is a powerful tool to learn about massive objects in the Galaxy. The amount of matter between the source and observer is typically described in terms of the microlensing optical depth, which is defined as the probability that a source flux will be gravitationally magnified by more than a factor of 1.34. Early predictions (Griest et al.\\ 1991; \\pac\\ 1991) of the optical depth towards the Galactic center included only disk lenses and found values near $\\tau = 0.5\\ten{-6}$. The early detection rate (Udalski et al.\\ 1993, 1994a) seemed higher, and further calculations (Kiraga \\& Paczy\\'{n}ski 1994) added bulge stars to bring the prediction up to $0.85 \\ten{-6}$.  The first measurements were substantially higher than this: $\\tau \\geq 3.3 \\pm 1.2 \\ten{-6}$ at $(l,b) \\approx (1 \\hbox{$.\\!\\!^\\circ$} 0, -3 \\hbox{$.\\!\\!^\\circ$} 9)$ (Udalski et al.\\ 1994b) based upon 9 events and $\\tau = 3.9^{+1.8}_{-1.2} \\ten{-6}$ at $(l,b) = (2 \\hbox{$.\\!\\!^\\circ$} 52, -3 \\hbox{$.\\!\\!^\\circ$} 64)$ (Alcock et al.\\ 1997a) based upon 13 clump-giant events and an efficiency calculation.  Many additional theoretical studies ensued, including additional effects, especially non-axisymmetric components such as a bar (e.g., Zhao, Spergel \\& Rich 1995; Metcalf 1995; Zhao \\& Mao 1996; Bissantz et al.\\ 1997; Gyuk 1999; Nair \\& Miralda-Escud\\'e 1999; Binney, Bissantz \\& Gerhard 2000; Sevenster \\& Kalnajs 2001; Evans \\& Belokurov 2002; Han \\& Gould 2003). Values in the range $0.8 \\ten{-6}$ to $2 \\ten{-6}$ were predicted for various models, and values as large $4\\ten{-6}$ were found to be inconsistent with almost any model.  More recent measurements using efficiency calculations found values of $\\tau = 2.43^{+0.39}_{-0.38} \\ten{-6}$ at $(l,b) = (2 \\hbox{$.\\!\\!^\\circ$} 68, -3 \\hbox{$.\\!\\!^\\circ$} 35)$ from 99 events in 8 fields using difference image analysis (Alcock et al.\\ 2000a; MACHO), $\\tau = 2.59^{+0.84}_{-0.64} \\ten{-6}$ at $(l,b) \\approx (1 \\hbox{$.\\!\\!^\\circ$} 0, -3 \\hbox{$.\\!\\!^\\circ$} 9)$ from 28 events using difference image analysis (Sumi et al.\\ 2003; MOA)\\footnote{We would like to notice here that the results from the difference image analyses cannot be easily compared to the ones from clump analyses. The analyses that use clump giant sources have reasonably good control over the location of sources they are sensitive to: the bulge clump giants dominate over any other possible locations. Difference image determinations are sensitive to all sources along the line of sight, and thus the determined optical depth is characteristic for the direction only and cannot be claimed to be the optical depth toward the bulge. A typically employed remedy is to assume that the only two populations of sources that matter in star counts are: the bulge one and the foreground disk one, and that only bulge sources contribute to the optical depth. Following on this assumption, one can correct a measured optical depth by a fudge factor that is supposed to account for the number of inefficient disk sources. It is not clear whether such correction is needed and the size of it is very uncertain. The corrections applied by both Alcock et al.\\ (2000a) and Sumi et al.\\ (2003) increase the optical depth by about 25\\%. The results we quote here do not include these controversial corrections.}, $\\tau = 2.0 \\pm 0.4 \\ten{-6}$ at $(l,b) = (3 \\hbox{$.\\!\\!^\\circ$} 9, -3 \\hbox{$.\\!\\!^\\circ$} 8)$ from around 50 clump-giant events in a preliminary version of this paper (Popowski et al.\\ 2001a; MACHO), and $\\tau = 0.94 \\pm 0.29 \\ten{-6}$ at $(l,b) = (2 \\hbox{$.\\!\\!^\\circ$} 5, -4 \\hbox{$.\\!\\!^\\circ$} 0)$ from 16 clump-giant events (Afonso et al.\\ 2003; EROS).  Blending is a major problem in any analysis of the microlensing data involving point spread function photometry. The bulge fields are crowded, so that the objects observed at a certain atmospheric seeing are blends of several stars. At the same time, typically only one star is lensed. In this general case, a determination of the events' parameters and the analysis of the detection efficiency of microlensing events is very involved and vulnerable to a number of possible systematic errors. If the sources are bright, one can avoid these problems.  First, parameters of the actual microlensing events are typically better constrained. Second, it is sufficient to estimate detection efficiency based on the sampling of the light curve alone.  This eliminates the need of obtaining deep luminosity functions across the bulge fields. Red clump giants are among the brightest and most numerous stars in the bulge, so we focus on these stars here. An optical depth determination using all observed microlensing events would be desirable, but would require an accurate calculation of the blending efficiency for non-clump stars, demanding much additional input including HST quality images and luminosity functions over much of the bulge.  The structure of this paper is the following. In \\S~\\ref{sec:data} we briefly describe the MACHO experiment as the source of the data used here. Section \\ref{sec:selection} is devoted to the selection of microlensing events. In particular, we discuss how we select our sample of 62 clump giant events (60 unique) and its relation to the catalog of over 500 events constructed by Thomas et al. (2004, companion paper). In \\S~\\ref{sec:blending} we test our sample for signatures of blending. Optical depth toward the Galactic bulge is derived in \\S~\\ref{sec:optical_depth}. In \\S~\\ref{sec:clustering} we consider spatial distribution of events and discuss the significance of some apparent clusters seen in the data. We summarize our results in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Using 7 years of MACHO survey data, we presented a new determination of the optical depth to microlensing towards the Galactic bulge.  We selected the sample of 62 microlensing events (60 unique) on clump giant sources (out of more than 500 total) and performed a detailed efficiency analysis. Using a subsample of 42 clump events detected during monitoring of about 739000 clump giant sources concentrated in just 4.5 deg$^2$, we found $\\tau =  2.17^{+0.47}_{-0.38} \\ten{-6}$ at $(l,b) = (1 \\hbox{$.\\!\\!^\\circ$} 50, -2 \\hbox{$.\\!\\!^\\circ$} 68)$, consistent with theoretical expectations.  We can make a comparison with previous MACHO collaboration measurements by using our current clump sample in the fields analyzed previously. For example in the 24 fields used by Alcock et al.\\ (1997a), we find 37 clump events giving $\\tau = 1.36^{+0.35}_{-0.28} \\ten{-6}$, which is two sigma below the earlier result of $\\tau = 3.9^{+1.8}_{-1.2} \\ten{-6}$. Over the 8 fields in our difference imaging study (Alcock et al.\\ 2000a), our current, almost independent sample gives $\\tau = 2.20^{+0.67}_{-0.52} \\ten{-6}$ at $(l,b) = (2 \\fdg 68,-3 \\fdg 35)$, which is in very good agreement with that earlier result of $\\tau = 2.43 ^{+0.38}_{-0.38} \\ten{-6}$ (the agreement would be somewhat worse if we introduced the fudge factor discussed in the Introduction).  Our new value of $\\tau = 2.17^{+0.47}_{-0.38} \\ten{-6}$ is also very consistent with the recent EROS collaboration (Afonso et al.\\ 2003) measurement. They found $\\tau = (0.94 \\pm 0.29) \\ten{-6}$ at $(l,b) = (2\\hbox{$.\\!\\!^\\circ$}5, -4\\hbox{$.\\!\\!^\\circ$}0)$ from 16 clump giant events and an efficiency calculation similar to ours. The values are different but they are reported at different position. The difference becomes entirely insignificant if we take into account either the numerical value of the gradient reported in \\S 5 or read off the approximate optical depth value at $b=-4.0$ from Figure~{\\ref{fig:taugrad}}. In addition, our new optical depth is in very good agreement with recent theoretical prediction of the Galactic bulge optical depth ($\\tau = 1.63 \\ten{-6}$) by Han \\& Gould (2003).  \\begin{figure*}[t] \\subfigure{\\epsfig{figure=f14left.eps, width = 0.48\\textwidth }} \\subfigure{\\epsfig{figure=f14right.eps,  width=0.48\\textwidth }} \\caption{Left panel shows the histogram of the distribution of Einstein crossing times. The right panel shows the same distribution with each event weighted by its inverse efficiency and renormalized. \\label{fig:thathist}} \\end{figure*}  We note that about 41\\% of the optical depth is in the events longer than 100 days (10 out of 62 events), and about 63\\% of the optical depth is in events longer than 50 days (21 out of 62 events). This may be at odds with some models of the Galactic structure and kinematics (e.g., Han \\& Gould 1996), but consistent with others (Evans \\& Belokurov 2002) which include streaming motion along the bar. The distributions of event durations are shown in Figure~\\ref{fig:thathist} where we plot histograms for uncorrected and efficiency-corrected cases. We note that the mean Einstein diameter crossing time of all our clump giant events is $\\VEV{\\that} = 56 \\pm 64$ days and for our fields close to the Galactic center $\\VEV{\\that}_{CGR} = 39 \\pm 31$ days. To facilitate proper comparison with theoretical models we have to weigh each event by its inverse efficiency (Table~\\ref{tab:taubyevent}). Then the averages become: $\\VEV{\\that}({\\rm eff}) = 40 \\pm 50$ days for all fields and $\\VEV{\\that}_{CGR}({\\rm eff}) = 30 \\pm 29$ days for fields in the CGR\\footnote{Because the distribution is non-Gaussian we also give the median and quartiles, which are (23.7, 12.0, 47.8) for all clump events and (18.7, 9.5, 46.9) for clump events in the CGR.}. Evans \\& Belokurov (2002) find average Einstein crossing times range from around 100 days in models that include bar streaming motion to 30 days in models without bar streaming. Our values are too uncertain to distinguish between these kinematic possibilities and are consistent with both types of models.  {\\bf Note added.} After this paper was completed, Sumi et al.\\ (2005) presented determination of the microlensing optical depth toward the Galactic bulge based on OGLE-II data. Using a sample of 32 events from the extended red clump giant region of the color-magnitude diagram, they obtained $\\tau = 2.37^{+0.53}_{-0.43} \\times 10^{-6}$ at $(l,b)=(1 \\hbox{$.\\!\\!^\\circ$} 16, -2 \\hbox{$.\\!\\!^\\circ$} 75)$. Their optical depth is fully consistent with our value of $\\tau =  2.17^{+0.47}_{-0.38} \\ten{-6}$ at $(l,b) = (1 \\hbox{$.\\!\\!^\\circ$} 50, -2 \\hbox{$.\\!\\!^\\circ$} 68)$ which is based on 42 clump giant events."
    },
    "0410/astro-ph0410172.txt": {
        "abstract": "We present  a  new catalog of   isolated galaxies obtained through  an automated systematic search.  These  2980 isolated galaxies were found in $\\approx$  2099 deg$^{2}$ of  sky in  the Sloan  Digital Sky Survey Data  Release  1  (SDSS  DR1) photometry.   The  selection  algorithm, implementing a variation on the criteria developed by Karachentseva in 1973, proved to be   very efficient and   fast.  This catalog will  be useful for  studies of the   general galaxy characteristics.   Here we report on our results. ",
        "introduction": "Over  the past  few decades,  observational  and theoretical work  has shown that  truly isolated ``field'' galaxies,  if they exist  at all, are a rarity in the universe, comprising less than 5\\% of all galaxies (\\citealt{Adams80}).  Rare  though  they  may  be, they  serve as   an important  comparison sample in studies of  the effects of environment on   galaxy    morphologies   and    star  formation     rates  (e.g., \\citealt{Adams80,Haynes80,Haynes84,Koopmann98}).      Truly   isolated galaxies, which may have experienced no major interactions in billions of years, can act as a zeropoint in these studies.  Furthermore, isolated  galaxies are  interesting in their   own right. Recent studies of   isolated galaxies include those by  \\citet{Aars02, Aars03}, who looked at the photometric and spectroscopic properties of extremely isolated  elliptical galaxies in the \\citet{Karachentseva73} catalog;  \\citet{Pisano02}  and   {\\citet{Pisano03}, who performed  an \\ion{H}{1} survey for the  gaseous  remnants of the  galaxy  formation process around  nearby ($\\la$  30$h^{-1}$~Mpc) isolated galaxies  they identified  in  the  \\citet{Tully88}'s  ``Nearby  Galaxies  Catalog''; \\citet{Sauty03}, who measured  the molecular gas  mass of 99  isolated late-type galaxies  in  the   \\citet{Karachentseva73}   catalog  using observations   of  the   $^{12}$CO(1-0)  line;  \\citet{Stocke04},  who measured  the the luminosity functions  of  isolated elliptical and S0 galaxies in the \\citet{Karachentseva73} catalog; and \\citet{Varela04}, who studied  the properties of disk galaxies  in a catalog of isolated galaxies  that they extracted from  a volume-limited sample of the CfA Redshift Survey \\citep{Huchra00}.  A  simple  method  of  identifying   isolated  galaxy  candidates  was described by \\citet{Karachentseva73}.    She selected all  galaxies in the   {\\em  Catalog    of    Galaxies  \\&    Clusters  of  Galaxies\\/} (\\citealt{Zwicky68}; Zwicky  Catalog) whose nearest neighbors, of size within a factor of four of the same  diameter, lie further than twenty diameters  away.  Then she compiled  her catalog from work with prints of the Palomar    Sky  Survey.  The original  catalog   contains  1052 candidate isolated galaxies,  which was later  reduced to 893 galaxies (\\citealt{Karachentseva80}).   She   also showed  that  a  significant number of   the final catalog  members  are located in regions  of low density in the periphery of superclusters.  (Note that \\citealt{Leon03} have  recently updated  the \\citealt{Karachentseva73} catalog with improved galaxy positions.)  \\citet{TurnerGott75} suggested    an    interesting  prescription  for isolating the classical uniform  field  population of galaxies in  the Zwicky catalog.  They placed  galaxies brighter than 14th magnitude in two  classes,  the  associated (A)  that  have  at least one  neighbor brighter than $m=14$ within 45$\\arcmin$, and the single (S) which have none.    The  S   galaxies, 43\\%   of   the   sample, appear uniformly distributed across the sky, as expected for field galaxies.  Later \\citet{HuchraThuan77}  determined   that only  12 out of 1088 S galaxies ($\\approx$1\\%) of the Turner \\& Gott sample can be considered as truly isolated.  Turner \\& Gott S  galaxies appeared single because the 14th magnitude cutoff misses fainter companions  and the angular separation criterion of $45\\arcmin$ mistakes members of nearby groups of galaxies with large angular extent for singles.  These earlier efforts for  identifying homogeneous samples of isolated galaxies   were magnitude limited to  $B<15$.    The Sloan Digital Sky Survey (SDSS; \\citealt{York00}), which will eventually cover up to one quarter of  the  sky with uniform photometry  in  five filters down to $g\\approx 22$, makes  for  an  obvious  hunting ground  for   isolated galaxies to much greater depths and volumes.  Using a modified version of Karachentseva's  isolation   criteria, we  have  extracted  such an objectively defined catalog of  isolated galaxies from the photometric catalog of the SDSS Data Release 1 (DR1;  Abazajian et al.\\ 2003).  We present this catalog as follows.  In \\S~\\ref{sec:DATA} we describe the region  of  the   sky   used  for  this    preliminary    search.   In \\S~\\ref{sec:SelectionCriteria}, we   describe the catalog construction techniques,  and we present  the  catalog in \\S~\\ref{sec:catalog}.  In \\S~\\ref{sec:conclusions} we conclude and describe our future plans.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusions} A key problem in astronomy involves the role of the environment in the formation and evolution of galaxies. In  order to answer this question it   is necessary to characterize  a  reference sample  with a minimum influence from the environment  so  that its evolution is   completely determined by nature.  A fundamental issue in  galactic evolution is the  relative importance of  initial   conditions vs.\\  environment.   To  address the  role of non-cluster environments,   we  present a  new   catalog  of  isolated galaxies in the  SDSS DR1 data. %Out  of  the 24,489029 galaxies  in DR1   best data, isolated galaxies %represent only 0.0001\\%, $\\approx$ 1.4/$deg^{2}$.  At a  detection rate of  1.4  isolated galaxies per square  degree, we expect that  the   final catalog, based upon   the  a completed   SDSS covering up  to one-quarter of the sky,  will contain on the  order of $\\approx$ 14,000 galaxies.   This catalog will allow statistical study of the properties of the  Interstellar Medium as  a function of galaxy environment,    and its  relation to  star   formation, morphology and luminosities, as well as nuclear  activity frequency (e.g.   Lisenfeld et al.  2003).  This catalog will also be useful for future studies of dark matter   density profile  of  isolated  galaxies (e.g.   Prada et al. 2003).   Finally, this  catalog  will offer a  sample  of galaxies which can greatly   aid in the investigation  of  galaxy evolution and galaxy formation.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0410/astro-ph0410405_arXiv.txt": {
        "abstract": "Recent observations of the deeply embedded L\\,1551 IRS\\,5 system permit the detailed examination of the properties of both the stellar binary and the binary jet. For the individual components of the stellar binary, we determine their masses, mass accretion rates, effective temperatures and luminosities. For the atomic wind/jet flow, we determine the mass loss rate, yielding observationally determined values of the ratio of the mass loss to the mass accretion rate, $f$. For the X-ray emitting region in the northern jet, we have obtained the jet-velocity and derive the extinction and the densities on different spatial scales. Examining the observational evidence within the framework of the x-wind theory leads us to conclude that these models are indeed potentially able to account for the observational data for this {\\it deeply embedded} source. ",
        "introduction": "Since their discovery nearly three decades ago, the unexpected phenomenon of outflows in star forming regions has remained essentially unexplained. In particular, the processes responsible for the acceleration and collimation of the flows present one of the major unsolved problems of modern astrophysics \\citep{lada85,shu87,shu00,eis00,kp00}. Considerable progress has been made in the theoretical field, but observational results had generally been of too poor quality to make direct comparisons meaningful. In this paper, we examine recent various observational results for the {\\it deeply embedded} source L\\,1551 IRS\\,5, which, when put together, finally permit the detailed comparison with theoretical models. Optically visible young stellar objects (lower mass loss and accretion rates) have previously been addressed by, e.g., \\citet[and references therein]{shang02}.  Over the years, IRS\\,5 has enjoyed a variable status of `the archetypical CO outflow source' and that of `a pathological case' \\citep[e.g.,~][]{padman97}. More recently, renewed interest in this object has arisen, in part due to the recognition of its duplicity, as binarity seems to be a very frequent phenomenon among powerful outflow drivers (M.\\,Barsony, private communication). If protostellar binarity is indeed intimately related to the physics of generating bipolar outflows, then one obviously wishes to understand the physics of the binary itself. We devoted therefore quite some effort to arrive at the understanding of the IRS\\,5 system. Based on the theory of protostellar structure and evolution, we find a solution capable of explaining observed characteristics of both IRS\\,5 itself and its associated jet flows, and which, when applied to x-wind models, successfully recovers the physical properties of the jets. Thus, in this particular case, a direct dependence on source binarity is not evidenced, as the theory of x-winds has been developed for single stellar objects.  The organisation of this paper is as follows: In Sect.\\,2, we review the evidence based on recent observations in both the optical, infrared and radio spectral regimes. In Sect.\\,3, we use published models of protostellar structure and evolution to derive the physical properties of the individual binary components and apply x-wind theory to derive some parameters relevant in the present context. In Sect.\\,4, we discuss our results and, finally, in Sect.\\,5, we briefly summarise our main conclusions. ",
        "conclusions": "Our main conclusions from this work can briefly be summarised as follows:  \\begin{itemize} \\item[$\\bullet$] Interpreting recent observational achievements within the framework of the theory of protostellar structure and evolution allows the derivation of the physical properties of the individual components of the binary protostar L\\,1551 IRS\\,5. \\item[$\\bullet$] We offer an alternate scenario to the FUOR-disk hypothesis to explain optical spectroscopic data of the scattered light from IRS\\,5, which can be put to test with currently available observing facilities. This offers potentially the unique opportunity to directly observe the surfaces of accreting protostars. \\item[$\\bullet$] The derived properties of the binary, in combination with their observed mass loss rate, lead to results which are consistent with detailed predictions by the theory of x-winds. \\end{itemize}"
    },
    "0410/astro-ph0410296_arXiv.txt": {
        "abstract": "{  Presolar SiC grains formed around Asymptotic Giant Branch (AGB) stars during their carbon-rich phase and contain heavy elements in trace amounts showing the signature of the $slow$ neutron capture process ($s$ process). Thanks to recent advances in analysis techniques, SiC data now provide extremely precise information on neutron capture cross sections and AGB models. For example, high-precision data for Mo in single SiC grains indicate that a revision of the $^{95}$Mo neutron capture cross section is needed, while data for Zr indicates that the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction cannot be a dominant neutron source for the $s$ process in AGB stars. We present model predictions for the composition of Fe-peak elements in AGB stars. These elements could be analysed in the near future thus providing further stringent constraints to our understanding of AGB stars. ",
        "introduction": "Among presolar meteoritic material, silicon carbide (SiC) grains have isotopic compositions very different to those commonly measured in the solar system. The majority of these grains are believed to have formed in the extended envelopes of carbon-rich Asymptotic Giant Branch (AGB) stars. A strong indication of this origin is the composition of elements heavier than Fe and present in very small amounts ($trace$), showing the unequivocal signature of the $slow$ neutron capture process ($s$ process) that occurs in AGB stars \\citep{gallino:98,busso:01}.  The bulk of $s$-process elements in AGB stars are produced neutrons released by the $^{13}$C($\\alpha$,n)$^{16}$O reaction. The $^{13}$C neutron source is activated in a tiny region in the top layers of the He intershell, assuming that some small-scale mixing occurs between the H-rich convective envelope and the $^{12}$C-rich He intershell at the end of each third dredge up. At H-reignition, proton captures on $^{12}$C produce $^{13}$C and a $^{13}$C {\\it pocket} is formed. The $^{13}$C is subsequently consumed via ($\\alpha,n$) reactions, neutrons are released and $s$-process isotopes are overproduced by factors of up to several thousand times their initial abundances. The pocket is then engulfed in the convective zone driven by the thermal instability, and diluted by a factor of $\\sim$ 1/20 with material from H-burning ashes and other He-intershell material.  During the thermal pulse a second small neutron exposure occurs, produced by the $^{22}$Ne($\\alpha,n$)$^{25}$Mg reaction, which is marginally activated at the bottom of the convective zone when the temperature reaches above $2.5 \\times 10^{8}$ K. The marginal activation of the $^{22}$Ne neutron source produces a small neutron burst of short duration ($\\sim$ 6 yr) but with a high average peak neutron density, $\\sim 10^{11} n$ cm$^{-3}$, affecting the final composition of branching-dependent isotopes. After the quenching of the thermal pulse the following third dredge up episode is mixes the $s$-processed material to the surface, where it is spectroscopically observed and it is included in SiC grains.  High-precision laboratory measurements of the isotopic compositions of SiC grains represent the most detailed record of the composition of AGB stars, and thus a major constraint for the theoretical models of these stars. In this paper we describe the technique used to analyse the composition of heavy elements in SiC grains (Section 2), a few main conclusions derived from comparing laboratory data to stellar models (Section 3) and predictions for the compositions in AGB envelopes of some Fe-peak elements (Section 4), which are good candidates for future high-precision measurements and will set further constraints to our knowledge of AGB stars. ",
        "conclusions": ""
    },
    "0410/astro-ph0410543_arXiv.txt": {
        "abstract": "We estimate the constraints that the recent high-redshift sample of supernovae type Ia put on a phenomenological interaction between dark energy and dark matter. The interaction can be interpreted as arising from the time variation of the mass of dark matter particles. We find that the coupling correlates with the equation of state: roughly speaking, a negative coupling (in our sign convention) implies phantom energy ($w_{\\phi }<-1$) while a positive coupling implies {}``ordinary'' dark energy. The constraints from the current supernovae Ia Hubble diagram favour a negative coupling and an equation of state $w_{\\phi }<-1$. A zero or positive coupling is in fact unlikely at 99\\% c.l. (assuming constant equation of state); at the same time non-phantom values ($w_{\\phi }>-1$) are unlikely at 95\\%. We show also that the usual bounds on the energy density weaken considerably when the coupling is introduced: values as large as $\\Omega _{m0}=0.7$ become acceptable for as concerns SNIa. We find that the rate of change of the mass $\\dot{m}/m$ of the dark matter particles is constrained to be $\\delta _{0}$ in a Hubble time, with $-10<\\delta _{0}<-1$ to 95\\% c.l.. We show that a large positive coupling might in principle avoid the future singularity known as {}``big rip'' (occurring for $w_{\\phi }<-1$) but the parameter region for this to occur is almost excluded by the data. We also forecast the constraints that can be obtained from future experiments, focusing on supernovae and baryon oscillations in the power spectra of deep redshift surveys. We show that the method of baryon oscillations holds the best potential to contrain the coupling. ",
        "introduction": "One of the main open problems in Cosmology is to determine the properties of the unclustered component called dark energy that is required to explain CMB, supernovae Ia, cluster masses and other observational data.  Although most work on dark energy assumes it to be coupled to the other field only through gravity, there is by now a rich literature on possible interactions to standard fields and to dark matter. Even before the evidences in favor of dark energy, a coupling between a cosmic scalar field and ordinary matter has been studied by Wetterich \\cite{wet88,wet90} while Damour, Gundlach and Gibbons \\cite{dam} investigated the possibility of a species-dependent coupling within the context of Brans-Dicke models. Generally speaking, any low-energy limit of higher-dimensional theories predicts the existence of scalar fields coupled to matter. More recently, several authors considered an explicit coupling between dark energy and matter: we denote this class of models as \\emph{coupled dark energy}. A partial list of works in this field is in refs. \\cite{cq,chim,post2002}. Other authors considered a coupling to specific standard model fields: to the electromagnetic field \\cite{carroll}, to neutrinos \\cite{horvat}, to baryonic or leptonic current \\cite{li}: in these cases, there are strong constraints from local observations or from variation of fundamental constants.  In this paper we confine our attention to the interaction with dark matter, as in ref. \\cite{cq}. Such a coupling is of course observable only with astrophysical \\cite{gf} and cosmological \\cite{aq} experiments involving growth of perturbations or global geometric effects. Notwithstanding the large number of papers dealing with variants of coupled dark energy, the works dedicated to constraining the interaction via supernovae Type Ia (SNIa) are very limited. The main reason is that there is yet no compelling form of the interaction (just as there is no compelling form of the dark energy equation of state). To make a step towards constraining the energy exchange between the dark components, instead of considering a specific coupling motivated by (or inspired by) some fundamental theory, we adopt here a phenomenological point of view. That is, we assume a general relation between dark energy and dark matter and derive its theoretical and observational properties. Our simple relation includes several previous coupled and uncoupled dark energy models but extends the analysis to cases which have not been tested so far. The main aim of this paper is to derive bounds on the interaction strength by analysing the recent SNIa sample of Riess et al. \\cite{riess2004}, which include several $z>1$ supernovae. This works extends previous investigations in ref. \\cite{Dalal}, which studied the same relation between dark matter and dark energy, and in ref. \\cite{agp}, which compared the recent SNIa data with a special class of coupled dark energy models motivated by superstrings and characterized by a constant ratio of densities. ",
        "conclusions": "This paper studied the behavior of a model with three components, baryons, dark matter and dark energy. Contrary to most similar analyses available in literature we included a general phenomenological coupling between dark matter and dark energy specified by a simple but rather general scaling relation. By this extension of dark energy models we have been able to put bounds on the present interaction between the two components or, equivalently, to the rate of change of the dark matter particle mass. Moreover, we can appreciate the extent to which the constraints on such fundamental quantitites as $\\Omega _{m}$ and $w_{\\phi }$ depend on the assumptions concerning the dark energy interactions.  We obtained several results that we summarize here.  \\begin{enumerate} \\item We find the constraints on the coupling constant $\\delta _{0}$, which can be interpreted as the rate of change of dark matter mass per Hubble time. We obtain $-10<\\delta _{0}<-1$ at 95\\% c.l. and $\\delta _{0}<0$ at 99\\%. \\item We find $-4.2<w_{\\phi }<-1$ (95\\% c.l.). Together with the bounds on $\\delta _{0}$ we conclude that the current SN data set favours negatively coupled phantom dark energy. In general, we find that the equation of state correlates with the coupling: positive coupling implies $w_{\\phi }<-1$, negative coupling implies phantom energy. Quite remarkably, imposing zero coupling peaks the cosmological constant as preferred value of $w_{\\phi }$. \\item Marginalizing over the coupling we derive new bounds on $\\Omega _{m0}$. In particular, we find $-4<w_{\\phi }<-1$ and $0.05<\\Omega _{m0}<0.65$ (95\\%) . The dark matter density allowed region widens considerably when a nonzero coupling is introduced in the model. \\item We find as best fit a value for the beginning of acceleration $z_{acc}=0.3$ but we cannot exclude earlier acceleration ($z>1$) at more than 95\\% c.l.. \\item We find that the addition of even a small fraction of uncoupled matter, e.g. baryons, modifies profoundly the asymptotic behavior of the model. In particolar, for $\\xi <0$ we find models in which the epoch of dark energy dominance (and acceleration) is only temporary. Notice that this occurs for $w_{\\phi }$ negative and constant. \\item The ratio of baryons-to-dark matter varies in our model. Depending on $\\xi $ it may increse or decrease with time. This may offer new methods of constraining a preferential coupling to dark matter. \\item The big rip that in uncoupled models occurs for $w_{\\phi }<-1$ can be prevented if $\\delta _{0}>3|w_{\\phi }|\\Omega _{\\phi ,0}$; however, these values are unlikely at more than 99\\% c.l. \\end{enumerate} Future experiments will constrain the equation of state and the coupling to a much better precision. We find that an experiment with the specification of SNAP might reduce the errors on $\\delta _{0},w_{\\phi }$ by a factor of five roughly. The method of the baryon oscillations could reduce the error on $\\delta _{0}$ by a factor of 25 roughly. Needless to say, these forecasts depend on the exact experimental setting; however, they give a feeling of the expected precision on the dark energy parameters that can be reached a few years from now. Whether the final outcome will show any trace of coupling (or, for that matter, any trace of deviation from a pure cosmological constant) is one of the most exciting question to ask."
    },
    "0410/astro-ph0410282_arXiv.txt": {
        "abstract": "We review the statistical properties of relativistic Doppler boosting relevant for studies of relativistic jets from compact objects based on radio--X-ray(--mass) correlations, such as that found in black-hole X-ray binaries in the low/hard state, or the ``fundamental plane'' of Merloni, Heinz, \\& DiMatteo.  We show that the presence of only moderate scatter in such relations does not necessarily imply low Lorentz factors of the jets producing the radio emission in the samples under consideration.  Applying Doppler beaming statistics to a large sample of XRBs and AGN, we derive a limit on the width of the Lorentz factor distribution of black holes with relativistic jets: If the X-rays are unbeamed (e.g., if they originate in the accretion disk or in the slower, innermost part of the jet), the width of the $\\beta\\Gamma$ distribution should be about one order of magnitude or less.  If the scatter about the ``fundamental plane'' is entirely dominated by relativistic beaming, a lower limit on the mean Lorentz factor $\\langle\\beta\\Gamma\\rangle > 5$ can be derived.  On the other hand, if the X-rays are boosted by the same factor as the radio emission, we show that the observed scatter cannot be reasonably explained by Doppler boosting alone. ",
        "introduction": "\\label{sec:introduction} Astrophysical jets from black holes in active galactic Nuclei (AGNs) and X-ray binaries (XRBs) are known to propagate with velocities close to the speed of light.  Evidence for relativistic bulk motion stems from variability and compactness limits in compact cores \\citep[][]{jones:74}, from superluminal motions of knots\\citep[][]{whitney:71,cohen:71}, and from doppler boostes jet-to-counter jet flux ratios \\citep[e.g.][]{perley:82}. While proper motion measurements typically only constrain pattern speeds, the body of evidence points towards a range of mildly to ultra-relativistic jet speeds.  As a result, the observed flux densities from these jets can be affected by strong Doppler boosting, which must be corrected for if one wants to infer the intrinsic luminosity of the jet, emitted in the rest frame of the plasma, which is important when calculating the physical parameters of the jet.  Recently, a correlation between the radio emission from steady jets and the hard (2-10 keV) X-ray emission has been found in black hole XRBs \\citep{corbel:03,gallo:03}.  This relation implies that in a given XRB in the low/hard state \\citep[see][for a review of the state classification in XRBs]{mcclintock:04}, the radio flux is proportional to the X-ray flux to the 0.7th power.  This trend has been observed for different sources.  The normalisation of the relation differs only by factors of a few for different XRBs.  Since the radio emission stems from the jet, it will be affected by Doppler boosting and its brightness will differ strongly for different viewing angles.  In the standard scenario, the X-rays come from the disk (or alternatively the inner, slow moving part of the jet) and are thus not significantly beamed.  The lack of strong scatter in the radio-X-ray relation has been interpreted as an indication that Doppler boosting is weak in the steady jets of XRBs in the low/hard state and that they must therefore move with only mildly relativistic speeds \\citep{gallo:03}.  Furthermore, \\cite{merloni:03} and \\citep{falcke:04} have found a strong correlation between radio luminosity, X-ray luminosity, and black hole mass for samples of black holes spanning a wide range in black hole masses and accretion rates (called the ``fundamental plane of black hole activity'', FP for short).  Again, this relation shows a certain amount of intrinsic scatter, part of which might be contributed by Doppler boosting.  Thus, studying this scatter can provide constraints on the presence or absence of Doppler boosting in the jets that produce the radio emission.  The nature of Doppler boosting has been studied exhaustively in the literature, specifically in regard to its statistical effects on samples of objects emitting beamed radiation \\citep{orr:82,urry:84,urry:91,urry:95,morganti:95,lister:97,lister:03}. Typically, one deals with a sample of radio sources that have been selected in the radio band and consider the effects of Doppler boosting on their luminosity function and on possible selection effects, not knowing what the unboosted flux of any particular source in the sample is.  What makes the situation considered in this paper different is that we actually {\\em have} an unbiased estimator of the unbeamed radiation: in the case of XRBs it is the X-ray luminosity, in the more general case of black holes of all masses it is the FP relation that links radio luminosity, X-ray luminosity, and black hole mass.  It is therefore worth considering the statistical properties of relativistic Doppler boosting under those conditions.  In \\S\\ref{sec:boosting} we will review the basic properties of Doppler boosting and define the statistical integrals necessary for the remainder of the paper.  In \\S\\ref{sec:binaries} we will apply these results to individual pairs of XRBs and argue that the observed moderate amount of scatter in the XRB radio-X-ray relation alone cannot be used to argue for low jet velocities.  In \\S\\ref{sec:agn} we apply the same method to the FP sample to derive constraints on the Lorentz factor distribution of the source in the sample.  Section \\ref{sec:conclusion} presents our conclusions. ",
        "conclusions": "\\label{sec:conclusion} We showed that the scatter in the radio-X-ray relation in XRBs and in the ``fundamental plane'' relation in accreting black holes can be used to constrain the width of the Lorentz factors distribution of the jets in these sources.  It cannot be used to put an upper limit on the mean Lorentz factor $\\langle \\Gamma \\rangle$ of the jets in the sample.  However, if all of the scatter is indeed due to relativistic Doppler boosting, we show that a lower limit can be put on $\\langle \\beta\\Gamma \\rangle$.  Both log-normal and log-flat distributions in $\\beta\\Gamma$ fit the observed scatter well.  We show that, if the X-rays are produced in the jet, they either have to originate in an unbeamed portion of the jet (close to the base) or other sources of scatter must dominate in the ``fundamental plane'' relation. \\vspace*{6pt}  \\thanks{We would like to thank Rob Fender, Sera Markof, Mike Nowak, and the anonymous referee for helpful insights and discussions. Support for this work was provided by the National Aeronautics and Space Administration through Chandra Postdoctoral Fellowship Award Number PF3-40026 issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the National Aeronautics Space Administration under contract NAS8-39073.}"
    },
    "0410/astro-ph0410411_arXiv.txt": {
        "abstract": "Although point caustics harbour a larger potential for measuring the brightness profile of stars during the course of a microlensing event than (line-shaped) fold caustics, the effect of lens binarity significantly limits the achievable accuracy. Therefore, corresponding close-impact events make a less favourable case for limb-darkening measurements than those events that involve fold-caustic passages, from which precision measurements can easily and routinely be obtained. Examples involving later Bulge giants indicate that a $\\sim 10\\,\\%$ misestimate on the limb-darkening coefficient can result with the assumption of a single-lens model that looks acceptable, unless the precision of the photometric measurements is pushed below the $1\\,\\%$-level even for these favourable targets. In contrast, measurement uncertainties on the proper motion between lens and source are dominated by the assessment of the angular radius of the source star and remain practically unaffected by lens binarity. Rather than judging the goodness-of-fit by means of a $\\chi^2$ test only, run tests provide useful additional information that can lead to the rejection of models and the detection of lens binarity in close-impact microlensing events. ",
        "introduction": "In order to resolve the surface of the observed star during a microlensing event, the magnification pattern created by the lens needs to supply a large magnification gradient. Two such configurations meeting this requirement have been discussed extensively in the literature: a point-caustic at the angular position of a single point-like lens \\citep{WM94,NemWick94,Gould94:extsrc,BC95,BC96,Witt95,GW96,GG:detld,Heyetal2000,Hey2003} and a line-shaped fold caustic produced by a binary lens \\citep{SchneiWei:QSO,SchneiWag, GG:detld,Rhie:ld,Do:SecondLD,Do:Fold,Do:FoldLD}. It has been pointed out by \\citet{GG:detld} that fold-caustic events are more common and their observation is easier to plan, whereas close-impact events where the source transits a point caustic can provide more information. However, I will argue that this apparent gain of information can usually not be realized due to potential lens binarity. In contrast to fold caustics which form a generically stable singularity, point caustics are not stable and do not exist in reality. Instead, there is always a small diamond-shaped caustic containing four cusps.  In this paper, the influence of lens binarity on the measurement of stellar limb-darkening coefficients and proper motion is investigated and the arising limitations of the power of close-impact events where the source passes over a single closed caustic are discussed. Sect.~\\ref{sec:cie} discusses the basics of close-impact microlensing events with the effect of source size, the potential of measuring stellar proper motion and limb darkening, and the effect of lens binarity. Sect.~\\ref{sec:binvsld} shows the influence of lens binarity on the extraction of information from such events. First, the effect of lens binarity on the light curves is demonstrated by means of two illustrative examples involving K and M Bulge giants. Subsequently, a simulation of data corresponding to these configurations is used to investigate the potential misestimates of parameters if lens binarity is neglected. Sect.~\\ref{sec:conclusions} presents the final conclusions and a summary of the results. ",
        "conclusions": "\\label{sec:conclusions}  Generally speaking, the potential binarity of the lens limits the accuracy of a limb-darkening measurement of the observed source star for close-impact microlensing events. The caustic near the center of the lens star is never exactly point-like, but always a small diamond with four cusps. If a point lens is assumed, the obtained limb-darkening coefficient is systematically offset, while the nature of the lens binarity may not be apparent. Although the inclusion of a binary lens in the modelling of the event will yield a proper limb-darkening measurement, the additional degrees of freedom still diminish the achievable precision. Unfortunately, such a computation is extremely demanding. Simulations for typical event configurations involving K or M Bulge giants show that single-lens models for binary lens events that involve an offset of  $\\sim\\,10~\\%$ on the limb-darkening coefficient $\\Gamma$ look acceptable both from $\\chi^2$ and from run tests, unless photometry significantly below the $1~\\%$-level is possible. The measurement uncertainty of the proper motion $\\mu = \\theta_\\star/t_\\star$ however is dominated by the determination of the angular source radius $\\theta_\\star$, involving stellar spectra, compared to which the accuracy limits on $t_\\star$ caused by lens binarity for close-impact microlensing events are negligible.  In contrast, the measurement of limb darkening from fold-caustic passages does not suffer from any of the problems encountered for close-impact events. A binary lens is assumed a-priori, but the measurement of limb darkening only depends on a smaller set of local properties rather than on the complete binary-lens parameter space. The use of the local approximation of the light curve in the vicinity of the fold-caustic passage \\citep[e.g.][]{PLANET:SMC,Do:Fold} makes the computation quite inexpensive and easy. Moreover, fold-caustic exits can be well predicted. Once a caustic entry has been observed, it is clear that a corresponding exit will occur, and the light curve on the rise to the caustic exit peak allows to predict the time of the caustic exit with sufficient precision usually more than a day in advance. A fair coverage of the caustic entry as well as the determination of the spectral type of the source star and an early measurement of the event time-scale $t_\\rmn{E}$ even allows a rough guess on the passage duration and therefore a proper a-priori assessment of the potential for measurements of the source brightness profile. While the characteristic properties of a fold-caustic exit, which provides a full scan of the source from the leading to the trailing limb, can therefore be estimated in advance, corresponding predictions for close-impact events are only possible after the leading limb has alreaded passed the caustic, leaving only the second half of the caustic passage involving the trailing limb. From fold-caustic passages, a linear limb-darkening coefficient can be obtained routinely and easily with a precision of less than $5~\\%$ \\citep[e.g.][]{Do:FoldLD}. Hence, events involving fold-caustic passages turn out to be the clear favourite over close-impact events for measuring limb-darkening coefficients apart from those where the source transits a cusp. The latter provide an opportunity to determine further limb-darkening coefficients beyond a linear law, which is quite challenging and usually impossible for fold-passage events \\citep{Do:SecondLD}.    However, events with fold-caustic passages make the worse case for proper motion measurements. Fits to the light curve in the vicinity of the caustic passage only yield $t_\\star^\\perp = t_\\star/(\\sin \\phi)$ \\citep[e.g.][]{EROS:SMC,PLANET:SMC,Do:Fold}, where $\\phi$ is the caustic-crossing angle, which needs to be determined from a model involving the full set of binary lens parameters. While a precise measurement of  $t_\\star^\\perp$ is possible, the uncertainty in $t_\\star$ is severely limited by degeneracies and appararent ambiguities for the caustic-crossing angle $\\phi$ \\citep{Do99:CR,PLANET:sol,joint}, which can have a comparable or even larger influence on the determination of the proper motion $\\mu$ than uncertainties in the angular stellar radius $\\theta_\\star$."
    },
    "0410/astro-ph0410627_arXiv.txt": {
        "abstract": "We have carried out a sensitive high-latitude ($|b|$ $>$ 15$^{\\circ}$) HI 21cm-line absorption survey towards 102 sources using the GMRT. With a $3\\sigma$ detection limit in optical depth of $\\sim$0.01, this is the most sensitive HI absorption survey. We detected 126 absorption features most of which also have corresponding HI emission features in the Leiden Dwingeloo Survey of Galactic neutral Hydrogen. The histogram of random velocities of the absorption features is well-fit by two Gaussians centered at V$_{lsr}$ $\\sim$ 0 km s$^{-1}$ with velocity dispersions of 7.6 $\\pm$ 0.3 km s$^{-1}$ and 21 $\\pm$ 4 km s$^{-1}$ respectively. About 20\\% of the HI absorption features form the larger velocity dispersion component. The HI absorption features forming the narrow Gaussian have a mean optical depth of 0.20 $\\pm$ 0.19, a mean HI column density of (1.46 $\\pm$ 1.03) $\\times$ 10$^{20}$ cm$^{-2}$, and a mean spin temperature of 121 $\\pm$ 69 K. These HI concentrations can be identified with the standard HI clouds in the cold neutral medium of the Galaxy. The HI absorption features forming the wider Gaussian have a mean optical depth of 0.04 $\\pm$ 0.02, a mean HI column density of (4.3 $\\pm$ 3.4) $\\times$ 10$^{19}$ cm$^{-2}$, and a mean spin temperature of 125 $\\pm$ 82 K. The HI column densities of these fast clouds decrease with their increasing random velocities. These fast clouds can be identified with a population of clouds detected so far only in optical absorption and in HI emission lines with a similar velocity dispersion. This population of fast clouds is likely to be in the lower Galactic Halo. ",
        "introduction": "\\label{sec:intro} A number of HI 21cm-line absorption surveys have been carried out in the last 50 years or so. While interferometric surveys are a better alternative for such studies, since it rejects the more extended HI emission, the lack of collecting area of the interferometers has limited the sensitivity of the HI absorption surveys. Until recently, the Very Large Array (VLA) used to be the only instrument with a collecting area comparable to large single dish telescopes. From the various HI absorption surveys carried out so far, more than 600 absorption spectra are available, but the optical depth detection limits of more than 75\\% of these are above 0.1 (see for eg. Mohan et al, \\nocite{me2004}2004 - hereafter paper I, for a summary of previous surveys).  The motivation for the present survey, the observing strategy, the sources observed, their HI 21cm line spectra and the parameters of the discrete HI line components are presented in \\nocite{me2004}paper I. Here, we discuss their interpretation. As was mentioned in \\nocite{me2004}paper I, the low optical depth regime of Galactic HI is largely unexplored, except for the HI absorption study by Dickey et al \\nocite{dst78}(1978) and by Heiles \\& Troland (\\nocite{ht03a}2003a, \\nocite{ht03b}b). Dickey et al. \\nocite{dst78}(1978) measured HI absorption/emission towards 27 extragalactic radio sources located at high and intermediate Galactic latitudes ($|b|$ $>$ 5$^{\\circ}$) using the Arecibo telescope. The rms optical depth in their spectra was $\\sim$0.005. These profiles can be considered the best in terms of signal-to-noise ratio, though in many of the profiles the systematics in the band dominate the noise. Despite these limitations, Dickey et al \\nocite{dst78}(1978) noted that the velocity distribution of HI absorption features is dependent on their optical depths. For the optically thin clouds ($\\tau$ $<$ 0.1), the velocity dispersion was $\\sim$11 km s$^{-1}$, whereas for the optically thick clouds ($\\tau$ $>$ 0.1) this value was $\\sim$6 km s$^{-1}$. Similar trend was also noticed in the Effelsberg-Green Bank survey (Mebold et al, \\nocite{egs82}1982). More recently, Heiles \\& Troland \\nocite{ht03b}(2003b) also found indications for an independent population of lower optical depth HI absorption features.  One of the first efforts to study the nature of the ISM was the observation of interstellar absorption in the optical line of singly ionized calcium (CaII) towards early type stars by Adams \\nocite{adams}(1949). He noted that the observed extent in the radial velocities exceeded 50 km s$^{-1}$ in the local standard of rest frame (LSR). Blaauw \\nocite{b52}(1952) analysed the random velocity distribution of interstellar absorption lines in Adams' data. One of the main conclusions of this analysis was that the random velocity distribution of interstellar absorption features cannot be explained by a single Gaussian distribution. Support for this conclusion came from the study of Routly \\& Spitzer \\nocite{rs52}(1952), who found the ratios of column densities of neutral sodium (NaI) to singly ionized calcium (CaI) to decrease systematically with increasing random velocities of the absorption features. This effect, called the ``Routly-Spitzer effect'' was later confirmed from a much larger sample of stars by Siluk \\& Silk \\nocite{ss74}(1974). Field, Goldsmith \\& Habbing  \\nocite{fgh69}(1969) modeled the ISM as cool dense concentrations of gas, often referred to as ``interstellar clouds'' (the Cold Neutral Medium or CNM), in pressure equilibrium with a warmer intercloud medium (the Warm Neutral Medium or WNM). While this initial model of the ISM has been refined considerably by later studies, the basic picture of the ISM with cold diffuse clouds and the warmer intercloud medium has remained. However, there has been a discrepancy in the velocity distribution of interstellar clouds (the CNM) obtained from the optical absorption line studies and from the HI 21-cm line observations. The optical as well as UV absorption line studies indicated the presence of features with larger spread in random velocities, which was absent in the 21-cm line observations (Mohan et al, \\nocite{me2001}2001 \\& the references therein). This is clear from Fig. \\ref{fig:ss74rad72}.  \\begin{figure} \\begin{center} \\includegraphics[width=7.5cm,height=13.0cm,angle=-90]{figures/ss74o_rad72.ps} \\end{center} \\caption{Histogram of radial (random) velocities from the optical absorption line studies by Siluk \\& Silk (1974), along with that from the HI absorption line survey by Radhakrishnan et al., 1972 (dashed line; shaded). The higher velocity tail is missing in the radio observations.} \\label{fig:ss74rad72} \\end{figure}  According to one hypothesis, the high velocity ($|v|$ $\\gtsim$ 15 km s$^{-1}$) optical absorption lines arise in interstellar clouds, shocked and accelerated by supernova remnants in their late phases of evolution (Siluk \\& Silk, \\nocite{ss74}1974; Radhakrishnan \\& Srinivasan, \\nocite{rs80}1980). Such a mechanism would naturally result in the higher random velocity clouds being warmer and having lower column densities due to shock heating and evaporation, as compared to the low random velocity clouds. Lower column density explains the non detections of such clouds in HI emission, since the intensity of the HI 21cm line emission is directly proportional to the HI column density. Higher temperature, along with lower column density ($N_{\\scriptscriptstyle{\\rm{HI}}}$) results in lower optical depth, since $\\tau_{\\scriptscriptstyle{\\rm{HI}}}$ $\\propto$ $N_{{\\scriptscriptstyle{\\rm{HI}}}}/T$, where $T$ is the excitation temperature of the spectral line (Siluk \\& Silk, \\nocite{ss74}1974; Radhakrishnan \\& Srinivasan, \\nocite{rs80}1980; Rajagopal et al, \\nocite{jd98b}1998). If this scenario is correct, then a sensitive HI absorption survey should detect those features with lower optical depth and higher random velocities which are presumably the counterparts of the higher velocity optical absorption lines. We have used the present dataset to address this scenario.  In the next section we discuss a method to estimate the contributions from the Galactic differential rotation to the observed radial velocities of the HI absorption features. To supplement the absorption spectra, the corresponding HI 21cm emission profiles from the Leiden-Dwingeloo HI emission survey (LDS, Hartman \\& Burton, \\nocite{lds}1995) were analyzed. Appendices B \\& C of paper I list the HI line profiles (absorption \\& emission), and their fitted line parameters respectively. In section \\ref{sec:old_radio+opt}, we compare the present dataset with the previous HI absorption surveys. In Section \\ref{sec:abs}, we discuss the statistics of HI absorption line parameters obtained from the GMRT. The velocity dispersion of the interstellar clouds estimated from the present survey and the low optical depth features are highlighted. In section \\ref{sec:discuss}, we compare our results with various other existing studies to discuss the location of the newly detected high random velocity and low optical depth features. ",
        "conclusions": "\\label{sec:discuss} The present high-latitude, high-sensitivity HI absorption survey from the GMRT has confirmed many of the known results. These results are summarised in the histograms (Fig. \\ref{fig:lineparams}) of random velocity, peak optical depth, and spin temperature of the absorption components. In addition, variation of peak optical depth w.r.t. spin temperature, random velocity and Galactic latitude are displayed in Figs. \\ref{fig:tauts}, \\ref{fig:velotau} and \\ref{fig:tau_b} respectively. However, a completely new result is the detection of the low optical depth absorption features forming the high velocity tail in the histogram shown in Fig. \\ref{fig:lineparams}a. Most of these low optical depth features are below the detection limit of the earlier HI absorption surveys. The velocity histogram (Fig. \\ref{fig:lineparams}a) is well-fit by two Gaussians indicating two populations of HI absorbing clouds identified by velocity dispersions $\\sim$ 7 and $\\sim$ 21 km s$^{-1}$ respectively (Fig. \\ref{fig:tau_histo_gfit}). While the slow clouds were first detected in the optical absorption lines and subsequently in the HI absorption and emission surveys, the fast clouds were only detected in optical absorption lines. The non-detection of the fast clouds remained a puzzle for a long time. The present observations have clarified the nature of the fast clouds to some extent.  There have been attempts in the past to explain the high velocity tail seen in the histogram of random velocities of the optical absorption lines ( Siluk \\& Silk \\nocite{ss74}(1974), Radhakrishnan \\& Srinivasan \\nocite{rs80}(1980), Rajagopal et al \\nocite{jd98b}(1998)). According to these authors, the high velocity optical absorption lines arise in interstellar clouds, shocked and accelerated by supernova remnants in their late phases of evolution. The fast clouds are therefore warmer and also of lower HI column density as compared to the slow clouds, due to shock heating and evaporation.  The present observations indicate that these fast clouds have three times larger velocity dispersion and ten times lower column densities compared to the slow clouds as might be expected if they were from a shocked population of clouds. The decrease in the HI column densities of the fast clouds as a function of their random velocities (Fig. \\ref{fig:nh_v_plt}) is also consistent with this scenario. The shocked HI clouds are also expected to be warmer than the slow clouds. The mean spin temperature of the fast clouds detected in the GMRT survey is similar to that of the standard slow clouds (Fig. 6). However, this might be a selection effect since for a given optical depth detection limit and an HI column density, clouds with lower spin temperature will be preferentially detected.  The fast clouds with three times higher dispersion are expected to have a scale height about ten times larger compared to the slow HI clouds. Given an effective thickness of 250 pc for the slow clouds, the fast clouds can have an effective thickness of $\\sim$ 2.5 kpc. Therefore, fast clouds can be part of the halo of the Galaxy. Alternative evidence supports the existence of atomic gas in the Galactic halo. Albert \\nocite{al83}(1983) selected lines of sight wherein a halo star and a nearby star are aligned one behind the other. Absorption lines of TiII, CaII and NaI were measured towards these stars. Though with a limited sample size of nine directions, the results of this study clearly indicate that the higher velocity absorption lines are seen only towards the distant star. She concludes that the high velocity tail seen in the optical line studies arise from the gas in the Galactic halo. Later, similar studies (Danly, \\nocite{danly89}1989; Danly et al, \\nocite{danly92}1992; Albert et al, \\nocite{al94}1994; Kennedy et al, \\nocite{gc01}1996, \\nocite{gc02}1998a, \\nocite{gc03}1998b) confirmed this result. The halo gas shows a larger spread in velocity, as compared to the gas in the Galactic disk. Furthermore, recent HI emission studies using the Green Bank telescope have led to the discovery of a population of discrete HI clouds in the Galactic halo with a velocity dispersion similar to that of the fast clouds reported here (Lockman \\nocite{jay02}2002). The mean HI column density of these clouds was estimated to be a few times 10$^{19}$ cm$^{-2}$. He concludes that a cloud population with a line of sight velocity dispersions of $\\sigma_v$ $\\sim$15 -- 20 km s$^{-1}$ is capable of explaining the observed velocity spread of these features. Lockman \\nocite{jay02}(2002) finds that many clouds in the halo have narrow line widths implying temperatures below 1000K. In addition, indications for a core-halo structure for these clouds was also found. The fast clouds detected in the present survey are likely to be part of the same halo gas detected in HI emission from Green Bank. The fast clouds, once detected only in optical absorption lines, have now been detected in both HI absorption and emission leading to a clearer picture of the interstellar medium.\\\\  \\newpage  \\noindent{{\\textbf{Acknowledgements:}}} \\\\ \\noindent{We wish to thank Shiv Sethi for useful discussions related to the numerical techniques. We thank the referee, Miller Goss, for detailed comments and constructive criticisms resulting in an improved version of this paper. We thank the staff of the GMRT who made these observations possible. The GMRT is operated by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research. This research has made use of NASA's Astrophysics Data System.}"
    },
    "0410/nucl-th0410063_arXiv.txt": {
        "abstract": "We study relativistic mean-field models with hadron masses and coupling constants depending self-consistently on a scalar meson field. We demonstrate that by the field redefinition some models can be equivalently transformed into each other. Thereby the large variety of scaling functions for masses and couplings can be reduced to a restricted set of functions with a constrained dependence on a scalar field. We show how by choosing properly the latter scaling functions one may stiffen or soften the equation of state at high densities and simultaneously increase the threshold density for the direct Urca process without any change of the description of nuclear matter close to the saturation density. The stiffening of the equation of state might be motivated by recent neutron star mass measurements, whereas the increase of the threshold density for the direct Urca process is required by the analysis of neutron star cooling data. We demonstrate that if a rho meson is included in a mean-field model as a non-Abelian gauge boson, then there is a possibility for a charged rho-meson condensation in dense nuclear matter. We show that such a novel phase can be realized in neutron star interiors already for sufficiently low densities, typically $\\sim 3\\div 4~n_0$, where $n_0$ is the nuclear saturation density. In the framework of the relativistic mean field model the new phase arises in a second-order phase transition. The appearance of a $\\rho^-$ condensate significantly alters the proton fraction in a neutron star but changes moderately the equation of state. The neutrino emissivity of the processes involving a $\\rho^-$ meson condensate is estimated. ",
        "introduction": "The baryon matter at densities relevant for neutron star (NS) interiors is convenient to describe in terms of relativistic mean-field (RMF) models, cf. ~\\cite{G00} and references therein. In these models the nucleon mass decreases with an increase of the nucleon density due to coupling to a scalar field, whereas the effective masses of scalar ($\\sigma$), vector ($\\omega$) and isovector ($\\rho$)  mesons, and their coupling constants, are usually assumed to be density and field independent quantities. The RMF models provide a simple and economical tool for construction of the equation of state (EoS) for dense baryonic matter. By tuning parameters of a RMF model at a saturation density one may get an EoS similar to the one that follows from more involved microscopic calculations, e.g., as that of Urbana--Argonne group (A18+$\\delta v$+UIX*)~\\cite{APR98}.  The latter calculation provides probably the most realistic EoS for nuclear densities up to $4\\,n_0$, but it uses a non-relativistic potential and violates causality at higher densities. As an extrapolation to higher densities RMF models are practical. However all RMF models have the very same an unpleasant feature. They produce  a  large fraction of protons in the NS matter, cf. ~\\cite{BB77,Boguta81}. This permits a very efficient cooling process, direct Urca (DU) process $n\\rightarrow pe\\bar{\\nu}$, for $n>n_{\\rm crit}^{\\rm DU}$, with a low threshold density $n_{\\rm crit}^{\\rm DU}\\lsim 3~n_0$. Ref.~\\cite{MSTV} emphasized this fact as an obvious shortcoming of RMF models. Oppositely, ref.~\\cite{LP} and some subsequent works~\\cite{YAKDU} tried to use this fact to develop a so-called DU-based scenario of a NS cooling (based on $n_{\\rm crit}^{\\rm DU}\\sim (2\\div 3)n_0$).  The NS cooling \"data\" (surface temperatures of NSs as a function of the NS age) can be subdivided in three groups, corresponding to \"slow cooling\", \"intermediate cooling\" and \"rapid cooling\" of neutron stars. If the density in the NS central region exceeds $n_{\\rm crit}^{\\rm DU}$, which corresponds to the mass $M_{\\rm crit}^{\\rm DU}$, the star cools down very fast, radiating mostly from its central region. Thus, the NSs with $M<M_{\\rm crit}^{\\rm DU}$ cool slowly, whereas the stars with the mass $M$ only slightly above $M_{\\rm crit}^{\\rm DU}$ would be already characterized by a very rapid cooling. Such a scenario was criticized in~\\cite{TTTTT02,T04,BGV04}. It is doubtful that many NSs, belonging to an intermediate cooling group, have the very similar masses. It has been already experimentally established that NSs may have essentially different masses. E.g., masses of two NSs in double NS systems are measured with the very high precision, $M_{\\rm B1913+16}=1.4408\\pm 0.0003 M_{\\odot}$, and $M_{\\rm J0737-3039B}=1.250\\pm 0.005 M_{\\odot}$~\\cite{S04}.  Authors of many RMF models do not care about the DU threshold at all and obtain so low values of $n_{\\rm crit}^{\\rm DU}$ that $M_{\\rm crit}^{\\rm DU}<1\\,M_{\\odot}$. However according to the standard mechanisms only the NSs with $M\\gsim 1\\,M_{\\odot}$ can be formed in supernova explosions. The mentioned above microscopic Urbana-Argonne model produces $n_{\\rm crit}^{\\rm DU}\\simeq 5~\\,n_0$, that corresponds to a heavy NS with $M_{\\rm crit}^{\\rm DU}\\simeq 2\\,M_{\\odot}$. If one used a standard RMF model  fitted in \\cite{KV} to reproduce the Urbana-Argonne EoS for $n<4n_0$, but,  as well as other RMF models, producing a low DU threshold ($n_{\\rm crit}^{\\rm DU}\\simeq (2.6\\div 2.7)\\,n_0$), one would conclude that the majority of the NSs seen in soft $X$ rays are low mass objects with  the mass $M< 1.3\\,M_{\\odot}$. The latter value is however below the averaged value of the NS mass ($M\\simeq 1.35\\pm 0.04\\,M_{\\odot}$) measured in NS binaries \\cite{TC99}.   Another problem is that there might exist heavy NSs. Recent measurements of NS masses in binary compact systems yielded $M_{\\rm J0751+1807}=2.2\\pm 0.2 M_{\\odot}$ \\cite{S04} and $M_{\\rm 4U1700-37}=2.44\\pm 0.27~ M_{\\odot}$ \\cite{Q03}. The latter object might be a black hole due to a lack of pulsation but the former object, J0751+1807, is definitely a NS. The confidence level of 95\\% still allows to have a smaller limiting NS mass than $2~M_{\\odot}$ (however the  limiting NS mass is surely higher than $1.6~M_{\\odot}$). Additional information on heavy NSs may be obtained from quasi-periodic oscillation (QPO) sources. These are NSs that are $X$-ray sources at frequencies of the orbiting accreting matter. If QPOs  originate from the innermost stable orbit~\\cite{Miller} of the accreting matter, their observed values imply that the accreting NS has the mass $\\simeq 2.2~M_{\\odot}$ in the case of 4U1820-30. The value $2.2~M_{\\odot}$ is the limiting mass of a NS following the Urbana--Argonne EoS. Unfortunately this EoS is acausal at high densities. Ref.~\\cite{HHJ} suggested an analytic fit of the Urbana--Argonne EoS for $n\\lsim 4 n_0$, which respects causality for larger densities. We will denote this parameterization as the HHJ EoS. As a price, the limiting mass decreased down to $\\simeq 2M_{\\odot}$. Moreover, all phase transitions, like pion and kaon condensations, and quark matter,  which are assumed to be possible in NS interiors, may essentially  soften EoS decreasing further the limiting mass of a NS\\footnote{The Urbana--Argonne calculation \\cite{APR98} already includes neutral pion condensate for $n>1.3~n_0$ that only slightly softens EoS in their case.}. Thus modern microscopically motivated EoS and possible phase transitions in NS interiors depend crucially on one precision measurement of a heavy NS  with $M>(2\\div 2.2)~M_{\\odot}$.  Within  a standard RMF model it would be hard to describe heavy NSs with $M>2.2~M_{\\odot}$ without either an essential increase of the compressibility parameter $K$ or a decrease of the effective nucleon mass at the saturation density. Both values are rather constrained by the atomic nucleus data. The problem originates from the fact that the RMF model parameters are fixed at the saturation density and no free parameters remain to control the pressure and the symmetry energy at higher densities. The problem could be in principle solved by introducing new terms with new fitting parameters.  There are strong indications that meson masses as well as the nucleon mass depend on the nucleon density. A general trend for such a dependence is governed by the property that  masses of all hadrons except Goldstone bosons, like pions and kaons,\\footnote{The kaon effective mass may also decrease with the density, whereas the nature  of  this decrease is different \\cite{NK,G00,KV}.} should decrease, when the nucleon density $n$ increases towards a critical point of the chiral symmetry restoration~\\cite{chiral}. According to the conjecture of Brown and Rho~\\cite{BR} the nucleon mass and the masses of vector $\\om\\,, \\rho$ and scalar $\\sigma$ mesons may obey an approximately universal scaling law (a so-called \"Brown--Rho scaling\") \\be \\frac{m_\\rho^* }{m_\\rho}\\simeq  \\frac{m_\\om^*}{m_\\om}\\simeq \\frac{m_\\sigma^*}{m_\\sigma}\\simeq \\frac{m_N^* }{m_N}=\\Phi_{BR}\\,, \\label{BRm}\\ee where asterisks indicate in-medium masses. According to ~\\cite{BR} $\\Phi_{BR}$ decreases with the increase of the density. In ref.~\\cite{SBMR97} an attempt was made to include density dependent masses in the RMF model in order to incorporate the Brown--Rho scaling, explicitly. A theoretical disadvantage of this model is that the  density dependence, $\\Phi_{BR} \\simeq (1+yn/n_0 )^{-1}$, $y\\simeq 0.28$, was inserted by hands into the model Lagrangian rather than it followed from the corresponding equations of motion. Insertion of a density dependence  into parameters of a mean-field Lagrangian violates conservation of the energy-momentum and spoils the thermodynamic  consistency. Therefore, such models require a special treatment~\\cite{consist}. Ref. \\cite{SBMR97} also realized that the scaling of meson masses (\\ref{BRm}) would result in a dramatic stiffening of the EoS, if one does not simultaneously scale the corresponding coupling constants in a similar way. As an argument for such a procedure, authors relied on the description of the flow in heavy ion collisions. Below we will show that this prescription relies on equivalence of some RMF models that allows to use different interpolating mean fields without changing physical quantities on a mean-field level in homogeneous matter.  In ref.~\\cite{manka} the density dependence of  effective hadronic masses  was incorporated into the mean-field Lagrangian via extra terms describing a hadron interaction with a scalar field. It resulted, however, in \\emph{an increase} of the $\\rho$ and $\\omega$  meson masses with increase of the nucleon density in disagreement  with (\\ref{BRm}). At the same time the $\\sigma$ mass was assumed to be constant. Thus, the inclusion of the hadron mass scaling into RMF models deserves a further consideration. Some other works, cf. \\cite{THIM}, played with a field-dependent parameters of generalized mean field models, permitting changes of both  masses and couplings.  Equivalence between different parameter choices  was not realized in these papers.   Appropriate knowledge of the  EoS of hot, dense, and isospin asymmetrical nuclear matter is essential for understanding the results of heavy-ion collisions. Flow, particle yields, fragmentation may give some constraints on the EoS and the symmetry energy, cf.~\\cite{dan,Gai}. Often these constraints disagree with those obtained from NS physics, e.g., yielding very low DU threshold density. The proper choice of the EoS is particularly important for the CBM (compressed baryon matter) experimental program at the future facility at GSI. Precision measurements of heavy ion collisions are planned in the whole energy range  from $\\sim 1$~GeV/A up to $\\sim 40$~GeV/A.   To  describe correctly the symmetry energy of the nucleon matter one needs to include the $\\rho$-meson field into a mean-field Lagrangian. The $\\rho$-meson field is traditionally assumed to have only one mean-field component, $(\\vec{\\rho\\,}_\\mu)^a=\\delta_{\\mu 0}\\, \\delta^{a 3}\\, \\rho_0^{(3)}$, where $\\mu=0,1,2,3$ is the Lorentz index, and $a=1,2,3$ is the isospin index. This field produces then the expected contribution to the energy density $\\propto (n_p-n_n)^2$\\,, where $n_p$ and $n_n$ are the proton and neutron densities. In ref.~\\cite{v97} it was realized that, if the $\\rho$-meson is introduced as a non-Abelian gauge boson,  e.g., in the framework of the hidden local symmetry model~\\cite{bando83}, then at a density $n>n_{\\rm c}^{\\rho}$ the charged $\\rho$-meson condensation occurs. In a new phase the symmetry energy  grows as $\\propto |n_p-n_n|$ only. The critical density was found to be $n_{\\rm c}^{\\rho}\\sim m_\\rho^{*3}$, being sensitively dependent on the value of the in-medium mass of the $\\rho$-meson ($m_{\\rho}^{*}$). The phenomenon of the charged $\\rho$-meson condensation is similar to the condensation of gauge bosons in QCD~\\cite{migdal} (gluon condensate) and to electro-weak~\\cite{linde} ($W$-boson condensate) sectors of the Standard model.   In this paper we first construct a generalized RMF model with all parameters, being field dependent (secs.~\\ref{sec:grmfm},\\ref{sec:mass}). We derive the energy density functional, and show how various standard cases can be reproduced. We discuss the equivalence relations among some cases. Then we demonstrate shortcomings  of the models with a hadron mass scaling  without a scaling of coupling constants. We show  that the RMF models with the universal scaling for masses and coupling constants can be reduced to the models without any scaling. Such models do not solve, however, the problem of the low $n_{\\rm crit}^{\\rm DU}$. Then we allow for a different scaling of hadron masses and coupling constants in order to generate a model that produces a stiffer EoS at high densities. This model supports a large limiting mass of the NS and rather high critical density for the DU process. On the other hand, it uses  appropriate values of compressibility and effective nucleon mass at the saturation nuclear density. Then we demonstrate another example of the RMF model with a non-universal scaling of masses and coupling constants that fits well the HHJ EoS including the correct density dependence of the proton concentration. In sec.~\\ref{sec:rho} we investigate a possibility for a charged $\\rho$-meson condensation. We demonstrate that in the framework of our generalized RMF models the charged $\\rho$-meson condensation occurs as a second-order phase transition. Emissivity of the NS cooling processes involving the charged rho condensate is estimated in sec.~\\ref{sec:rhocool}. Results are summarized in the concluding section~\\ref{sec:concl}. In Appendix we show how one can introduce inhomogeneous electric potential. Throughout the paper we use units $\\hbar =c=1$. ",
        "conclusions": "\\label{sec:concl}  In this paper we studied generalized relativistic mean field models that include in-medium changes of meson masses and coupling constants. We demonstrated equivalence of a certain class of the models. Practically we showed that the energy density for homogeneous matter depends only on three independent combinations of mass and coupling scaling functions. Often this equivalence is not realized  and one considers some models as different, whereas in reality they are  identical for the description of homogeneous matter on the mean field level.  We showed that a naive modification of solely boson masses would lead to severe drawbacks, like discontinuous changes of mean fields and hadron masses with the density. To get rid off these problems the mass scaling should be essentially compensated  by the scaling of the corresponding coupling constants.  For linear RMF models, i.e., without a non-Abelian  $\\rho$ interaction and a non-linear $\\sigma$ potential $U$, we demonstrated equivalence of the models with universal scaling $g_i^*\\,m_i/g_i\\,m_i^*\\simeq 1$\\,, $i=\\sigma, \\om,\\rho$, and the models without any scaling of those masses and couplings. For non-linear RMF models ($U\\neq 0$) the above equivalence would require a certain relation between corresponding non-linear potentials.   We have demonstrated the efficiency of the generalized relativistic mean field models based on the modified Walecka (MW) model with a non-universal scaling of meson masses and couplings (MW(nu)). It allows to enlarge the threshold density for the direct Urca process and produces a stiffer equation of state without any changes of the EoS near the saturation density. These possibilities are favored  by \"the nuclear medium cooling scenario\" for the neutron star cooling, cf. \\cite{BGV04} and by recent measurements of heavy neutron stars, cf. \\cite{S04}.  Our generalized mean field models with a non-universal scaling of meson masses and couplings might be useful in discussion of possibilities of different phase transitions in neutron star interiors, since they allow to compensate a softening of the equation of state due to a phase transition by a stiffening due to a stronger decrease of masses compared to a decrease of couplings. Especially the model might be helpful in discussion of hybrid stars. With the standard hadron models the quark matter either begins to appear already at rather low neutron star mass, or does not appear at all, cf. \\cite{BGV1}. In the former case the limiting mass of the star is found to be rather low (typically $\\lsim 1.8~M_{\\odot}$). If one used a model assuming a smaller nucleon mass at the saturation, one could get a still  higher value of the limiting neutron star mass.  We also demonstrated an example of the relativistic mean field model that fits well the modern Urbana-Argonne equation of state (more precisely, the fit \\cite{HHJ} that cures the causality problem), yielding the same threshold density for the direct Urca process.  The MW(nu) models can be useful for the description of properties of finite nuclei. Here equivalence is absent due to the spatial inhomogeneity of the field profiles. In standard MW models one is forced to decrease the $\\sigma$ meson mass in order to appropriately fit the proton density profile and the binding energy as a function of the atomic number, cf. \\cite{VYTkaon}. In the framework of the MW(nu) models there is still a room to fit the scaling laws for effective masses and couplings.  It would be interesting to apply the model to the description of heavy-ion collisions. Early naive studies of the pion yield in heavy-ion collisions required extremely hard EoS, with the compressibility $K\\sim 1000$~MeV. Further it was demonstrated, cf. \\cite{MSTV,V83,KV96} and references therein, that the problem might be naturally solved within the standard modified Walecka model (MW) when one includes contributions from  soft pion modes. These modes yield very small contribution to thermodynamical characteristics of the cold nuclear matter but produce a large contribution for the hot nuclear matter. The physical reason is obvious: soft modes are easily excited with the temperature increase. Due to a decrease of $\\sigma$, $\\omega$, $\\rho$ and nucleon masses related modes become to be softer. Thus our mean field models may give an appropriate basis to go beyond the mean field level considering thermal excitations.   Recent studies of the collective isospin flow \\cite{THIM} extracted the density dependence of the symmetry energy but did not introduce any astrophysical constraints and a possible soft mode contribution. It is therefore a challenge to fit the corresponding data using the constructed above models with inclusion of thermal contribution from soft modes.   If one introduces a rho-meson field as a non-Abelian gauge boson (cf.~\\cite{bando83}) in a relativistic mean-field model which supports the decreasing $\\rho$-meson mass, one should take care of the possibility of the charged rho meson condensation in neutron star matter. Following~\\cite{v97}, we found that at density $n>n_{\\rm c}^{\\rho\\,,\\rm II}$ the charged $\\rho$-meson mean field appears. We demonstrated that in the framework of our relativistic mean field models the novel phase  arises by a second order phase transition. The critical density $n_{\\rm c}^{\\rho\\,,\\rm II}$ is rather low in the neutron star matter (typically $\\sim 3\\div 4\\,n_0$), if one allows for a decrease of the rho meson mass. We also included the corresponding decrease of the effective nucleon-rho coupling, as well as the scaling of other meson masses and couplings. If one used a model assuming a smaller nucleon mass at the saturation, one could get a still  smaller value of the critical density for the charged $\\rho$ condensation. Charged $\\rho$ condensation slightly softens the equation of state. In the neutron-enriched matter the charge of the $\\rho$-meson condensate is negative. We evaluated the emissivity of the direct Urca-like processes on the condensate $\\rho^-$ meson. The resulting rate behaves similar to that for the $\\pi^-$ condensate. The analysis of the cooling data indicates that the value $n_{\\rm c}^{\\rho\\,,\\rm II}$ should be not too small, $n_{\\rm c}^{\\rho\\,,\\rm II}\\gsim 2.5\\div 2.7~n_0$.  Our non-Abelian $\\rho$ meson field is in some sense analogous to the gluon field. Dropping of the effective $\\rho$ meson mass enforces a similarity. Effective nucleon and other meson masses also decrease with increase of the density. Decrease of the corresponding effective couplings motivates  a future  asymptotic freedom in the quark-gluon phase. All these features  might be in favor of a smooth transition from the hadron  degrees of freedom to the quark-gluon degrees of freedom. Up to now there is no appropriate description of a cross-over temperature region seen from the lattice results \\cite{Fodor}. Formulated above models might be helpful in construction of such a description.   \\noindent {\\bfseries"
    },
    "0410/astro-ph0410694_arXiv.txt": {
        "abstract": "We present high spatial resolution ($2\\farcs3\\times1\\farcs9$ or 43 pc $\\times$ 36 pc at D = 3.9 Mpc) 100 GHz millimeter-wave continuum emission observations with the Nobeyama Millimeter Array toward an expanding molecular superbubble in the central region of M82. The 100 GHz continuum image, which is dominated by free-free emission, revealed that the four strongest peaks are concentrated at the inner edge of the superbubble along the galactic disk. The production rates of Lyman continuum photons calculated from 100 GHz continuum flux at these peaks are an order of magnitude higher than those from the most massive star forming regions in our Galaxy. At these regions, high velocity ionized gas (traced by H41$\\alpha$ and [\\ion{Ne}{2}]) can be seen, and H$_{2}$O and OH masers are also concentrated. The center of the superbubble, on the other hand, is weak in molecular and free-free emissions and strong in diffuse hard X-ray emission. These observations suggest that a strong starburst produced energetic explosions and resultant plasma and superbubble expansions, and induced the present starburst regions traced by our 100 GHz continuum observations at the inner edge of the molecular superbubble. These results, therefore, provide the first clear evidence of self-induced starburst in external galaxies. Starburst at the center of the superbubble, on the other hand, begins to cease because of a lack of molecular gas. This kind of intense starburst seems to have occurred several times within $10^{6}-10^{7}$ years in the central region of M82. ",
        "introduction": "\\label{intro}  Messier 82 (NGC 3034, Cigar Galaxy) is known as one of the nearby \\citep[3.9 Mpc;][]{sak99}, bright irregular galaxies, and, therefore, has been studied by many authors, at many wavelengths, and over many years. Optical observations \\citep{lyn63,sho98,ohy02} show remarkable features of kpc-scale filaments extending above and below the disk of this galaxy. Soft X-ray \\citep*{wat84,kro85,bre95,tsur97,str97}, molecular gas \\citep*{nak87,mat00,tay01,wal02}, and dust \\citep*{kun97,alt99} observations also show filamentary structures outside the disk region. These observations suggest that the filamentary structures are the result of explosions, which occurred within the nuclear region. The mass-to-luminosity ratio, $M/L$, of this galaxy is extremely low ($<0.02$, \\citealp{rie79}; $\\approx0.04$, \\citealp{tel93}) as compared with normal galaxies \\citep[$5-10$,][]{rie78} and can be explained by bursts of star formation or a starburst \\citep{rie78,rie80}.  The starburst in M82 is believed to be triggered by a tidal interaction with the nearby two-armed spiral galaxy M81. Neutral hydrogen (\\ion{H}{1}) emission observation toward the M81 group shows that there is a tidal bridge between M81 and M82, suggesting that there is a close encounter between these two galaxies $\\sim10^{8}$ years ago \\citep*[e.g.,][]{cot77,yun94}. Optical and near-infrared imaging and spectroscopic observations toward the disk region ($0.5-1$ kpc outside from the galactic center) of M82 indicate that the starburst occurred at this region $\\sim10^{8}$ years ago, corresponds to the episode of the tidal interaction with M81 \\citep*{oco78,gal99,deg01,smi01}. Since the present starburst region is located at the center, they suggested that the global starburst in M82 has propagated inward from the disk to the nuclear regions. The evolution of starburst at the nuclear 500 pc region is, on the other hand, still unclear. Some authors suggest the nuclear starburst propagated outward \\citep[e.g.,][]{sat95,sat97,for03}, and some inward \\citep[e.g.,][]{she95}.  Interferometric molecular gas observations show an arc-like deviation from rigid rotation in position-velocity (PV) diagrams at $\\sim140$ pc westward from the galactic nucleus, and this deviation implies the existence of an expanding molecular superbubble \\citep{nei98,wei99,wil99}. The existence of the superbubble was confirmed by imaging the shell-like structure with the size of 260 pc $\\times$ 170 pc \\citep*{mat00,wil02}. The age and the total energy of the superbubble are estimated to be about $(1-2)\\times10^{6}$ yr and $\\sim(0.5-2)\\times10^{55}$ ergs, which corresponds to the total energy of $\\sim10^{3}-10^{4}$ supernovae. At the center of the superbubble, there is a ``2.2 $\\mu$m secondary peak'' \\citep[a primary peak corresponds to the nucleus;][]{die86,les90}, which is dominated by luminous supergiants \\citep*{joy87}. Since luminous supergiants are the late phase of OB stars, it is reasonable to suppose that a starburst has occurred at this near-infrared peak in the recent past. Actually, calculations of the initial mass function of this peak suggest that the number of supernovae, which have already exploded within the age of the superbubble, can explain the total energy of the superbubble. It is therefore concluded that this superbubble has been produced as a result of localized starburst at the ``2.2 $\\micron$ secondary peak'' \\citep{mat00}.  Detailed information about the nature of starburst in this region is, however, still not clear, because the central region of M82 is deeply obscured by dust even in the K-band (e.g., $A_{\\rm K}\\sim1.4$, \\citealp{mcl93}; $A_{\\rm K}\\sim1.0$, \\citealp{sat95}; $A_{\\rm K}\\lesssim0.8$, \\citealp{tel91}: These correspond to $A_{\\rm V}\\sim14$, $\\sim10$, and $\\lesssim8$, respectively, if we assume $A_{\\rm V}/A_{\\rm K}=10$; \\citealp{dra89}). Hence it is very hard to derive detailed information --- such as the precise flux of recombination lines --- in the optical or near-infrared.  On the other hand, 100 GHz continuum observations are generally dominated by the free-free emission and would not be affected by the dust absorption, as is the case for optical or near-infrared observations. In fact, spectral energy distribution (SED) studies from the infrared to radio wavelength region of M82 indicates that the 100 GHz continuum emission in M82 is dominated by free-free emission \\citep{con92,mat98}. A detailed SED study for small-scale structures toward the central region of M82 by \\citet{car91} based on their millimeter and centimeter interferometric (BIMA and VLA) observations concluded that the free-free emission dominates even in the small-scale structures of 100 GHz continuum emission (i.e., 100 GHz continuum emission from all the peaks are dominated by free-free emission). Hence, 100 GHz continuum observations can have great advantages for the observation of the central starburst regions in M82, where a large amount of dust exists.  In this paper, we report the structures and properties of the present starburst regions and environment in M82 with our high spatial resolution 100 GHz observations and compare them with previously published data. We also discuss the triggering of new starbursts by the past energetic starburst and the resultant superbubble expansion, namely, self-induced starburst. ",
        "conclusions": "\\label{conclude}  Our new high spatial resolution 100 GHz continuum (i.e., free-free emission) image in the central region of M82 clearly shows that present starburst regions are strongly concentrated at the inner edge of the expanding molecular superbubble. The Lyman continuum photon numbers from the starburst regions (i.e., the strength of the starburst) are an order of magnitude larger than those in the active high-mass star forming regions in our Galaxy (NGC 3603 and W51) and nearby galaxies (IC 342 and M51).  The internal structure of the present superbubble, from the center, shows a red supergiant star cluster (``2.2 $\\mu$m secondary peak''), overpressurized hot thermal plasma or high energy electrons (diffuse hard X-ray emission), high velocity ionized gas (H41$\\alpha$ and [\\ion{Ne}{2}]), shocked molecular gas (masers and SiO emission) and the present starburst regions (free-free emission), and the expanding molecular superbubble located furthermost outside. This structure and the physical properties of these objects/mediums suggest that the present starburst has been induced by the past starburst that made the superbubble, namely, self-induced starburst. The existence of kpc-scale outflows, a few hundred parsec-scale superbubble, and induced starburst inside the superbubble suggest that the starbursts in M82 have occurred several times in this $10^{6}-10^{7}$ years. Since the appearance of the starburst (e.g., outflows) in M82 is typical in other starburst galaxies, this starburst mechanism will be very important for the studies of other starburst regions/galaxies and galaxy formation."
    },
    "0410/astro-ph0410007_arXiv.txt": {
        "abstract": "The velocity dispersion of 7 Seyfert 1 galaxies at $z\\sim0.37$ is measured using high signal-to-noise Keck spectra.  Black hole (BH) mass estimates are obtained via an empirically calibrated photoionization method. We derive the BH mass velocity dispersion relationship at $z\\sim0.37$. We find an offset with respect to the local relationship, in the sense of somewhat lower velocity dispersion at a fixed BH mass at $z\\sim0.37$ than today, significant at the 97\\% level. The offset corresponds to $\\Delta \\log \\sigma$ = $-0.16$ with rms scatter of 0.13 dex. If confirmed by larger samples and independent checks on systematic uncertainties and selection effects, this result would be consistent with spheroids evolving faster than BHs in the past 4 Gyrs and inconsistent with pure luminosity evolution. ",
        "introduction": "The correlation of the mass of the central black hole (M$_{\\rm BH}$) with the spheroid velocity dispersion $\\sigma$ (Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000; hereafter BHS) links phenomena at widely different scales (from the pcs of the BH sphere of influence to the kpcs of the bulge).  This demonstrates that galaxy formation and AGN activity are connected and several physical explanations have been proposed (e.g. Kauffmann \\& Haenhelt 2000; Monaco et al.\\ 2000; Volonteri et al.\\ 2003; Haiman, Ciotti \\& Ostriker 2004).  Interesting clues can be obtained from the cosmic evolution of empirical relations. Different scenarios -- all reproducing the local BHS relation -- predict different evolution. For example, in a pure luminosity evolution scenario of spheroids, $\\sigma$ would not change with time, while M$_{\\rm BH}$ would increase as a result of accretion. A typical $10^7 M_{\\odot}$ BH accreting at an average rate $0.01 M_{\\odot} yr^{-1}$ (e.g., Sun \\& Malkan 1989), would double its mass in a Gyr. This would predict the \\mbh\\, at given $\\sigma$ to increase with time, possibly changing the slope of the BHS relation if the accretion rate is a function of M$_{\\rm BH}$ (e.g. Small \\& Blandford 1992). In contrast, if spheroids grew faster than BH, $\\sigma$ at fixed M$_{\\rm BH}$ could increase with time.  In this {\\it Letter} we present the first results from an observational program aimed at measuring the cosmic evolution of the BHS relation. We measure the BHS relation for a sample of 7 Seyfert 1s at $z\\sim0.37$. This redshift is high enough to provide a time baseline over which we might expect evolution, yet low enough to be observationally practical. ",
        "conclusions": "We have measured the velocity dispersion of the hosts of distant active nuclei. In combination with \\mbh\\, estimates using the ECPI method, we have obtained a first estimate of the BHS relation at $z\\sim0.37$. Since M$_{\\rm BH}$ cannot decrease, this measurement, taken at face value, suggests the $\\sigma$ of bulges of Sy1s increased by roughly 40\\% in the past 4 Gyrs. Having provided a first illustration of the method, it is now necessary to collect data for a large sample of objects spanning a larger range of \\mbh\\, and $\\sigma$, and to test systematic uncertainties by obtaining independent measures of \\mbh\\, and spatially resolved information on the bulge and disk morphology and kinematics.  {"
    },
    "0410/astro-ph0410231_arXiv.txt": {
        "abstract": "Accelerating superluminal motions were detected recently by multi-epoch Very Long Baseline Interferometry (VLBI) observations. Here, a Linear Accelerating Superluminal Motion (LASM) model is proposed to interpret the observed phenomena. The model provides a direct and accurate way to estimate the viewing angle of a relativistic jet. It also predicts that both Doppler boosting and deboosting effects may take place in an accelerating forward jet. The LASM model is applied to the data of the quasar 3C 273, and the initial velocity, acceleration and viewing angle of its three components are derived through model fits. The variations of the viewing angle suggest that a supermassive black hole binary system may exist in the center of 3C273. The gap between the inner and outer jet in some radio loud AGNs my be explained in terms of Doppler deboosting effects when the components accelerate to ultra-relativistic speed. ",
        "introduction": "Superluminal motions were first detected by Very Long Baseline Interferometry (VLBI) in 1970s \\citep{coh77}. The observations carried out since then showed that such motions are quite common in blazars. Many models were proposed to explain the phenomenon, which were reviewed by \\citet{bla77}. The relativistic jet model \\citep{bla79} has become the {\\it de facto} standard model to describe the superluminal motion.  From a simple physical consideration, acceleration is required to drive the subrelativistic material in the accretion disk up to the relativistic speed observed in the jets. Acceleration is also necessary to compensate the radiation losses in jets. There is increasing evidence that the central engine of acceleration in AGNs is a rotating supermassive black hole surrounded by a geometrically thin accretion disk, which gives rise to the formation of a pair of relativistic jets. The mechanism of acceleration may be due to the centrifugal and shear effects \\citep{web94,rie02}. In recent years, multi-epoch VLBI monitoring indeed led to the discovery parsec-scale accelerating superluminal motion in the quasar 3C 273 \\citep{kri01} and in other sources \\citep{hom01}.  Currently, quadratic or cubic polynomials are usually employed to fit the accelerating superluminal motion curves \\citep{kri01}. Such mathematical fits, however, have no real physical meaning. Actually, a physical model is required to describe the observed phenomena. Here, we consider a simple model where a component is ejected out collinearly, with intrinsic constant acceleration $g$. Such acceleration may exist in rotating magnetized jets \\citep{rie02}, especially at the outsets of jets where the resistance of the intergalactic medium (IGM) can be neglected. Also, some theoretical work shows that the jets in pulsars indeed move outward with linear acceleration \\citep{con02}. In the frame of special relativity, this model can interpret the accelerating superluminal motions, and we name it as a Linear Accelerating Superluminal Motion (LASM) model.  In this Letter, the LASM model will be introduced in detail. The basic kinematic equations of an accelerating relativistic component and their solutions are listed in Section 2. In Section 3, the proper motion of 3C\\,273 monitored with VLBI is model-fitted, and the viewing angle, initial velocity and acceleration of three components are estimated. Possible application of the LASM model, with emphasis on the estimates of the viewing angle and the evolution of the Doppler factor, are discussed in Section 4.  Throughout this Letter, we adopt a flat, accelerating cosmology with Hubble constant $H_0 = 71 {\\rm km\\, s^{-1}\\, Mpc^{-1}}$, $\\Omega_\\Lambda=0.73$ and $\\Omega_{\\rm m} = 0.27$. ",
        "conclusions": "\\subsection{Viewing angle} The viewing angle plays an important role in the unification for radio-loud AGN. It is believed that all radio-loud AGNs have a similar physical structure: a supermassive black hole (probably spinning) located at their centre, whose gravitational potential energy is the ultimate origin of the AGN luminosity; an accretion disk surrounding the black hole, with relativistic jets emitted perpendicular to it; broad line and narrow line regions, and a dusty torus located at a larger distance from the innermost nucleus. The appearance of AGNs, however, depends so strongly on orientation that our current classification schemes are dominated by random pointing directions instead of more grounded physical properties. Thus, to understand the intrinsic mechanism of the AGNs, beaming effects induced by the viewing angle should be removed \\citep{urr95}.  Besides the  apparent velocity measured by VLBI, other observational properties are needed to estimate the viewing angle. The apparent velocity $\\beta_{\\rm app}=\\beta\\sin\\theta/(1-\\beta\\cos\\theta)$ depends both on $\\theta$ and $\\beta$, so $\\theta$ cannot be derived only from $\\beta_{\\rm app}$. If both apparent velocity $\\beta_{\\rm app}$ and Doppler boosting factor $D$ are known, then the viewing angle can be calculated \\citep{ghi93}. $\\beta_{\\rm app}$ can be directly observed by VLBI. Currently, there are three ways to estimate $D$. The first one is based on synchrotron self-Compton (SSC) model \\citep{mar87}. The second one depends on equipartition brightness temperature \\citep{rea94}. The last one can be derived from the flux density variability of blazars \\citep{lah99}. Compared to the LASM model, all of these three models need extra physical parameters which are probably difficult to assess. For example, SSC model requires the flux at the turnover frequency, which is usually substituted by the observed VLBI frequency in most cases, with an additional source of errors. Furthermore, these three models just provide the average values of the Doppler factors of jets, while the LASM model can trace the evolution of the Doppler factor as well as of the apparent velocity of the individual components.  The LASM model provides a direct and reliable way  to estimate the viewing angle of a relativistic jet, provided that there are enough high angular resolution VLBI monitoring data. As shown in Section \\ref{sec22}, the shape of a proper motion curve is controlled by the viewing angle and other two parameters. Through nonlinear fitting, all the parameters could be evaluated. Furthermore, the acceleration will also influence the observed flux densities of jet components (see Equation \\ref{doppler}). If flux density data  are incorporated into model fitting, the accuracy of the estimation of viewing angle should be improved.  The model fitting results of 3C 273 show that the jet components C11 C12 and C13 have different viewing angles (See Table \\ref{tab1}). This is clear from the  VLBI images \\citep{kri01}, where different jet components move out along projected trajectories with different position angles. The maximum difference of position angles can reach 40 degree. The difference of position angles as well as viewing angles may indicate that there probably exists a super-massive binary black hole system \\citep{sun97} where the orbital motion of the relativistic jet emerging from one massive black hole causes the variation of viewing angles \\citep{rie00}. Another possible explanation is precession \\citep{gow82}, and the precession angle and period can be used to estimate the masses of the black holes \\citep{rom00}.  \\subsection{Doppler factor\\label{sec42}} During the accelerating stage, the evolution of the Doppler factor will affect the flux density variability of a jet component. The observed flux density $S_{\\rm ob}$ of a component is related to its intrinsic flux density $S$ by $S_{\\rm ob}=S D^{3+\\alpha}$ \\citep{bla79}, where $\\alpha$ is the spectral index. A small variation of $D$ will cause a large change of the observed flux density. As seen in Figure \\ref{fig2}, during the acceleration, the evolution of a Doppler factor is divided into two stages. In the first stage, $D$ increases as the intrinsic speed of a component increases. In the second stage, as the intrinsic speed keeps raising, $D$ will gradually decrease, going to values lower than 1. Assuming $\\alpha=0.5$, the observed flux densities of C11, C12 and C13 are boosted up to the maximum level of 1267, 140 and 820 times respectively. If we set $D=1$ as the detection threshold of VLBI, then the traceable lengths of the components C11, C12 and C13 range roughly from 3 mas to 60 mas. This is in agreement with the observations.  These considerations allow us to interpret the gaps between the inner and outer jets of some radio loud AGNs, such as for instance 3C\\,273 and M87. For 3C\\,273, the inner jet is about 100 mas long at 1.67GHz \\citep{dav91}. The outer jet is about $13''$ away from the inner jet and extends about $10''$ \\citep{sam01,mar01}. Inside this gap, no radiation is detected. The possible reason is that continued acceleration of components will cause Doppler deboosting. Therefore, at some distance, these components  will disappear under the detection level of present observation. However, the acceleration will end when the internal energy is mostly converted to kinetic energy.  From then on, due to the interaction between the jets and the medium of their host galaxy at sub-kpc scales, the components will decelerate and and their Doppler factors will increase. Consequently, the Doppler boosted components will appear again in the outer jet region."
    },
    "0410/astro-ph0410117_arXiv.txt": {
        "abstract": "{ In an analysis of \\xmm\\ archival observations of the bright Local Group spiral galaxy \\m31\\ we study the population of X-ray sources (X-ray binaries, supernova remnants) down to a 0.2--4.5 keV luminosity of 4.4\\ergs{34}. EPIC hardness ratios and optical and radio information are used to distinguish between different source classes. The survey detects in an area of 1.24 square degree 856 sources. We correlate our sources with earlier \\m31\\ X-ray catalogues and use information from optical, infra-red and radio wavelengths. As \\m31\\ sources we detect 21 supernova remnants (SNR) and 23 SNR candidates, 18 supersoft source (SSS) candidates, 7 X-ray binaries (XRBs) and 9 XRB candidates, as well as 27 globular cluster sources (GlC) and 10 GlC candidates, which most likely are low mass XRBs within the GlC. Comparison to earlier X-ray surveys reveal transients not detected with \\xmm, which add to the number of \\m31\\ XRBs. There are 567 sources classified as hard, which may either be XRBs or Crab-like SNRs in \\m31\\ or background AGN. The number of 44 SNRs and candidates more than doubles the X-ray detected SNRs. 22 sources are new SNR candidates in \\m31\\ based on X-ray selection criteria. Another SNR candidate may be the first plerion detected outside the Galaxy and the Magellanic Clouds. On the other hand, six sources are foreground stars and 90 foreground star candidates, one is a BL Lac type active galactic nucleus (AGN) and 36 are AGN candidates, one source coincides with the Local Group galaxy M 32, one with a background galaxy cluster (GCl) and another is a GCl candidate, all sources not connected to \\m31. ",
        "introduction": "In the \\xmm\\ survey of the Local Group Sc galaxy M~33, \\citet[][hereafter Paper I]{2004A&A...426...11P} detected 408 sources in a 0.8 square degree field combining the counts of all EPIC instruments. The use of X-ray colours and optical and radio information allowed them to identify and classify the X-ray sources and proved to be efficient in separating super-soft X-ray sources (SSSs) and  thermal supernova remnants (SNRs) in  M~33 from Galactic stars in the foreground and ``hard\" sources. These hard sources may be either X-ray binaries (XRBs) or Crab-like SNRs in M~33 or active galactic nuclei (AGN) in the background of the galaxy.  The Andromeda galaxy \\m31\\ is located at a similar distance as \\me33\\ \\citep[780 kpc,][ i.e. 1\\arcsec\\ corresponds to 3.8 pc and the flux to  luminosity conversion factor is 7.3\\expo{49} cm$^2$]{1998AJ....115.1916H,1998ApJ...503L.131S} and -- compared to the near face-on view of \\me33\\ -- is seen under a higher inclination (78$^{\\circ}$). The optical extent of the massive SA(s)b galaxy can be approximated by an inclination-corrected $D_{25}$ ellipse with large diameter  of 153\\farcm3 and axes ratio of 3.09 \\citep{1991trcb.book.....D,1988ngc..book.....T}. With its moderate Galactic foreground absorption  \\citep[\\nh = 7\\hcm{20}, ][]{1992ApJS...79...77S} \\m31\\ is well suited to study the X-ray source population and diffuse emission in a nearby spiral similar to the Milky Way. As \\m31\\ is seen through a comparable Galactic absorbing column as \\me33\\ one can use the same methods and similar source selection criteria that proved to be successful in the M~33 analysis. On the other hand, there is a big number of low mass XRBs (LMXBs) identified as bright X-ray sources in \\m31\\ from earlier observations (see below). The properties of these sources may help to better classify some of the ``hard\" sources in \\me33.  \\m31\\ was a target for many previous imaging X-ray missions. The \\ein\\ X-ray observatory detected 108 individual X-ray sources brighter than $\\sim$5\\ergs{36}, 16 of which were found to vary between \\ein\\ observations \\citep{1979ApJ...234L..45V, 1990ApJ...356..119C,1991ApJ...382...82T}. The sources were identified with young stellar associations, globular clusters (i.e. LMXBs) and SNRs \\citep[see e.g.][]{1981ApJ...247..879B, 1984ApJ...284..663C}. With the ROSAT HRI, \\citet{1993ApJ...410..615P} reported 86 sources brighter than $\\sim$\\oergs{36} in the central area of \\m31, nearly half of which were found to vary when compared to previous \\ein\\ observations. With a separation of about one year between the \\m31\\ surveys, the ROSAT PSPC covered the entire galaxy twice and detected altogether 560 X-ray sources down to a limit of $\\sim$5\\ergs{35} \\citep[][hereafter SHP97, SHL2001]{1997A&A...317..328S,2001A&A...373...63S}. The intensity of 34 of the sources varied significantly between the observations, and SSSs were established as a new class of \\m31\\ X-ray sources \\citep[see also][]{1999A&A...344..459K}.  With the new generation of X-ray observatories, \\chandra\\ and \\xmm, up to now only parts of \\m31\\ were surveyed. Deep \\chandra\\ ACIS-I and HRC observations of the central region (covered areas of 0.08 and 0.27 square degree) resolved 204 and 142 X-ray sources, respectively \\citep{2000ApJ...537L..23G,2002ApJ...577..738K,2002ApJ...578..114K}. Three \\m31\\ disk fields, spanning a range of stellar populations, were covered with short \\chandra\\ ACIS-S observations to compare their point source luminosity functions to that of the galaxy's bulge \\citep{2003ApJ...585..298K}. A synoptic study of \\m31\\ with the \\chandra\\ HRC covered in 17 epochs ``most\" of the disk (0.9 square degree) in short observations and resulted in mean fluxes and long-term light curves for the 166 objects detected \\citep{2004ApJ...609..735W}. In these observations, several \\m31\\ SNRs were spatially resolved \\citep{2002ApJ...580L.125K,2003ApJ...590L..21K} and bright XRBs in globular clusters and SSSs and quasisoft sources (QSSs) could be characterized \\citep[][]{2002ApJ...570..618D,2004ApJ...610..247D,2004ApJ...610..261G}. During the \\xmm\\ guaranteed time program there were four observations of the central area of  \\m31\\ and three aimed at the northern and two at the southern disk. These observations were used to investigate the bright and variable sources and diffuse emission \\citep[][]{2001A&A...365L.195S,2001A&A...378..800O,2001ApJ...563L.119T,TKP2004}, and to derive source luminosity distributions \\citep{2002ApJ...571L..17T}. In addition, the time variability and spectra of several individual XRBs have been studied in detail \\citep[e.g.][]{2002ApJ...581L..27T,2003A&A...411..553B,2003A&A...405..505B,2004A&A...419.1045M}.  Here we present X-ray images and a source catalogue for the archival  observations of \\m31\\ combining the three \\xmm\\ EPIC instruments and using only times of low background. For  the source catalogue and source population study of these observations we  analyzed the individual pointings and the merged data of the central area simultaneously in five energy bands  in a similar way as described in our \\me33\\ analysis in Paper I. The covered area of 1.24 square degree and limiting sensitivity is a significant improvement compared to the \\chandra\\ surveys, however, up to now only covers about 2/3 of the optical \\m31\\ extent ($D_{25}$ ellipse) with a rather inhomogeneous exposure. ",
        "conclusions": "We present a catalogue of 856 X-ray sources based on archival \\xmm\\ observations covering an area of 1.24 square degree. We correlate our sources with earlier \\m31\\ X-ray catalogues and use information from optical, infra-red and radio wavelengths and in addition their X-ray properties for source classification. We applied a source classification scheme similar to the one successfully used in Paper I for \\me33. As \\m31\\ sources we detect 21 SNRs and 23 SNR candidates, 18 SSS candidates, 7 XRBs and 9 XRB candidates, as well as 27 GlCs and 10 GlC candidates, which most likely are low mass XRBs within the GlC. Comparison to earlier X-ray surveys reveals transients not detected with \\xmm, which add to the number of \\m31\\ XRBs. The number of 44 SNRs and candidates more than doubles the X-ray detected SNRs. 22 sources are new SNR candidates in \\m31\\ based on X-ray selection criteria. Another SNR candidate may be the first plerion detected outside the Galaxy and the Magellanic Clouds. On the other hand, six sources are foreground stars and 90 foreground star candidates, one is a BL Lac type AGN and 36 are AGN candidates, one source coincides with the Local Group galaxy M 32, one with a background galaxy cluster and another is a GCl candidate, all sources not connected to \\m31.  There are 567 sources classified as hard, which may either be XRBs or Crab-like SNRs in \\m31\\ or background AGN. If source variability is detected SNRs may be excluded from source identification. Thirty-eight sources remain unidentified or without classification.  The archival \\xmm\\ \\m31\\ observations allowed us to probe the point source population of the covered area significantly deeper than during the ROSAT PSPC surveys. However, the \\xmm\\ EPIC coverage is rather inhomogeneous and a significant part of the \\m31\\ disks has not been observed at all. To get a full census of the \\m31\\ X-ray point sources it would be very important to fully cover this nearby galaxy with \\xmm\\ observations."
    },
    "0410/astro-ph0410321_arXiv.txt": {
        "abstract": "We present results from a program of optical spectroscopy for 23 nearby galaxies with emission-line nuclei.  This investigation takes advantage of the spatial resolution of the {\\em Hubble Space Telescope} to study the structure and energetics of the central $\\sim 10 - 20$ pc, and the resulting data have value for quantifying central black hole masses, star formation histories, and nebular properties.  This paper provides a description of the experimental design, and new findings from the study of emission lines.  The sample targets span a range of nebular spectroscopic class, from \\ion{H}{2} to Seyfert nuclei.  This data set and the resulting measurements are unique in terms of the sample size, the range of nebular class, and the investigation of physical scales extending down to parsecs.  The line ratios indicative of nebular ionization show only modest variations over order-of-magnitude differences in radius, and demonstrate in a systematic way that geometrical dilution of the radiation field from a central source cannot be assumed as a primary driver of ionization structure. Comparisons between large- and small-aperture measurements for the \\ion{H}{2}/LINER transition objects provide a new test that challenges conventional wisdom concerning the composite nature of these systems. We also list a number of other quantitative results that are of interest for understanding galaxy nuclei, including (1) the spatial distribution/degree of concentration of H$\\alpha$ emission as a function of nebular type; (2) the radial variation in electron density as a function of nebular type; and (3) quantitative broad H$\\alpha$ estimates obtained at a second epoch for these low-luminosity nuclei. The resulting measurements provide a new basis for comparing the nuclei of other galaxies with that of the Milky Way.  We find that the Galactic Center is representative across a wide span of properties as a low-luminosity emission-line nucleus. ",
        "introduction": "The physical properties of galaxy nuclei are of larger interest in light of growing indications that evolution in the central few parsecs is linked to the global properties of the surrounding galaxy.  This connection is evident particularly in correlations between the mass of a central black hole, and bulge mass (e.g., Kormendy \\& Richstone 1995; H\\\"aring \\& Rix 2004) or velocity dispersion (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000).  Star formation in the centers of galaxies is also known to take on unusual forms that may be in symbiosis with a central mass concentration, or the larger structure of the galaxy (i.e., bulge or bar).  A rich phenomenology is seen in the central few parsecs of galaxy nuclei within the Local Group, but our knowledge of nuclei at greater distances is limited by a corresponding reduction in spatial resolution. This problem has particularly afflicted ground-based optical spectroscopic studies, that typically sample nuclei with apertures encompassing hundreds of parsecs or more.  To investigate the structure and constituents of galaxy nuclei on small scales, we therefore undertook a Survey of Nearby Nuclei with STIS (SUNNS), a program of {\\em Hubble Space Telescope} ({\\em HST}\\/), narrow-slit spectroscopy for a sample of nearby galaxies outside the Local Group.  The target objects were selected on the basis of nebular emission in their centers, and are thus ``active'' in some sense, although at low levels that are very common in bright galaxies.  The SUNNS data are appropriate for addressing several major questions, as described in part in our previous papers in this series.  The long-slit measurements of gas kinematics can be used to measure or place limits on the mass of a central black hole (Sarzi et al. 2001, 2002). The spectra coupled with {\\em HST} imaging can be used further to probe the structure of nuclear gas disks, which we find to be influenced in many cases by nongravitational effects (Ho et al. 2002).  The use of small apertures is very helpful in isolating nuclear emission components that are lost in the contamination by circumnuclear sources, as demonstrated by our discovery of unusual broad-line emission in two galaxies (Shields et al. 2000; Ho et al. 2000).  Small apertures are also important for obtaining a description of the stellar population and star formation history within the central tens of parsecs (Sarzi et al. 2005).  In the current paper we present a summary description of the SUNNS experimental design, and focus on new results for the nuclear emission-line characteristics of the sample.  The spatial distribution and excitation properties of nebular emission can be used to draw conclusions about the dominant energy sources in galaxy centers, and as part of our investigation we explore the significance of aperture effects in determining the spectroscopic appearance of galaxy nuclei. Understanding the energetics of emission-line nuclei is of fundamental importance for tracing rates of star formation and black hole growth in these low-power but nearly ubiquitous objects.  The data reported here also provide an improved basis for comparing energetic phenomena and structure in the  Milky Way's nucleus with the centers of other similar galaxies. ",
        "conclusions": "Ground-based spectroscopic surveys indicate that the majority of galaxies host emission-line nuclei, and the SUNNS results provide new insights into the structure and energetics of these nebular sources. The findings presented here indicate that the emitting gas is spatially distributed on scales of $\\sim 100$ pc, with accretion-powered objects showing a significant degree of concentration in the observed emission.  Within these regions, modest radial gradients in optical emission-line ratios are seen; density gradients evidently play an important role in driving these trends, although there is some evidence as well for ionization gradients in the plasma.  LINER/\\ion{H}{2} transition nuclei show only mixed indications of a simple composite spatial structure, contrary to expectations, suggesting that distributed sources of ionization beyond normal O-type stars may be involved in powering these common objects.  These findings provide a framework for gauging the extent to which the center of the Milky Way is a typical galaxy nucleus.  The Galactic Center is very similar to other emission-line nuclei in terms of its nebular luminosity, surface brightness distribution, and ionization. Combining the present results with other published findings, it is additionally clear that the Galactic Center resembles the nuclei of other systems in terms of its metallicity and stellar population, and likewise its black hole mass and probable spectroscopic class, considered in relation to the bulge properties of the host galaxy.  As a laboratory for the study of galaxy nuclei, the center of the Milky Way thus provides a very representative example."
    },
    "0410/astro-ph0410098_arXiv.txt": {
        "abstract": "{ We present observations of CO overtone bandhead emission toward four massive Young Stellar Objects (spectral type O6--B5). The high signal-to-noise ratio $K$-band spectra were obtained with {\\em VLT-ISAAC} at a resolution $\\Delta v$ = 30 \\kms, sufficient to resolve the bandheads, but not the individual $J$-lines.  We are able to explain the shape of the lines by assuming a simple isothermal keplerian disk model seen at different inclinations. The gas temperature ranges from 1500 to 4500~K and the CO column density is between 0.1 and 4 $\\times$ 10$^{21}$ cm$^{-2}$.  The emission probably arises within the first few astronomical units of the disk, consistent with the high gas temperature.  Our results indicate that molecules can survive close to a hot star and suggest that dense ($n_{\\rm H}>10^{10}$ cm$^{-3}$) inner disks may be relatively common at an advanced stage of high-mass star formation. ",
        "introduction": "Disks are found around many young low-mass stars as a natural byproduct of angular momentum conservation during the star formation process. Observational evidence for disks around their more massive counterparts remains sparse, although they may be essential during the formation of high-mass stars \\citep{Churchwell02}.  Extended disks have been found around massive young stellar objects (YSOs) in the Hot Core phase \\citep[e.g.][]{Shepherd01Science,Beltran04,Garay99,Chini04}, albeit those millimeter observations solely probe gas beyond 1000~AU. In the later Herbig Be stage, disks have only been recently discovered \\citep{Fuente03}.  Evidence for the presence of a gaseous disk around the early B star \\object{NGC2024-IRS2} is presented by \\citet{Lenorzer04}.  They argue that the infrared excess observed toward \\object{NGC2024-IRS2} is best explained by the emission from a dense gaseous disk as close as 0.3~AU to the star.  At the high densities and temperatures (1000--2000~K) characteristic of disks at 0.1--5~AU around young stars, molecules are expected to be sufficiently excited to produce a rich ro-vibrational spectrum in the near- (overtone transitions $\\Delta v = \\pm2$) and mid-infrared (fundamental transitions $\\Delta v = \\pm 1$). CO first overtone emission is most likely emitted by a disk and/or a wind and has been detected and analyzed by, e.g., \\citet{Scoville79,Geballe87,Carr89,Chandler93,Chandler95} toward a number of other high- and low-mass stars.  Their studies were, however, hampered by the small number of observed $J$-lines or a low resolving power and moderate signal-to-noise.  \\citet{Brgspec04} examine $K$-band spectra of  a sample of massive YSOs associated with ultra-compact H{\\sc ii} regions (UCHII). Here, we present a simple disk model that describes the CO bandhead emission detected in four of these objects. The observations are presented in Sect.~\\ref{sec:obs}. The disk model and the results are shown in Sect.~\\ref{sec:model}. In Sect.~\\ref{sec:discussion} we discuss the location and origin of the CO. The implications of our findings on the process of high-mass star formation are summarized in Sect.~\\ref{sec:conclusion}.  \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{Ge272_f1.ps}} \\caption{$K$-band observed (upper spectra) and modeled (lower spectra) spectra of the CO 2--0 bandhead. The name of the object and the derived inclination is displayed.  By definition $0 \\degr$ inclination corresponds to a disk seen pole-on. } \\label{fig:COfits} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion}  We have modeled the CO first-overtone bandhead emissions of 4 massive YSOs. The different profiles of the CO bandheads are all well explained by a simple keplerian disk. Regardless of the gas excitation mechanism, the information on the gas kinematics given by the shape of the CO bandheads shows that the CO is located close to the star (0.1--5~AU).  The CO overtone emission likely arises in a narrow torus where the gas temperature is between 1500 and 4500~K. This range of temperature reflects the difference in  effective temperature of the central objects. The decrease in excitation temperature for one object from the $v=$~2--0 bandhead to the $v=$~5--3 bandhead indicates that the higher $v$-levels are sub-thermally populated. Excitation mechanisms other than LTE, like UV or infrared pumping and photo-electric heating cannot be excluded by the simple model used here.  The detected disks may be the remnants of massive accretion disks. In particular, it is the first evidence of a disk around a young O6--O8 star, suggesting that these objects are recently formed as the disks are not yet dispersed by the stellar FUV photons and stellar wind.   Spatially resolved observation of CO bandhead emissions with {\\em Amber} at the {\\em VLT} interferometer will provide further constraints on the size and the geometry of the CO emitting region."
    },
    "0410/astro-ph0410501_arXiv.txt": {
        "abstract": "Distant supernovae have been observed to be fainter than what is expected in a matter dominated universe. The most likely explanation is that the universe is dominated by an energy component with negative pressure -- dark energy. However, there are several astrophysical processes that could, in principle, affect the measurements and in order to be able to take advantage of the growing supernova statistics, the control of systematic effects is crucial. We discuss two of these; extinction due to intergalactic grey dust and dimming due to photon-axion oscillations and show how their effect on supernova observations can be constrained using observed quasar colours and spectra. For a wide range of intergalactic dust models, we are able to rule out any dimming larger than 0.2 magnitudes for a type Ia supernova at $z=1$. The corresponding limit for intergalactic Milky Way type dust is 0.03 mag. For the more speculative model of photons mixing with axions, we find that the effect is independent of photon energy for certain combinations of parameter values and a dimming as large as 0.6 magnitudes cannot be ruled out. These effects can have profound implications for the possibility of constraining dark energy properties using supernova observations. ",
        "introduction": "Observations of distant type Ia supernovae (SNIa) show that they are fainter than what the theory predicts for a universe described by a standard flat cold dark matter model. The most popular explanation is that the universe not only consists of luminous and dark matter, but also of large amounts of dark energy, a generic name for a homogeneously distributed energy component with negative pressure, $p<-\\rho /3$\\footnote{We do not consider any fluctuations of the dark energy component in this paper.}. This is also verified by other cosmological probes such as the power spectra obtained from the temperature fluctuations in the cosmic microwave background \\cite{Bennett:2003bz,Spergel:2003cb} and (on smaller scales) from large scale galaxy surveys \\cite{2dFGRS,Tegmark:2003ud,Tegmark:2003uf}.  Several mechanisms have been suggested in order to explain the faintness of distant SNIa without having to introduce a dark energy component. Since the existence of a large dark energy component has been independently detected by a number of cosmological probes, it is unlikely that all the dimming is caused by dust or photon-axion oscillations. However, it is crucial to control all possible systematic effects in current and future SN surveys focussing on constraining the time evolution of the dark energy component.  {\\em Supernova evolution:} Some of the apparent dimming of distant SNIa could be due to the fact that their luminosities are redshift dependent and that the SNe thus cannot be used as standard candles. This could, e.g., be caused by evolution of the SN progenitor \\cite{Dominguez:1998jt,Riess:1998cb}. Evolutionary effects can be searched for by comparing the light curve shape and spectra of SNe at different redshifts and environments. Most studies show that the luminosity properties are rather similar \\cite{Perlmutter:1998np,Sullivan:2002ub}, but others have found variations indicating evolutionary effects \\cite{Dominguez:1998jt,Hoeflich:1998uk,Riess:1999ti}.  {\\em Gravitational lensing:} Since matter is inhomogeneously distributed in the universe, most line of sights to distant objects will probe a matter density significantly lower than the average density and the corresponding images will be demagnified relative to a homogeneous universe. It is generally accepted that gravitational lensing cannot be responsible for all the observed dimming of distant SNIa, however the effect may still lead to a non-negligible bias that need to be corrected for, e.g., by using information from the magnitude distribution \\cite{Amanullah:2002xh}. It has also been suggested that the large scale expansion of the universe can be affected by the inhomogeneities in the matter distribution. However, perturbative calculations show that this effect is probably negligible \\cite{Kolb:2004,Rasanen:2004}.  {\\em Photon-axion oscillations:} It has been proposed that the photon flux from SNIa may be diminished because of photons oscillating into axions in much the same way as the electron neutrino ($\\nu_e$) flux, produced in the sun, is diminished because of oscillations into muon neutrinos ($\\nu_{\\mu}$) and tau neutrinos ($\\nu_{\\tau}$). Due to differences in methods and assumptions regarding the magnetic field strengths and plasma densities, different groups have reached different conclusions about the possible magnitude and the energy dependence of the effect \\cite{Csaki:2001yk,Deffayet:2001pc,Csaki:2001jk,Mortsell:2002dd, Christensson:2002ig,Csaki:2004}.  {\\em Intergalactic grey dust:} It has been suggested that intergalactic dust could explain the observed faintness of high redshift SNIa \\cite{Aguirre:1998ge, Aguirre:1999dc}. If there is dust along the line of sight, it would extinguish some of the SN light through absorption and scattering. The intergalactic dust has to be ``grey'' in order not to cause a significant reddening that would have been discovered if it existed.  In this paper, we are going to discuss intergalactic dust extinction and photon-axion oscillations and show how the effects can be constrained using quasars in the first data release (DR1) from the Sloan Digital Sky Survey (SDSS). Similar studies have been performed using quasars from the SDSS early data release (EDR) for intergalactic dust \\cite{Mortsell:2003vk} and for photon-axion oscillations \\cite{Mortsell:2003ya}. For intergalactic dust, a maximum allowed amount of dimming of 0.2 magnitudes for SNIa at $z=1$ was found. In this paper we redo the analysis using the larger and higher quality data set of DR1. We also extend the range of dust models studied. In the analysis of photon-axion oscillations, there was an erroneous factor involving the photon energy in the expression for the density matrix evolution. We redo this analysis using the larger data set of DR1 with the correct expression for the density matrix evolution.  In Section \\ref{sec:oscillations} we constrain the effects from photon-axion oscillations and in Section \\ref{sec:dust} we discuss limits on the intergalactic dust attenuation. In Section \\ref{sec:summary} we summarize our results. ",
        "conclusions": "\\label{sec:summary}  Observations of high redshift SNIa show that the universe is dominated by dark energy accelerating the expansion rate. In order to take advantage of the increased SN statistics and to be able to use SNIa data to determine the dark energy properties, it is necessary to control the systematic effects involved. We have used quasar colours and spectra from the SDSS DR1 to constrain the effects from intergalactic dust attenuation and photon-axion oscillations, respectively.  We have found that the largest possible dimming of a SNIa at $z=1$ caused by intergalactic dust is 0.2 magnitudes. The addition of new quasars to the intergalactic dust analysis will not improve the limits since the analysis is limited by systematic errors, in particular the template errors. One way of improving the template errors is to use different template spectra for different groups of quasars. For example, a median spectrum created from low redshift quasars will be unaffected by possible intergalactic dust and photon-axion oscillations. If the template errors were half their size, we would be able to give a maximum dimming limit of 0.05 magnitudes for a SNIa at $z=1$. Also, improving the resolution of the FIRB experiments could lead to the rejection of more dust scenarios.  The calculations of the dimming caused by intergalactic dust is based on mean extinction laws for the dust in the Milky Way \\cite{Cardelli:1989fp,Fitzpatrick:1999}. It is not known if these parametrisations are valid for an intergalactic dust population where $R_V$ might be much larger than in the interstellar medium. Also, we have assumed that the reddening parameter, $R_V$, is independent of direction and redshift. However, this is a fair approximation since the simulated SN attenuation is almost independent of $R_V$ (see e.g. figure~\\ref{fig:dust_ftz_uzri}).  For photon-axion oscillations we have found a maximum dimming of 0.1 magnitudes for SNIa at $z=1$ if $B\\omega/(M_{\\rm a}n_e) \\lesssim 9 \\cdot 10^{-21}{\\rm ~cm^3G}$, where $B$ is the magnetic field strength, $\\omega$ is the photon energy, $M_{\\rm a}$ is the inverse coupling between the photon and the axion, and $n_e$ is the electron density. However, when $B\\omega/(M_{\\rm a}n_e) \\gtrsim 9 \\cdot 10^{-21}{\\rm ~cm^3G}$ (corresponding to $B\\omega/n_e \\gtrsim 8\\cdot 10^{-2}~{\\rm eVcm^3G}$ using the CAST lower limit on $M_{\\rm a}$ \\cite{Andriamonje:2004}), the oscillation probability is independent of photon energy and a dimming as large as 0.6 magnitudes cannot be excluded. We have also shown that for certain combinations of $n_{e}$ and $B/M_{\\rm a}^{11}$, specifically $B\\omega/(M_{\\rm a}n_e) \\gtrsim 9 \\cdot 10^{-21}{\\rm ~cm^3G}$ and $B_0/M_{\\rm a}^{11}\\sim 10^{-12} - 10^{-11}{\\rm G}$, the attenuation from oscillations evolve with redshift in a manner very similar to the observed SNIa dimming. The region where the effect is independent of photon energy could be constrained using the observed dispersion in SNIa magnitudes. An inhomogeneous magnetic field strength and/or electron density will cause different lines of sight to correspond to different phases in the photon-axion oscillation. Thus we expect a dispersion in integrated broad band magnitudes of high redshift sources if oscillations are present. We have shown that it is not possible to use current data to obtain any further constraints on this region. However, a future SNIa sample could yield better constraints."
    },
    "0410/astro-ph0410671_arXiv.txt": {
        "abstract": "We report the results of \\chandra\\ and \\xmm\\ observations of a new outburst of an ultraluminous supersoft X-ray source in M101. \\sss\\ was observed in a low luminosity state ($L_X\\sim 10^{37}$\\lum) between 2004 January and May. The low-state X-ray spectra were relatively hard; the combined low state spectrum can be fitted with a combination of a power-law with photon index of 1.4 and a blackbody of 63 eV. During 2004 July, the source underwent a strong outburst and the peak 0.3--7 keV luminosity reached $3\\times10^{40}$\\lum, with a bolometric luminosity of about $10^{41}$\\lum. The outburst spectra were very soft and can generally be fitted with a blackbody model with temperatures of 50--100 eV. In two of the observations, absorption edges at 0.33 keV, 0.56 keV, 0.66 keV, and 0.88 keV were found. An \\xmm\\ observation was also performed during the decay of the outburst, and a power-law tail was seen in addition to the supersoft spectrum. We consider different accretion models; one involving an intermediate-mass black hole can explain the observations. ",
        "introduction": "Recent high resolution X-ray observations reveal that there are large number of ultraluminous X-ray sources (ULXs) in many nearby galaxies. ULXs are usually defined as very bright ($L_X \\approx 10^{39-41}$\\lum) off-nucleus, variable X-ray sources in galaxies. Since the luminosity exceeds the Eddington luminosity of a $10M_\\odot$ black hole (BH), ULXs are considered to be candidates for intermediate-mass black holes (IMBHs). More recently, a number of ULXs have been found to be supersoft; their X-ray spectra can be fitted with a blackbody model with temperature of less than 100 eV (e.g., Swartz et al. 2002; Mukai et al. 2003; Fabbiano et al. 2003; Di\\,Stefano \\& Kong 2003; Kong \\& Di\\,Stefano 2003).  In this Letter, we report a series of \\chandra\\ observations of the ultraluminous SSS in M101 (\\sss) from the low state (Kong \\& Di\\,Stefano 2004) to outburst (Kong 2004a, 2004b). Shortly after the discovery of the outburst, we initiated a multiwavelength campaign to monitor the source.  We here also present the results of an \\xmm\\ follow-up observation conducted about two weeks after the first outburst observation. The primary result of radio observations by the Very Large Array was reported by Rupen, Sjouwerman, \\& Kong (2004). We will present the detailed analysis of the optical and radio observations in a follow-up paper. ",
        "conclusions": "Although over $100$ ULXs have been discovered (Miller \\& Colbert 2003), only a handful of sources have $L_X\\gaeq 10^{41}$ ergs s$^{-1}$, assuming isotropic radiation (Matsumoto et al. 2001; Gao et al. 2003; Davis \\& Mushotzky 2004; Soria \\& Motch 2004). A luminosity this high is difficult to achieve in an X-ray binary unless the accretor has a mass greater than roughly $10\\, M_\\odot.$ We do have evidence that \\sss\\ is likely to be an X-ray binary, since its luminosity has been observed to change by a factor of $\\gaeq 10$ on a time scale of hours (Mukai et al. 2003; Di\\,Stefano \\& Kong 2003). While it is possible that our luminosity estimates are higher than the true luminosity, the effects that lead to overestimates, such as beaming, or various anisotropies, tend to change estimated luminosities downward by a factor of roughly $10.$ Its unusually high bolometric luminosity ($1.2\\times 10^{41}$ ergs s$^{-1}$), coupled with its short-time-scale time variability, therefore make \\sss\\ a good candidate for an accreting IMBH.  For instance, King et al. (2001) proposed that anisotropic X-ray emission can result super-Eddington luminosity for a stellar-mass BH. However, in order to achieve such a high luminosity for a stellar-mass BH, extreme beaming is required and the disk is expected to be much hotter. Similarly, the temperature of radiation pressure-dominated accretion disk model proposed by Begelman (2002) is too high. The pure blackbody spectrum also makes relativistically beaming (K{\\\"o}rding et al. 2002) unlikely.  Its soft spectrum in the high state is another important piece of evidence. An accretion disk around a very massive BH is expected to produce supersoft X-ray emission (see Di\\,Stefano \\& Kong 2002). For instance, if the supersoft emission comes from the inner disk, a 70 eV SSS at luminosity of $10^{40}$ \\lum\\ corresponds to a $\\sim 6000~ M_{\\odot}$ BH. If we used the 90\\% lower limits of the inner disk temperature derived from the disk blackbody fits, the BH mass is estimated to be $> 2800 M_\\odot$ (Makishima et al. 2000). This is consistent with prior work on IMBH models for ULXs (see e.g., Miller \\& Colbert 2003). It also complements other work on evidence for cool disks in ULXs that has also been considered as evidence for IMBH models (Miller et al. 2003,2004) although a power-law component usually contributes a significant fraction of X-ray emission. Furthermore, the high state luminosity is approximately $16\\%$ of the Eddington luminosity for a $6000\\, M_\\odot$ BH; we therefore expect the inner disk to be optically thick, which is consistent with the IMBH interpretation.  \\begin{table*} \\centering{ \\small \\caption{Best-fitting Blackbody Parameters for the High State} \\begin{tabular}{lcccccccccc} \\hline \\hline Date & $N_H$ & $kT$ & E$_{edge 1}$ & $\\tau_1$& E$_{edge 2}$ & $\\tau_2$ & $\\chi^2_{\\nu}/$dof & $L_{0.3-7}$ & $L_{bol}$ \\\\ & ($10^{21}$ cm$^{-2}$) & (eV) & (keV) & & (keV) & & & ($ 10^{40}$\\lum) & ($10^{40}$\\lum)\\\\ \\hline Jul 5  & $2.2^{+0.2}_{-0.2}$ & $70^{+15}_{-15}$ &--- &--- &---&---&0.90/8 & 0.74 & 1.97\\\\ Jul 6  & $2.9^{+0.2}_{-0.2}$ & $104^{+22}_{-22}$ & $0.66^{+0.02}_{-0.02}$ &$1.6^{+0.6}_{-0.6}$ & $0.88^{+0.03}_{-0.03}$ & $2.8^{+1.9}_{-1.9}$ & 0.97/26 & 1.04 & 1.82\\\\ Jul 7  & $3.8^{+2.8}_{-2.8}$ & $64^{+17}_{-17}$ & ---&--- &--- &--- & 0.98/14 & 3.82 & 12.40\\\\ Jul 8  & $1.4^{+0.6}_{-0.6}$ & $56^{+4}_{-4}$ & $0.33^{+0.05}_{-0.05}$ & $5.3^{+4.0}_{-4.0}$ & $0.57^{+0.03}_{-0.03}$ & $1.2^{+0.6}_{-0.6}$ & 1.19/17 & 0.81 & 8.71\\\\ Jul 11 & $1.2^{+1.3}_{-1.0}$ & $59^{+11}_{-10}$ & ---&---  &--- &--- & 1.26/10 & 0.17 & 0.63\\\\ \\hline \\end{tabular} } \\end{table*}  The state changes of \\sss\\ span a luminosity range larger than a factor of $1000.$ \\sss\\ was first detected by \\rosat\\ in 1996 November (Wang et al. 1999), with an extrapolated unabsorbed 0.3--7 keV luminosity of $8\\times10^{38}$ \\lum\\ (blackbody model with $N_H=10^{21}$ cm$^{-2}$, $kT=75$ eV). Higher absorption ($3\\times10^{21}$ cm$^{-2}$) would imply a luminosity up of $4\\times10^{39}$\\lum. The source was not detected in other \\rosat\\ observations (see Wang et al. 1999). In a 108 ks \\rosat\\ HRI observation taken in 1996 May/June, we can set a $3\\sigma$ detection limit at $6.8\\times10^{37}$\\lum\\ (power-law model with $N_H=1.2\\times10^{20}$ cm$^{-2}$ and $\\alpha=2$). Setting the $N_H$ at $1.5\\times10^{21}$ cm$^{-2}$, the luminosity limit becomes $1.5\\times10^{38}$\\lum. In 2000 March and October, \\chandra\\ detected the SSS in the very soft state with luminosities of $\\sim 10^{39}$\\lum\\ (Mukai et al. 2003; Di\\,Stefano \\& Kong 2003). In 2002 June, \\xmm\\ observed M101 and did not detect the SSS (Jenkins et al. 2004). We re-analysed the image and found that the SSS may be barely detected with possible contamination from nearby sources. We derived the 0.3--7 keV luminosity of the SSS as $\\sim 10^{37}$\\lum, by assuming a power-law model. The source was also in the low state between 2004 January and May with luminosity of $\\sim10^{37}$\\lum. It is clear that the SSS has had at least three major outbursts ($L_X > 10^{39}$\\lum). On the other hand, the very low luminosities ($\\sim 10^{37}$\\lum) during 2002 and 2004 indicate that the source was in the low state. The source varies by as much as a factor of 1000 between the low state and the high state. This amplitude is even greater than many Galactic BHs. Remarkably, the source also shows spectral changes. It is very clear that there is a power-law component ($\\alpha=1.4$) in the composite low state spectrum, while the high state spectra are supersoft. Our \\xmm\\ observation taken during the decay of the outburst revealed that the blackbody component was still strong and there was a very hard power-law tail, similar to the low state. While it is difficult to make direct comparison with Galactic BHs because of the blackbody emission in the low state and the pure blackbody spectrum in the high state (see however Zdziarski \\& Gierlinski 2004), it is unusual for a SSS to have high/very soft state and low/hard state transition. With a peak bolometric luminosity $\\approx 10^{41}$\\lum, it is likely that IMBH is the central engine of the system. Interestingly, the blackbody component is always seen while the power-law component becomes stronger when the source is at lower luminosities. It may indicate that the power-law component is due to Comptonization of soft photons (Wang et al. 2004). The photon index is, however, harder than that of Galactic BHs in the low state ($\\alpha\\sim1.7$). None of the Galactic BHs has similar spectrum (McClintock \\& Remillard 2003). The closest example is Galactic microquasar V4641 Sgr for which the photon index was measured to be between 0.6 and 1.3 (Miller et al. 2002; McClintock \\& Remillard 2003).  Another important feature of our X-ray spectral fits is the presence of absorption edges, which we found at 0.33, 0.57, 0.66, and 0.88 keV in two of the high state spectra. These are consistent with C {\\small V}, N {\\small VI}, N {\\small VII}, and O {\\small VIII} edges. We note, however, that the 0.33 keV edge may be due to calibration of the ACIS-S near the carbon edge around 0.28 keV. These features may signal the presence of highly ionized gas in the vicinity of the accretor (e.g., warm absorber), and may be consistent with an outflow from the source. In fact, outflow models have been suggested for ULXs. King and Pounds (2003) suggested that the soft component of some ULXs is due to outflow of a stellar-mass BH. Mukai et al. (2003) also used similar argument to explain the super-Eddington luminosity of \\sss. However, the new data which reveal the changes of temperature and bolometric luminosity, the extremely high luminosity, and the state transition of the source will be difficult to explain by such a model. It remains a puzzle that we saw different edges in the two observations, possibly related to the geometry of the system. Nevertheless, outflows from IMBHs may be expected.  It is also interesting to note that similar absorption edges are expected and have been observed in white dwarf (WDs) systems such as the recurrent nova U Sco (Kahabka et al. 1999), CAL~87 and RX~J0925.7--4758 (Ebisawa et al. 2001). SSS radiation  with luminosities in the range from $10^{37}$ ergs s$^{-1}$ to roughly the Eddington limit for a $1.4\\, M_\\odot$ WD are expected for nuclear burning WDs (van den Heuvel et al. 1992). Indeed, such models have been proposed to explain SSSs in the Galaxy and Magellanic Clouds. Note, however, that neutron  star and BH models have also been proposed. The high state luminosity of this source is almost two orders of magnitude larger than that of any other known SSS, and it rules out steady nuclear burning on hot WD models as an explanation for \\sss . A similar argument can be made for neutron star models. A nova explosion could explain some features of the data. It is not favored, however, because the photospheric radius and lack of hard radiation at the peak of the outburst have no obvious explanation.  It may be impossible to conclusively establish, using current technology, that any X-ray source in an external galaxy is an IMBH. A system like \\sss\\ is therefore particularly valuable, because it provides $4$ different pieces of evidence that together make a consistent argument in favor of an IMBH interpretation. Its high-state luminosity, its short-time-scale variability, its soft high-state spectrum, and its pronounced spectral changes, all suggest that \\sss\\ is a strong IMBH candidate."
    },
    "0410/astro-ph0410392_arXiv.txt": {
        "abstract": "We consider the optimum depth of a cluster survey selected using the Sunyaev-Zel'dovich effect. By using simple models for the evolution of the cluster mass function and detailed modeling for a variety of observational techniques, we show that the optimum survey yield is achieved when the average size of the clusters selected is close to the size of the telescope beam. For a total power measurement, we compute the optimum noise threshold per beam as a function of the beam size and then discuss how our results can be used in more general situations. As a by-product we gain some insight into what is the most advantageous instrumental set-up. In the case of beam switching observations one is not severely limited if one manages to set the noise threshold close to the point which corresponds to the optimum yield. By defining a particular reference configuration, we show how our results can be applied to interferometer observations.  Considering a variety of alternative scenarios, we discuss how robust our conclusions are to modifications in the cluster model and cosmological parameters. The precise optimum is particularly sensitive to the amplitude of fluctuations  and the profile of the gas in the cluster. ",
        "introduction": "\\label{intro}  The Sunyaev-Zel'dovich (SZ) effect is due to the re-scattering of  cosmic microwave background (CMB) photons by the hot electrons situated in intracluster media~\\cite{SZ}. It has been realized that it should provide a reliable mechanism for selecting a sample of clusters~\\cite{HMH,HHM} and a number of instruments sensitive to the microwave frequency range (15GHz-350GHz) are under construction to perform this task~\\cite{AMI,SZA,AMIBA,BOLOCAM,OCRA,SPT,ACT}. Since the cluster mass function is sensitive to cosmological parameters~\\cite{TC,OB,SCO,B,VL,ECF} such surveys are likely to be an important part of the program to probe the late-time history of the universe, searching for the origin of the apparently accelerating expansion of the Universe.  The number of clusters as a function of redshift and flux is a strong function of $\\Omega_{\\rm m}$, the mass density of the universe in units of the critical density, and $\\sigma_8$, the normalization of the matter power spectrum on scales of $8h^{-1}{\\rm Mpc}$, and is a weaker function of the Hubble constant $H_0=100h\\,{\\rm km}\\,{\\rm sec}^{-1}{\\rm Mpc}^{-1}$, the equation of state $w_0$ of the dark energy, Q, and its evolution $w_1$. Here, $w_0$ and $w_1$ are defined by $P_{\\rm Q}/\\rho_{\\rm Q}=w=w_0+w_1z$. It has been shown that the degeneracies within the space of cosmological parameters are such that cluster surveys are complementary to those which might be obtained using observations of Type Ia Supernovae at high redshift~\\cite{WBK,BW}.  Since the distribution of the clusters obeys Poisson statistics the errorbars deduced on the cosmological parameters will be very sensitive to the number of clusters found in a given survey. Therefore, we anticipate that finding the largest number of clusters possible for a given integration time  will be close to giving the best constraints on cosmological parameters, although this may also weakly depend on their distribution with redshift. Clearly a survey which will find $\\sim 1000$ clusters will constrain cosmological parameters much better than one which finds $\\sim 100$. Having a much larger sample would also seem desirable in terms of reducing the systematic errors due to the poorly understood gas physics. Fortunately, we will show that, by  assuming we can model the gas physics well, one can get to within a factor of two of the optimal situation for a total power measurement. This is more tricky in more complicated experimental situations such as beam switching whereby an ill-judged selection of the survey depth can lead to one finding very few clusters at all.  The goal of this paper, therefore,  is to discuss issues related to optimizing the yield of a given survey; the optimum being defined as the survey in which the largest number of objects are found in a given amount of observing time. This is a question one is always faced with in any kind of astronomical measurement: how long does one integrate for on a given patch of sky? Obviously the shorter the required integration time for individual patches, the more sky area can be covered. However, if one doesn't integrate long enough only the very brightest objects will be found. Clearly, one is always guided in such situations by one's expectations of what is there to be observed.  For the present case of clusters, one can attempt to make a simple estimate by assuming (1) the clusters are point-like objects in the telescope beam, (2) their comoving number density is constant and (3) the Universe is Euclidean. In this case the number of objects with flux density, $S$, greater than the limiting flux density of a survey $S_{\\rm lim}$ is $N(S>S_{\\rm lim})\\propto \\Delta\\Omega/S^{3/2}_{\\rm lim}$, where $\\Delta\\Omega$ is the angular coverage. Since for fixed total integration time $S_{\\rm lim}\\propto \\sqrt{\\Delta\\Omega}$, one can deduce that $N\\propto \\Delta\\Omega^{1/4} \\propto S_{\\rm lim}^{1/2}$, and hence a survey with largest possible area (that is, the whole sky) would be optimal in this case. Under these assumptions the survey would be flux density limited and hence an instrument with the largest possible collecting area would be optimal in terms of making the largest number of detections in the shortest period of time.  Unfortunately, in most cases clusters are unlikely to be point-like, nor is their comoving density likely to be constant due to evolution and, of course, the Universe is not Euclidean. As we will describe in section~\\ref{distn}, the distribution of clusters as a function of mass and redshift can be modelled in a sensible way for a given set of cosmological parameters. However, the key to understanding the issue at hand here is how the limiting mass of a survey is related to the instrumental noise and the angular resolution of the telescope.  Shallow surveys, as suggested by the simple estimate above, will only find clusters which are very large and close by, and such clusters cover many beam areas. Since the same signal is spread over a number of beam areas, such objects can dip below the detection threshold  even though they are very large. It is clear that this effect will compete with that due to Euclidean source counts in our initial estimate, leading to an optimum at a particular value of the noise and angular coverage. We will show that this optimum will depend on the size of the beam: larger beam sizes, say $8^{\\prime}$, will require a deeper\\footnote{We define deep to mean a low noise level in terms of temperature.} survey to achieve the optimum, whereas a much shallower survey will be optimal for much smaller beam sizes, say $1^{\\prime}$. Essentially the optimal situation is for most of the clusters in the survey to be the same size as the beam as explained in Fig.~\\ref{fig:beams}.  \\begin{figure} \\setlength{\\unitlength}{1cm} \\centerline{\\psfig{file=beamcluster.eps,width=8cm,height=8cm}} \\caption{The top row is the negative cluster profile for an observation frequency of less than 217GHz. The bottom row are three hypothetical beam profiles: (a) the beamsize is much smaller than the cluster and the observation resolves it. In this case the observation is surface brightness temperature limited. (b) the beamsize exactly fits the cluster profile. This is the optimal situation. (c) the beamsize is larger than the cluster profile. In this case the observation is flux density limited.} \\label{fig:beams} \\end{figure}  Once the beam is sufficiently small to resolve all the clusters, the survey becomes surface brightness temperature limited, rather than being flux density limited, and increasing the collecting area no longer acts to  increase the rate of detection. This illustrates an interesting feature of cluster survey, which is in contrast to intuition one might have from considering standard point sources and, as we shall see, leads to some other somewhat counter-intuitive conclusions.  In section~\\ref{mass} we will discuss how to derive a relation for the limiting mass of the survey as a function of redshift and the limiting flux, incorporating the effects of a Gaussian beam profile. Our initial motivation is to model idealized total power measurements using a single dish with a circular aperture. We will allow the observations to be smoothed using a window function which is exactly tuned to every possible cluster size above the maximum resolution of the telescope. Under these assumptions we show in section~\\ref{total} that for a particular minimum beam size there exists a unique value of the noise which corresponds to the optimum in terms of the total of clusters detected, and for all beam sizes less than around $2^{\\prime}$ (FWHM) this is equivalent in terms of the number of clusters found in the survey, although the precise optimum noise will be different for different beam sizes.  Accurate total power measurements are only feasible in space or at a very dry site (for example, the South Pole) on the earth's surface, mainly due to the effects of water vapour. In section~\\ref{real} we discuss how our calculations can be adapted to situations closer to those experienced in reality. In particular, we will illustrate how our calculations can be applied to techniques which minimize the effects of the atmospheric contamination which is ubiquitous at the relevant frequencies at most sites. These techniques involve either use of a beam switch or an interferometric array. For the most part our conclusions are similar to those for a total power experiment, although more care is required in interpreting them.  Obviously our goal has some clear dangers since one is, by necessity, required to assume a specific distribution of clusters. In situations where the source population is completely unknown, attempts at optimizing observations in this way are completely useless! However, the distribution and precise characteristics of clusters are becoming much better understood; something which is only likely to improve as observations progress. We have, therefore, used the so-called concordance cosmological parameters~\\cite{OS} and have attempted to model most of the relevant features of the distribution of clusters. In section~\\ref{robust} we test the robustness of our predictions to variations in the cosmological parameters and in the cluster model. Small variations do not affect the optimum significantly but, not surprisingly, large variations substantially change our predictions. We quantify this in terms of the parameters. Of particular concern is the uncertainty in the detailed modelling of the poorly understood  gas physics within the cluster. It is clear that more sophisticated modelling of this physics would be required before we can have full confidence in our predicted optima; they should, nonetheless, provide a useful guide.  For the purposes of this paper, we will ignore the effects of confusion due to radio and infra-red (IR) point sources, the CMB anisotropies and that due the distribution of clusters themselves which are below the flux limit. These will be discussed in a future publication. Suffice to say that the limiting flux densities  and temperatures that we suggest here are compatible with confusion being negligible except for the most extreme cases. Very much deeper surveys run the risk of becoming confusion limited. We note that some of our conclusions concerning the most advantageous instrumental setup do not take into account possible strategies to mitigate problems which might occur due to confusion backgrounds.  We should note that our optimization  are only of relevance to ground-based experiments specifically designed to perform high resolution CMB measurements and SZ surveys at a single frequency. For general purpose, multi-frequency instruments such as PLANCK, the observational strategy (in the case of PLANCK, an all sky survey) has been designed to take into account a wide range of possible applications, not just the ability to detect the SZ in clusters.  Throughout this paper we will work in natural units where $c=\\hbar=k_{\\rm B}=1$. Standard conversion factors can be found, for example, in Kolb \\& Turner (1989). ",
        "conclusions": "\\label{concl}  We have discussed in detail the optimal depth of an SZ cluster survey and have as a by-product gained some insight in the optimal experimental configuration. Ignoring the effects of confusion noise, for every beamsize considered we have shown that there  exists an optimum depth at which one will find the largest  number of clusters for a fixed integration time. The number of clusters found as a function of the flux density or temperature limit is typically quite broad, which is fortunate since the calculations we have made contain many uncertainties. We have shown that the most significant of these are those relating to the $M-T$ relation and the cluster profile. Even within the known uncertainties in these parameters it should be possible to get within $50\\%$ of the optimum.  For an instrumental setup with $\\theta_{\\rm FWHM}^{\\rm min}<2^{\\prime}$ and assuming that beam smoothing is possible, one is brightness temperature limited and it is inefficient to make very deep surveys searching for clusters which match a much smaller beam. One is always in a position to perform shorter integrations on many small beam areas and smooth them down to the lower resolution where the maximum number of clusters are available. These conclusions are particularly sensitive to some of the more sophisticated experimental techniques which are used to counteract the effects of the atmosphere. We have discussed two: in the case of beam switching, we have shown that one is not severely affected if the beam switch is not so small as to prevent one detecting the most abundant clusters which are which have a virial diameter of $\\sim 2^{\\prime}$.  The discussion of the optimal design of an interferometer and the survey strategy is complicated by the extra freedom due to density of visibilities in the $u$-$v$ plane. We have discussed how our methods can be applied in that case, but, other than asserting that the most filled interferometer would have the optimal design, we have not made any statements on the optimal strategy. Given a particular situation it might be possible to do this in more detail.  We hope that our calculations can be used to assist observers in their choice of survey strategy. We acknowledge numerous times that our predictions are model dependent. One could imagine a method by which one could search adaptively for the optimum depth of a survey as part of the process of performing the observation. The basic idea is to take some small patch of sky, compared to the potential total survey volume, and compute the number of sources as a function of $S_{\\rm lim}$. When the logarithmic gradient of this quantity is greater than -2, the optimum has been achieved, and that value of $S_{\\rm lim}$ should be used on the rest of the survey area. Clearly, this will only be possible in scenarios where one is expected to find a large number of objects since one will have to find a sufficiently large number in the initial test area to make the estimate of the gradient accurate enough.  \\subsection*{ACKNOWLEDGEMENTS}  RAB is funded by PPARC and thanks the members of the CMB group at Jodrell Bank Observatory, in particular Peter Wilkinson, Ian Browne and Clive Dickinson, for many useful discussions on the content of the is paper. JW is supported by the DOE and the NASA grant NAG 5-10842 at Fermilab.  \\subsection*{Appendix A}  Let us consider computing the integral ${\\cal I}(B,\\zeta)$ when $\\zeta(\\theta)=F(y)$ where $y=\\theta^2/\\theta_{\\rm vir}^2$ for $\\theta\\le\\theta_{\\rm vir}$ and zero for $\\theta>\\theta_{\\rm vir}$. The function is defined so that  $F(0)=1$. It is easy to see that for the isothermal $\\beta$-model with $\\beta=2/3$ \\be F(y)=\\left(1+\\alpha^2y^2\\right)^{-1/2}{\\tan^{-1}\\left[{\\alpha(1-y^2)^{1/2}\\over 1+\\alpha^2y^2}\\right]\\over \\tan^{-1}\\alpha}\\,, \\ee and it is likely that most sensible profile functions will have this property.  If we define $R=4(\\theta_{\\rm vir}/\\theta_{\\rm FWHM})^2\\log 2$ then \\be {\\cal I}(B,\\zeta)={\\int_0^1\\,e^{-Ry}\\,F(y)\\,dy\\over \\int_0^1\\,F(y)\\,dy}\\,. \\ee One can see that in the limit $R\\rightarrow 0$ the function ${\\cal I}\\rightarrow 1$ which corresponds to the point source limit. The opposite limit, $R\\rightarrow\\infty$, is more interesting since it corresponds to that when the cluster is much larger than the beam. One can show that the leading order term in the asymptotic expansion in $1/R$ is \\be {\\cal I}\\sim \\left[R\\int_0^1\\,F(y)\\,dy\\right]^{-1}\\propto \\theta_{\\rm vir}^{-2}\\,. \\ee From (\\ref{smt}) we see that $S\\propto \\langle T_{\\rm e}\\rangle_{\\rm n}M_{\\rm vir}{\\cal I}d_{\\rm A}^{-2}$ and since $R_{\\rm vir}=\\theta_{\\rm vir}d_{\\rm A}$ we can deduce that \\be S\\propto {\\langle T_{\\rm e}\\rangle_{n}M_{\\rm vir}\\over R_{\\rm vir}^2}\\,, \\ee which is finite as $z\\rightarrow 0$. Hence, we see that the limiting mass tends to a finite value at $z=0$ in contrast to the case of a point source.  \\subsection*{Appendix B}  Let us consider ${\\cal I}(B_{\\rm eff},\\zeta)$ for $B_{\\rm eff}(\\theta)$ given by that for a switch beam system (\\ref{switch}). We will assume that $F(y)$ and $R$ are defined  as in Appendix A, and will define $Q=4(\\theta_{\\rm s}/\\theta_{\\rm FWHM})^2\\log 2$. Using these definitions we find that \\be {\\cal I}(B_{\\rm eff},\\zeta)={\\int_0^1\\,e^{-Ry}\\left[1-e^{-Q}I_0\\left(2\\sqrt{RQy}\\right)\\right]\\,F(y)\\,dy\\over \\int_0^1\\,F(y)\\,dy}\\,. \\ee Understanding the asymptotic behaviour of this integral is more difficult than in the standard case. In order to make some progress let us write $F(y)$ as a power series \\be F(y)=\\sum_{n=0}^{\\infty}A_{m}y^{m}\\,. \\ee where $A_0=1$ and $A_1<0$. From this we can deduce that \\be {\\cal I}(B_{\\rm eff},\\zeta)={\\sum_{m=0}^{\\infty}A_{m}\\left(B_{m}-e^{-Q}C_{m}\\right)\\over \\int_0^1F(y)dy}\\label{ll} \\ee where \\begin{eqnarray} B_m&=&\\int_0^1e^{-Ry}y^mdy\\,, \\\\ C_m&=&\\int_0^1e^{-Ry}I_0\\left(2\\sqrt{RQy}\\right)y^mdy\\,. \\end{eqnarray} One can write $C_{m}$ in terms of $B_{m}$ by writing $I_0(x)$ as a power series in $x$, \\be C_{m}=\\int_{n=0}^{\\infty}{1\\over (n!)^2}(RQ)^{n}B_{m+n}\\,. \\ee Moreover, it can be shown that $B_m\\sim m!/R^{m+1}$ and hence $C_0\\sim e^{Q}/R$ and $C_1\\sim (Q+1)e^{Q}/R^2$. From this we see that in the asymptotic expansion, the first term in (\\ref{ll}) cancels and therefore \\be {\\cal I}(B_{\\rm eff},\\zeta)\\sim {-A_1Q\\over R^2\\int_0^1F(y)dy}\\,. \\ee This cancellation is enough to ensure that ${\\cal I}\\rightarrow0$ as $z\\rightarrow 0$  sufficiently fast so that \\be S\\propto {\\langle T_{\\rm e}\\rangle_nM_{\\rm vir}d_{\\rm A}^2\\over R_{\\rm vir}^4}\\,, \\ee in this limit. It is easy to see that this forces the limiting mass to tend to infinity as $z\\rightarrow 0$."
    },
    "0410/astro-ph0410386_arXiv.txt": {
        "abstract": "{  We study the two main corrections generally applied to narrow-band \\Ha\\ fluxes from galaxies in order to convert them to star formation rates, namely for \\NII\\ contamination and for extinction internal to the galaxy.  From an imaging study using carefully chosen narrow-band filters, we find the \\NII\\ and \\Ha\\ emission to be differently distributed. Nuclear measurements are likely to overestimate the contribution of \\NII\\ to total narrow-band fluxes. We find that in most star formation regions in galaxy disks the \\NII\\ fraction is small or negligible, whereas some galaxies display a diffuse central component which can be dominated by \\NII\\ emission. We compare these results with related studies in the literature, and consider astrophysical explanations for variations in the \\NII /\\Ha\\ ratio, including metallicity variations and different excitation mechanisms.  We proceed to estimate the extinction towards star formation regions in spiral galaxies, firstly using Br$\\gamma$/\\Ha\\ line ratios.  We find that extinction values are larger in galaxy nuclei than in disks, that disk extinction values are similar to those derived from optical emission-line studies in the literature, and that there is no evidence for heavily dust-embedded regions emerging in the near-IR, which would be invisible at \\Ha.  The numbers of galaxies and individual regions detected in Br$\\gamma$ are small, however, and we thus exploit optical emission line data from the literature to derive global \\Ha\\ extinction values as a function of galaxy type and inclination. In this part of our study we find only a moderate dependence on inclination, consistent with broad-band photometric studies, and a large scatter from galaxy to galaxy.  Typical extinctions are smaller for late-type dwarfs than for spiral types.  Finally, we show that the application of the type-dependent extinction corrections derived here significantly improves the agreement between star formation rates calculated using \\Ha\\ fluxes and those from far-infrared fluxes as measured by {\\it IRAS}. This again supports the idea that heavily dust-embedded star formation, which would be underestimated using the \\Ha\\ technique, is not a dominant contributor to the total star formation rate of most galaxies in the local Universe. ",
        "introduction": "\\label{sec:intro}  The \\Ha\\ luminosity observed in spiral and irregular galaxies is believed to be a direct tracer of the ionisation of the inter-stellar medium (ISM) by the ultraviolet (UV) radiation which is produced by young high-mass OB stars.  Since only high-mass ($>$10~\\Msolar) and, therefore, short-lived ($<$20~Myr) stars contribute significantly to the integrated ionising flux, the \\Ha\\ emission line thus provides a nearly instantaneous measure of the star formation rate (SFR), independent of the previous star formation history.  The two main advantages of using \\Ha\\ to detect star formation are the direct relationship between the nebular line emission and the massive SFR, and also the high sensitivity.  A small telescope can map star formation down to low levels at high angular resolution, even in faint, low surface-brightness galaxies.  The \\Ha\\ Galaxy Survey (\\Ha GS) is a study of the star formation properties of a representative sample of galaxies in the local Universe using this technique.  We have imaged 334 nearby galaxies in both the \\Ha\\ line and the $R$-band continuum.  The sample consists of all Hubble types from S0/a to Im with recession velocities between 0 and 3000~km~s$^{-1}$.  All galaxies were observed with the, now unfortunately decommissioned, 1.0 metre Jacobus Kapteyn Telescope (JKT), part of the Isaac Newton Group of Telescopes (ING) situated on La Palma in the Canary Islands.  The selection and the observation of the sample are discussed in \\citet{paper1}, hereafter Paper I.  A key element of this project is to extend the study to fainter galaxies than has been done in previous large surveys.  The two major limitations of narrow-band \\Ha\\ imaging are contamination by the \\NII\\ line doublet, and uncertainties in the extinction corrections to be applied to each galaxy.  In this paper we will investigate refinements to the corrections applied by previous authors.  The \\Ha GS sample contains more faint, dwarf galaxies than most previous studies, and thus we aim to derive corrections applicable to both these objects, and the brighter spiral galaxies.  The \\NII\\ lines are located either side of the 6563\\AA\\ \\Ha\\ line, at 6548\\AA\\ and 6583\\AA\\, and will be referred to as \\NII-6548 and \\NII-6583 in this paper.  The widths of the narrow-band \\Ha\\ filters range from 44\\AA\\ to 55\\AA\\, thus including both \\NII\\ lines - as do almost all narrow-band \\Ha\\ observations of galaxies.  The \\NII-6548 line is weaker than the \\NII-6583 line, with \\NII-6583/\\NII-6548 $\\sim$ 3.  The contribution to the total flux for both lines together (\\NII-total henceforth) causes a scatter when using \\Ha\\ measurements to infer SFRs.  In order to improve calculations of SFRs it is therefore important to be able to correct for the \\NII\\ contamination.  Whilst the ratio of \\NII\\ to \\Ha\\ in the nuclei of spiral galaxies has been well studied (e.g., \\citealt{keel83,rubin86}) and found to be high (often greater than 1), \\citet{kk83} find that nuclear \\NII\\ emission rarely dominates the integrated photometry of galaxies and hence global corrections for \\NII\\ contamination are much smaller.  The most commonly used corrections for entire galaxies are those derived by \\citet{ken83} and \\citet{kk83}. Spectrophotometric \\NII/\\Ha\\ ratios of individual extragalactic \\HII\\ regions from the literature (see \\citealt{kk83} and references therein) were compiled from 14 spiral galaxies (mostly of type Sc) and 7 irregular galaxies.  The average \\Ha/(\\Ha\\ + \\NII-total) ratio was found to be fairly constant, spanning the ranges 0.75 $\\pm$ 0.12 for the spirals, and 0.93 $\\pm$ 0.05 for the irregulars (\\citealt{ken83}).  In terms of the ratio \\NII-total/\\Ha\\, this corresponds to a median value of 0.33 for spirals and 0.08 for irregulars. These values were calculated by finding the \\NII-total/\\Ha\\ ratio of the brightest \\HII\\ regions, averaging for each galaxy and then determining the mean value for spiral and irregular types.  This implicitly assumes that all \\HII\\ regions have the same proportion of \\NII-total to \\Ha\\ emission as those regions measured.  However, the integrated galaxy spectroscopy presented by \\citet{ken92} shows generally stronger \\NII\\ emission than the earlier studies.  Mean values of \\NII-6583/\\Ha\\ are 0.75 (st. dev 0.55) for types Sa--Sbc, 0.36(0.16) for types Sc--Sm, and 0.15(0.11) for Im galaxies (omitting the, probably misclassified, galaxy IC~883 from the Im sample).  Including the \\NII-6548 emission would make the discrepancy larger; however, these ratios were taken from deblended spectra as the \\Ha\\ and \\NII\\ lines were only marginally resolved.  Similarly, \\citet{mcq95} find high mean \\NII-total/\\Ha\\ ratios, of 0.606$\\pm$0.432 for 3 normal spirals, 0.909$\\pm$1.34 for 12 barred spirals and 0.838$\\pm$0.431 for 5 Type 2 Seyfert galaxies.  The 3 blue compact dwarf galaxies in their sample have a much lower mean \\NII-total/\\Ha\\ ratio, 0.104$\\pm$0.018, and 3 blue compact galaxies have an intermediate mean ratio, 0.276$\\pm$0.036.  The one LINER in their sample has very strong \\NII\\ emission, with \\NII-total/\\Ha\\ $=$ 3.627.  These ratios are derived from circular apertures of 13.5~arcsec diameter, thus sampling only the central regions of most of these galaxies.  \\citet{jan00} use integrated spectra of 196 nearby galaxies to show that the principal determinant of \\NII/\\Ha\\ ratio is galaxy luminosity rather than type. In their study, the \\NII/\\Ha\\ ratios for galaxies brighter than $M_B = $ -19.5 are in agreement with the values found by \\citet{ken92}, but fainter than this a striking trend is seen towards much lower values of this ratio.  These fainter galaxies are dominated in the \\citet{jan00} sample by late-type Sd--Im galaxies, but interestingly the same trend is clearly seen for earlier type spirals.  At $M_B = $ -16, \\NII/\\Ha\\ ratios of $\\sim$0.1 are typical. A similar trend in log(\\NII-total/\\Ha\\ ) vs galaxy luminosity is found by \\citet{gav04}.  In this paper, a new method of investigating the relative strengths of \\Ha\\ and \\NII\\ is introduced and explored.  This method will enable both the integrated ratio over the whole galaxy and the ratios of individual regions to be calculated.  Particular attention will be paid to radial variations in the \\NII-total/\\Ha\\ ratio, both for the practical reason that many literature determinations of this ratio have been based on measurements of the central regions of galaxies, and because there are theoretical reasons to expect significant variation with radius, due to metallicity changes and different excitation processes.  The major source of systematic error in the conversion of \\Ha\\ fluxes to SFRs is due to the effects of extinction within the galaxy being observed.  There is a great deal of uncertainty as to the exact magnitude of this extinction and how it varies with galaxy luminosity and type.  Studies using radio data (which are not affected by extinction) of large samples of individual \\HII\\ regions in nearby galaxies yield mean extinction values ranging from {\\it A}(\\Ha) = 0.5 mag to {\\it A}(\\Ha) = 1.8 mag (e.g. \\citealt{caplan86,kaufman87,hulst,caplan96,niklas}). \\citet{ken98} adopts a single correction of 1.1 mag for the effective \\Ha\\ extinction of entire galaxies for all Hubble types.  These samples, however, mainly contain bright galaxies, whereas the \\Ha GS also includes many faint and irregular galaxies. For very small irregular galaxies, extinction may be systematically lower due to the smaller quantity of ISM which the \\Ha\\ flux must traverse. Alternatively, most of the extinction could be associated with the star forming region itself, in which case galaxy type and size will be of lesser importance.  In this paper we will determine the extinction towards known star formation regions, and search for highly embedded star formation which may not be apparent through optical observations, using 3 techniques. The first will present new \\Brg\\ imaging data for galaxies in our sample, and use the \\Brg /\\Ha\\ ratio to estimate extinction towards individual \\HII\\ regions in these galaxies. The number of galaxies detected in the \\Brg\\ line is small, however, so a further investigation is presented making use of literature \\Ha\\ and \\Hb\\ line fluxes for a large sample of spiral galaxies.  This sample enables type-dependent extinction corrections to be derived and the inclination dependence of extinction to be investigated.  Finally, a statistical analysis is made of star formation rates for galaxies in the present sample, using both optical \\Ha\\ fluxes and far-IR fluxes from the IRAS survey.  This provides a test of the type-dependent extinction corrections derived here, and enables a further search for heavily dust-embedded star formation which may be missed by optical or even near-IR observations.  The organisation of the rest of this paper is as follows. In Section \\ref{sec:nii} of this paper we will investigate the fraction of \\NII\\ emission in the total \\Ha\\ + \\NII\\ flux, how this varies between bulges and disks of spiral galaxies, and suggest possible reasons for this variation. In Section \\ref{sec:brg}, we use \\Brg\\ flux measurements from a sample of galaxies to investigate internal extinction.  In Section \\ref{sec:ucm}, we use data published by the Universidad Complutense de Madrid (UCM) survey collaboration to investigate the dependences of extinction on galaxy inclination and morphology. In Section \\ref{sec:sfrcomp}, the SFRs calculated using both \\Ha\\ and far-infrared (FIR) data are compared, used to check the type-dependent extinction corrections, and to constrain deeply-embedded star formation.  Section \\ref{sec:conc} summarizes the main conclusions of this paper. ",
        "conclusions": "\\label{sec:conc} \\renewcommand{\\labelitemi}{$\\bullet$}  We have studied the two major corrections which need to be applied to narrow-band \\Ha\\ fluxes from galaxies in order to convert them to SFRs, i.e. \\NII\\ contamination removal and correction for extinction internal to the galaxy in question. \\begin{itemize} \\item From an imaging study using carefully-chosen narrow-band filters, we find that the \\NII\\ emission is generally very differently distributed to the \\Ha\\ emission, and that in particular nuclear measurements (e.g. from slit spectroscopy) can significantly over-estimate the contribution of \\NII\\ to total narrow-band fluxes.  In most star formation regions in galaxy disks, the \\NII\\ fraction is small or negligible, and \\NII\\ corrections applied in most previous studies may significantly under-estimate disk star formation rates as a result. \\item We estimate the extinction towards star formation regions in spiral galaxies from Br$\\gamma$/\\Ha\\ line ratios.  The main results from this study are that extinctions are larger for regions in galaxy nuclei compared with those in disks; disk extinction values are similar to those derived from optical emission-line ratios; and there is no evidence for heavily dust-embedded regions emerging in the near-IR, which would be invisible at \\Ha.  However, the numbers of galaxies and individual regions detected using this method are small, and we thus exploit optical emission-line data for the UCM study to derive global \\Ha\\ extinction values as a function of galaxy type and inclination. \\item We find evidence supporting the overall size of corrections ($\\sim$1~mag.) applied by previous authors, and determine typical extinctions to be smaller for late-type dwarfs than for spiral types. \\item We show that the application of type-dependent extinction corrections derived here significantly improves the agreement between star formation rates calculated using \\Ha\\ fluxes and those from far-IR fluxes as measured by the IRAS satellite. \\item These findings support the idea that heavily dust-embedded star formation, which would be underestimated using the \\Ha\\ technique, is not a dominant contributor to the total star formation rate of most galaxies. \\end{itemize}"
    },
    "0410/astro-ph0410453_arXiv.txt": {
        "abstract": "{\\src\\ is thought to be a 321s orbital period (the shortest known) double white dwarf binary system. According to the double degenerate binary (DDB) scenario this source is expected to be one of the strongest Gravitational Wave (GW) emitter candidates.  In the last years \\src\\ has been studied in great details, through multiwavelength observational campaigns and from the point of view of data analysis result interpretations. We present here the timing results obtained thanks to a 3.5-year long optical monitoring campaign carried out by the Very Large Telescope (VLT) and the Telescopio Nazionale Galileo (TNG) which allowed us to detect and study the orbital period derivative (spin--up at a rate of $\\sim$10$^{-3}$s yr$^{-1}$) of the 321s modulation, to detect the linear polarisation (at a level of about 2\\%), and to study the broad band energy spectrum. The VLT/TNG observational strategy we used allowed us to rely upon a P-\\.P coherent solution which we finally extended backward to the 1994 ROSAT observation of \\src. ",
        "introduction": "\\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{PuzzR_VLT.eps} \\includegraphics{PuzzX_XMM.eps}} \\caption{Left panel: March 2003 VLT FORS1 R-band image of the sky region including the position of the optical counterpart to \\src\\ (marked with the circle; effective exposure time of 40 minutes). Right panel: 0.2-10.0\\,keV EPIC-pn \\xmm\\ image of \\src\\ (brightest object; exposure time of $\\sim$26000\\,s). } \\label{images} \\end{figure*}  \\src\\ was discovered in 1990 with the \\R\\ satellite during the All-Sky Survey \\citep[RASS;][]{beu99}. However, it was only in 1999 that a periodic signal at 321\\,s was detected in its soft X-ray flux with the \\R\\ HRI \\citep[][hereafter I99]{io99} \\defcitealias{io99}{I99}. The detection of X-ray pulsations was also reported independently by \\citet{bur01}. Subsequent deeper optical studies allowed to unambiguously identify the optical counterpart of \\src, a blue $V=21.1$ ($B=20.7$) star \\citep[][hereafter I02B]{io02a,io02b}\\defcitealias{io02b}{I02B}. $B$, $V$ and $R$ time-resolved photometry revealed the presence of a $\\sim 15$\\% modulation at the $\\sim 321$\\,s X-ray period \\citep{io02b,ram02a}. \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{phase_noppdot.eps} \\includegraphics{xo_jun04.eps}} \\caption{Left panel: Results of the phase fitting technique used to infer the P-\\.P coherent solution for \\src: the linear term (P component) has been corrected, while the quadratic term (the \\.P component) has been kept for clarity.  The best \\.P solution inferred for the optical band is marked by the solid fit line.  Right panel: 1994--2002 phase coherently connected X--ray folded light curves (largest pulsed fraction) of \\src, together with the 2001-2004 phase connected folded optical light curves. Two orbital cycles are reported for clarity. } \\label{timing} \\end{figure*} The VLT spectral study revealed a blue continuum with no intrinsic absorption lines \\citepalias{io02b}. Broad ($\\rm FWHM\\sim 1500~\\rm km~s^{-1}$), low equivalent width ($EW\\sim -2\\div-6$ \\AA) emission lines from the He~II Pickering series (plus additional emission lines likely associated with He~I, C~III, N~III, etc.) were instead detected. These findings, together with the period stability and absence of any additional modulation in the 1\\,min--5\\,hr period range, are interpreted in terms of a double degenerate He-rich binary \\citep[a subset of the AM CVn class; see][for a review]{warn95} with an orbital period of 321\\,s, the shortest ever recorded. Moreover, \\src\\ was noticed to have optical/X-ray properties similar to those of \\srcm, a 569\\,s modulated soft X-ray source proposed as a double degenerate system \\citep{ram00,ram02b}.  More recently, \\citet{hak03,hak04} reported on the detection of spin--up, at a rate of $\\sim$6.2$\\times$10$^{-11}$\\, s~s$^{-1}$, for the 321\\,s orbital modulation, based on optical data taken form the Nordic Optical Telescope (NOT) and the VLT archive, and by using incoherent timing techniques. Similar results were reported also by \\citet{stro03} for the X-ray data (ROSAT and Chandra) of \\src\\ spanning over 10 years and based on the \\citet{hak03} results. Finally, \\citet{io03b} carried out a nearly-simultaneous Chandra/VLT observational campaign reporting the presence of nearly anti correlation between the optical and X--ray modulation. The X-ray spectral study also allowed to characterise the emission mechanism; a black body with a temperature of $\\sim$60\\,eV \\citep{io03b}.  The study of \\src\\ and \\srcm\\ has posed, in the last years, serious questions about their possible origin \\citep[see][for a review of the proposed theoretical models]{crop03}. Moreover, there is a number of additional nice reasons for studying \\src\\ and the related objects. Among other are the study of the gravitational wave, and the possibility that \\src\\ is a progenitor of the He-accreting DDBs also named AM CVns.  For this reason we requested and performed in the last 3.5 years several observations ranging from IR to X-ray band in order to unveil the nature of \\src\\ and, correspondingly, of \\srcm. Here we report on the results obtained from (a) the VLT/TNG time-resolved photometric monitoring aimed at accurately measure the P and \\.P of \\src\\ based on coherent timing techniques, (b) VLT-FORS1 polarimetry observations, and (c) an \\xmm\\ pointing.   \\section {Optical Observations}  After the successful first optical time-resolved observations of \\src\\ during 2001 (January 1st at TNG, and November 11th), we started a relatively long-term project aimed at monitoring the source 321\\,s modulation. We obtained 21 pointings between November 2002 and May 2004 (11 at VLT and 10 at TNG) scheduled in a way such that it was possible to keep the phase coherence among observations (the first observations were obtained at 1st-2nd nights, 9-10th nights, 19-20th nights, 49-50th nights, etc.). Four different optical bands (B, V, R and I) have been used in order to study and monitor the pulse shape and pulsed fraction as a function of wavelength and time (left panel of figure~\\ref{images} shows the R-band image of the region around \\src).  Such a strategy resulted quite efficient in reaching the purposes of the timing analysis, and allowed us also to extend the coherent solution backward to the 2001 optical observations. The best optical solution we found for P-\\.P is for P=321.53040(4)\\,s, \\.P=-3.6(1)$\\times$10$^{-11}$\\,s~s$^{-1}$ \\citep[90\\% uncertainties are reported; for more details see][see also figure~\\ref{timing}, left panel]{io04}. Moreover, we found a slightly energy--dependent pulse shape with the pulsed fraction increasing toward longer wavelengths, from $\\sim$12\\% in the B-band to nearly 14\\% in the I-band \\citep{io04}.  \\begin{figure}[bt] \\centering \\resizebox{\\hsize}{!}{\\rotatebox[]{-90}{\\includegraphics{QU.eps}}} \\caption{ Q and U Stokes parameters for all the objects detected within 3\\arcmin\\ from the \\src\\ position (on-axis). Values inferred for \\src\\ are clearly marked (filled circle).} \\label{QU}% \\end{figure} \\begin{figure}[th] \\centering \\resizebox{\\hsize}{!}{\\rotatebox[]{-90}{\\includegraphics{fitxmm.eps}}} \\caption{The results of the \\xmm\\ phase-resolved spectroscopy (PRS) analysis for the absorbed blackbody spectral parameters: absorption, blackbody temperature, blackbody radius (assuming a distance of 500\\,pc), and absorbed (triangles) and unabsorbed (asterisks) flux. Superposed is the folded X-ray light curve.} \\label{xmm}% \\end{figure}  The relatively high accuracy obtained for the optical phase coherent P-\\.P solution (in the January 2001 - May 2004 interval) allowed us to extrapolate it backward to the ROSAT observations without loosing the phase coherency, i.e.  only one possible period cycle consistent with our P-\\.P solution. Given the wider time interval spanned by the X--ray observations (9.5 years) an even more accurate solution was possible.  After taking into account the ROSAT photon arrival time spacecraft clock time - UTC correction (in order to compare the ROSAT data with the Chandra and \\xmm\\ observations), our best (optically selected) X-ray solution is P=321.53033(2)\\,s, \\.P=-3.67(1)$\\times$10$^{-11}$\\,s~s$^{-1}$ \\citep[for more details see][]{io04}. Figure~\\ref{timing} (right panel) shows the optical (2001-2004) and X--ray (1994-2002) light curves folded by using the above reported P-\\.P coherent solution, confirming the amazing stability of the X--ray/optical anti correlation first noted by \\citet{io03b}.   In February 2003 we also collected VLT FORS1 polarimetry data for the optical counterpart to \\src\\ in the B-band. Due to unforeseen large FORS1 overheads only phase-resolved linear polarisation observations were obtained.  In figure~\\ref{QU} we reported the U and Q Stokes parameters for all the detected objects within 3\\arcmin\\ from the \\src\\ optical counterpart position (objects at larger off--axis angles are indeed affected by the presence of spurious polarisation: more details on this problem can be found at www.eso.org/instruments/fors/pola.html). \\src\\ was found to be linearly polarised at a level of $\\sim$2.0$\\pm$0.3\\% (after correcting the the field average polarisation, $\\sim$0.7\\%). It is worth noting that the marginal detection of circular polarisation, at level of about 0.5\\%, has been recently reported by \\citet{rei04}, and ascribed to a $\\sim$10$^{6}$\\,Gauss magnetic field. ",
        "conclusions": "In this contribution we briefly listed the main results obtained thanks to a four years optical and X-ray monitoring of \\src, a DDB with an orbital period of only 5.4min. Even though the optical monitoring is still active (in fact is quite important to continue the \\.P study to look for variations or, more interestingly, a d\\.P/dt component) a number of important implications can be considered.  It is now assessed that the orbital period is decaying  at a rate that is nearly consistent with that predicted by the ultracompact binary model in which the loss or orbital angular momentum is dominated by gravitational radiation and in which there is no mass exchange between the two stars \\citep[see also][]{stro03}. In this respect the unipolar inductor model proposed by \\citet{wu02} is a possible scenario to account for the observed X--ray flux from \\src.  The X--ray and optical emissions are anti correlated at a level of stability (no variations on a baseline of at least 4 years) which is only observed in phase-locked binary systems. The anti correlation is by far and so far the strongest (indirect) suggestion of the binary nature of the 321\\,s, where the X--rays illuminate the companion surface originating the phase--shifted optical modulation.  The detection of linear and circular polarisation might imply the presence of a relatively faint magnetic field, however we note that the low polarisation level is also consistent with other possibilities, such as the Thomson scattering (sometimes observed in eclipsing binaries)."
    },
    "0410/astro-ph0410179_arXiv.txt": {
        "abstract": "\\noindent The determination of the primary cosmic ray mass composition from the longitudinal development of atmospheric cascades is still an open problem. In this work we propose a new method of the multiparametric topological analysis and show that if both $X_{max}$ - the depth of shower maximum and $N_{max}$ - the number of charged particles in the shower maximum are used, reliable results can be obtained. ",
        "introduction": "The study of the longitudinal profile of individual atmospheric cascades started in the early eighties with the development of the fluorescent light detection technique implemented for the first time within the framework of the Fly's Eye experiment \\cite{Baltr}. After these pioneering efforts, there has been only one experimental array continuing this type of studies: HiRes \\cite{Matth} and only recently a new and much more powerful detector has started to collect data: the Fluorescent Detector (FD) of the Pierre Auger Observatory \\cite{Bluem}. This instrument will produce a large data flow over the next decades and is therefore calling for new and accurate data analysis procedures capable to fully exploit the large amount of information contained in the FD data.  It is rather surprising, that while there are many methods, both parametric and non-parametric (KNN, Bayesian methods, pattern recognition, neural nets etc.), used to discriminate individual cascades on the basis of ground--based information, very little has been done to fully exploit the amount of information contained in FD data. To our knowledge, the only method available so far makes use of $X_{max}$ - the depth of the maximum cascade development \\cite{Gaiss} and derives the observed {\\em mean} mass composition as a function of the primary energy. Since there is a minimum bias in the detection of cascades of different origin, the observed mass composition coincides practically with the primary composition. It has to be stressed also that this approach relies on statistical grounds and therefore does not allow the identification of the primary particle for each individual cascade. Furthermore, even though in the longitudinal profile the $X_{max}$ parameter is the most sensitive to the mass of the primary particle, its sensitivity is still weak. For instance, at a primary energy of 1 EeV (10$^{18}$eV) the mean iron induced cascade has $X_{max}$ only (11-12)\\% lower than for a proton induced one, i.e. a difference which is of the same order of magnitude as the intrinsic fluctuations of $X_{max}$.  As we shall discuss below, such unsatisfactory situation improves drastically if other, seemingly less significant parameters, are taken into account. Among them, we might have: $N_{max}$ - the number of particles (mostly electrons) in the maximum of the cascade, the speed of rise in the particle number etc. For instance, at fixed primary energy, $N_{max}$ is about the same for all cascades, and even though iron induced cascades produce more muons and less energy is carried out by the electrons, the effects on $N_{max}$ are very small. However, due to the lower energy per constituent nucleon in the primary iron nucleus, the cascade development and the rise of the cascade curve are on average faster than for cascades originating from protons. This useful information is neglected when only $N_{max}$ is taken into account. This type of arguments triggered our efforts to find methods capable to fully exploit the information contained in the cascade curves and to allow the identification (at least in terms of probabilities) of the cascade origin also for individual showers.  Although tailored for possible applications in the context of the Pierre Auger experiment, methods described below are quite general and may find application in other similar experiments. However, in what follows we shall focus on data similar to those expected from the Fluorescent Detector (FD) of the Pierre Auger Observatory: namely on the longitudinal profile of each atmospheric cascade, i.e. on the number of charged particle $N_{ch}$ as a function of the atmospheric depth $X$. As it will be shown below, this profile carries more information than $X_{max}$ or $N_{max}$ alone.  We would like to stress that even though the fraction of the hybrid events, in which the information on the shower from the Pierre Auger Surface Array is supplemented by the Fluorescent Detector data, will hardly exceed 10\\% of the total statistics accumulated by the Surface Array, these events need to be properly handled since they contain the maximum information. In this paper we restrict ourselves to the analysis of the longitudinal development of cascades. The treatment of hybrid data will be addressed in subsequent papers. ",
        "conclusions": "We proposed and tested the Multiparametric Topological Analysis method for the determination of the {\\em mean} primary cosmic ray mass composition on the basis of measurements of the longitudinal development of atmospheric cascades. The method employs more information about the longitudinal development of atmospheric cascades than just the depth of the cascade maximum. Definitely the method needs further developments before being used for processing real experimental cascades, but the first results are very encouraging. The next efforts should include the increase of simulation statistics, the study of the effect of the shower energy spectrum and angular distribution, the accurate estimate of the reconstruction errors. The great advantage of the proposed MTA method is that it is extremely easy to use and generalize for a larger number of observables.  \\vspace{5mm}  \\noindent{\\bf Acknowledgements} \\medskip  \\noindent Authors thank M.Risse for simulations of cascades used in this work and K.-H.Kampert and L.Perrone for useful discussions. \\medskip"
    },
    "0410/astro-ph0410723_arXiv.txt": {
        "abstract": " ",
        "introduction": "The space weather effects of large solar flares motivate flare prediction. For example, the soft X-ray flux due to large flares causes increased ionisation of the upper atmosphere, which can interfere with high frequency radio communication. There is considerable interest in knowing when such short-wave fadeouts are likely to occur, and Australia's Ionospheric Prediction Service (IPS) issues predictions on this basis.\\footnote{see http://www.ips.gov.au} Other agencies issuing flare predictions include the US National Oceanic and Atmospheric Administration (NOAA)\\footnote{see http://www.sec.noaa.gov/ftpdir/latest/daypre.txt} and NASA\\footnote{see http://beauty.nascom.nasa.gov/arm/latest/}.  Existing methods of prediction are probabilistic, and rely e.g.\\ on the classification of physical characteristics of active regions, and historical rates of flaring for regions with a given classification (McIntosh 1990; Bornmann \\& Shaw 1994). One weakness of classification based approaches is that regions with a given classification may exhibit a wide range of flaring rates. The method of McIntosh (1990) also considers other information including the number of large flares already produced by an active region (the tendency of an active region which has produced large events to subsequently produce large events is called persistence), but this is done in an ad hoc way. No consideration is given to the important information in the number of small events already observed.  A new approach to flare prediction (Wheatland 2004a) exploits the history of observed flaring together with simple phenomenological rules of flare statistics to make a prediction, or to refine an existing prediction. The basic method is as follows. It is well known that the size distribution of flares (e.g.\\ the distribution of peak soft X-ray flux) follows a power law (e.g.\\ Crosby, Aschwanden \\& Dennis 1993): \\begin{equation}\\label{eq:fdist} N(S)=\\lambda_1(\\gamma-1)S_1^{\\gamma-1}S^{-\\gamma}, \\end{equation} where $N(S)$ is the number of events per unit size $S$ and per unit time, $\\lambda_1$ is the total rate of events above size $S_1$, and $\\gamma$ is the power-law index. Suppose we are interested in the probability of a large event ($S\\geq S_2$) occurring in a time $\\Delta T$. The expected rate of events above $S_2$ is, according to Equation~(\\ref{eq:fdist}) \\begin{equation}\\label{eq:rate_big} \\lambda_2=\\lambda_1 \\left( \\frac{S_1}{S_2}\\right)^{\\gamma-1}. \\end{equation} Flare occurrence may be described on short timescales as a Poisson process in time (e.g.\\ Moon et al.\\ 2001), and on longer timescales as a time-dependent Poisson process (e.g.\\ Wheatland 2001). According to Poisson statistics, the probability of at least one large event in time $\\Delta T$ is \\begin{equation}\\label{eq:prob_big} \\epsilon = 1-\\exp(-\\lambda_2 \\Delta T). \\end{equation} To apply these formulae it is necessary to estimate $\\lambda_1$ and $\\gamma$ from data, and hence to estimate $\\epsilon$. We adopt a Bayesian approach, in which `estimating' a parameter means calculating a posterior probability distribution for the parameter, given available data and any prior information (e.g.\\ Jaynes 2003). We assume that a a sequence of $M$ events with sizes $s_1,s_2,...,s_M$ (all larger than $S_1$) have been observed to occur at times $t_1< t_2< ...< t_M$ respectively. The power-law index $\\gamma$ may be approximated by the maximum likelihood value (Bai 1993) \\begin{equation}\\label{eq:gam_ML} \\gamma^{\\ast}=\\frac{M}{\\ln \\pi}+1, \\quad {\\rm where} \\quad \\pi=\\prod_{i=1}^M\\frac{s_i}{S_1}. \\end{equation} To estimate the rate we adopt a piecewise constant Poisson model. Hence we need to identify the most recent interval $T^{\\prime}$ during which the rate is constant, and we assume that $M^{\\prime}\\leq M$ events occurred during that time. One approach to the determination of the interval $T^{\\prime}$ is to use the `Bayesian blocks' procedure (Scargle 1998), which is discussed in more detail below. Based on $M^{\\prime}$, $T^{\\prime}$ and $\\gamma^{\\ast}$, the posterior probability distribution for $\\epsilon$ is (Wheatland 2004a) \\begin{eqnarray}\\label{eq:pbig_simp} P(\\epsilon ) &=& C\\left[-\\ln (1-\\epsilon ) \\right]^{M^{\\prime}} (1-\\epsilon )^{\\left( T^{\\prime}/\\Delta T\\right) \\left(S_2/S_1\\right)^{\\gamma^{\\ast}-1}-1} \\nonumber \\\\ &\\times & \\Lambda \\left[-\\frac{\\ln (1-\\epsilon )}{\\Delta T} \\left(\\frac{S_2}{S_1}\\right)^{\\gamma^{\\ast}-1} \\right], \\end{eqnarray} where $\\Lambda (\\lambda_1)$ is the prior distribution for $\\lambda_1$, i.e.\\ the distribution we would assign to $\\lambda_1$ in the absence of any data, and $C$ is the normalisation constant, determined by the requirement $\\int_{0}^{1}P(\\epsilon)d\\epsilon=1$. The mean of $P(\\epsilon )$ provides the estimate of the probability of at least one large event within time $\\Delta T$, and the standard deviation of the distribution provides an estimate of the associated uncertainty. ",
        "conclusions": "A Bayesian approach to flare prediction (Wheatland 2004a) has been tested for whole-Sun prediction of GOES soft X-ray events, based on the NOAA Solar Event Lists for 1975-2003. The method is found to over-predict events of M class and above, but performs quite well for prediction of X class events.  There are several possible reasons for the over-prediction of events above M class. One possibility is that the method is systematically late in detecting the decline in rate associated with the decay of a large active region, or the rotation of a large active region off the disk. The method uses the Bayesian blocks procedure to detect rate changes, and is always trying to `catch up' with what the Sun is doing. A specific limitation of the Bayesian blocks method is that it must have at least one event in a block, so that the rate is never identically zero. This limitation leads to overestimation of the rate at times of very low activity. It should also be noted that the existing Bayesian blocks procedure is not guaranteed to find an optimal decomposition --- in future this procedure will be replaced by an optimal algorithm recently devised by Scargle (Scargle, private communication, 2003). Another possibility is that there is a bias in each individual prediction which becomes apparent in the analysis of a large number of predictions. Such a bias will become less serious if each individual prediction is more accurate. The fractional error in each prediction goes as $(M^{\\prime})^{-1/2}$ (Wheatland 2004a), where $M^{\\prime}$ is the number of events associated with the estimation of the rate above size $S_1$. This error becomes smaller with increasing $M^{\\prime}$, i.e.\\ if a smaller $S_1$ can be chosen. It should be noted that the GOES event lists are a relatively poor choice for the present purpose because the time varying soft X-ray background means that a relatively large $S_1$ must be used (see Figure~1).  The GOES event lists also have other shortcomings as a basis for prediction from event statistics. Besides the departure from a power law at small sizes, it is likely the lists are incomplete above the nominal threshold $S_1$, e.g.\\ due to the difficulty of distinguishing two flares occurring close together in time (Wheatland 2001). This is unlikely to be important for the test described here, provided that the size distribution obeys a power law above the threshold, since both prediction and validation of the prediction rely on the same (possibly incomplete) lists. Another problem with the GOES event lists is that the GOES peak fluxes are not background subtracted, so that the intrinsic flux due to an M class event at a time of low activity, when the background is low, is larger than the intrinsic flux due to an M class event at a time of high activity, when the background is high.  However, again this is not particularly important to the present method provided that the size distribution for observed events is not distorted from a power-law form by this effect. The use of a previous year of data to make a prediction allows the possibility of incorporating variation in the power-law index with the solar cycle (e.g.\\ due to this effect) but in fact we find no evidence of such a variation. In subsequent work the method will be applied to more accurate event catalogs.  The present test is limited to whole-Sun prediction. In the future the method will also be applied to individual active regions, and other prior information on the rate (e.g.\\ the McIntosh classification of the associated sunspots) will be incorporated. However we note that, even in this simple form, the method has out-performed NOAA predictions for X class events.  Finally we note that automated predictions, based on this method, are now published on the web.\\footnote{see http://www.physics.usyd.edu.au/$\\sim$wheat/prediction/prediction.html} Predictions are made each day using the latest NOAA solar event lists. The web pages include a running score of how reliable the published predictions are, in the form of automatically updated reliability plots."
    },
    "0410/astro-ph0410209_arXiv.txt": {
        "abstract": "The aim of this paper is to set constraints of the epochs of early-type galaxy formation through the 'archaeology' of the stellar populations in local galaxies.  Using our models of absorption line indices that account for variable abundance ratios, we derive ages, total metallicities, and element ratios of 124 early-type galaxies in high and low density environments.  The data are analyzed by comparison with mock galaxy samples created through Monte Carlo simulations taking the typical average observational errors into account, in order to eliminate artifacts caused by correlated errors. We find that all three parameters age, metallicity, and \\aFe\\ ratio are correlated with velocity dispersion. We show that these results are robust against recent revisions of the local abundance pattern at high metallicities. To recover the observed scatter we need to assume an intrinsic scatter of about 20 per cent in age, $0.08\\;$dex in [\\ZH] and $0.05\\;$dex in [\\aFe].  All low-mass objects with $M_*\\la 10^{10}\\;\\Msun$ ($\\sigma\\la 130\\;$km/s) show evidence for the presence of intermediate-age stellar populations with low \\aFe\\ ratios.  About 20 per cent of the intermediate-mass objects with $10^{10}\\la M_*/\\Msun\\la 10^{11}$ ($110\\la\\sigma/{\\rm km/s}\\la 230$, both ellipticals and lenticulars) must have either a young subpopulation or a blue horizontal branch. Based on the above relationships valid for the bulk of the sample, we show that the Mg-$\\sigma$ relation is mainly driven by metallicity, with similar shares from the \\aFe\\ ratio ($23$ per cent) and age ($17$ per cent).  We further find evidence for an influence of the environment on the stellar population properties. Massive early-type galaxies in low-density environments appear on average $\\sim 2\\;$Gyrs younger and slightly ($\\sim 0.05-0.1$~dex) more metal-rich than their counterparts in high density environments. No offsets in the \\aFe\\ ratios, instead, are detected. With the aid of a simple chemical evolution model, we translate the derived ages and \\aFe\\ ratios into star formation histories. We show that most star formation activity in early-type galaxies is expected to have happened between redshifts $\\sim 3$ and 5 in high density and between redshifts 1 and 2 in low density environments. We conclude that at least 50 per cent of the total stellar mass density must have already formed at $z\\sim 1$, in good agreement with observational estimates of the total stellar mass density as a function of redshift.  Our results suggest that significant mass growth in the early-type galaxy population below $z\\sim 1$ must be restricted to less massive objects, and a significant increase of the stellar mass density between redshifts 1 and 2 should be present caused mainly by the field galaxy population.  The results of this paper further imply vigorous star formation episodes in massive objects at $z\\sim 2$--$5$ and the presence of evolved ellipticals around $z\\sim 1$, both observationally identified as SCUBA galaxies and EROs, respectively. ",
        "introduction": "\\label{intro} The most direct way to constrain the formation and evolution of galaxies certainly is to trace back their evolution with redshift \\citep[e.g.][]{Araetal93,BZB96,ZB97,SED98,KBB99,Pogetal99,Zieetal99,Keletal00,vDoketal00,Sagetal00}. The price to be paid, however, is that high-redshift data naturally have lower quality and are therefore more difficult to interpret.  A clear complication is the so-called progenitor bias, which implies that galaxies observed at low and high redshift are not necessarily drawn from the same sample \\citep{vDoketal00}.  The alternative approach is the detailed investigation of the stellar populations in local galaxies, which has been pioneered by analyzing slopes and scatter of color-magnitude and scaling relations of early-type galaxies \\citep[e.g.][]{Dreetal87,DD87,BLE92,BBF92,BBF93,RC93}, followed by a number of detailed studies of absorption line indices (e.g., \\citealt{P89}; \\citealt*{GAS90}; \\citealt*{WFG92}; \\citealt*{DSP93}; \\citealt{CD94}; \\citealt{Rosetal94}; \\citealt{BP95}; \\citealt*{FFI95}; \\citealt*{JFK95}; \\citealt{Greggio97}; \\citealt{Vazetal97}; \\citealt*{TCB98}; \\citealt{KD98}; \\citealt{Mehetal98}; \\citealt{Worthey98}; \\citealt{KA99}; \\citealt{VA99}; \\citealt{Jor99}; \\citealt{Teretal99}; \\citealt{Kun00}; \\citealt{Kunetal01}; \\citealt{Lonetal00}; \\citealt{Traetal00b}; \\citealt{Pogetal01a,Pogetal01b}; \\citealt{Davetal01}; \\citealt{Caretal02}; \\citealt{PS02}; \\citealt{TF02}; \\citealt{Sagetal02}; \\citealt{Cenetal03}; \\citealt*{CRC03}; \\citealt{Mehetal03}; \\citealt*{TMB03b}; and others).  We call this the 'archaeology approach'.  The confrontation with predictions from models of galaxy formation is certainly most meaningful, when the two approaches, the mining of the high-redshift universe and the archaeology of local galaxies set consistent constraints. In this paper we follow the latter approach.  The main challenge in the archaeology of stellar population is the disentanglement of age and metallicity effects. The use of absorption line indices to lift this degeneracy is powerful \\citep{Fabetal85,G93,Worthey94}, but has been up to now hampered by the fact that different metallic line indices yield different metallicities and therefore different ages (see references above).  We have solved this problem by developing stellar population models that include element abundance ratio effects and now allow for an un-ambiguous derivation of age, total metallicity, and element ratios from (Lick) absorption line indices \\citep*[][hereafter TMB]{TMB03a}. A complication that remains, however, is the degeneracy between age and horizontal branch morphology, which stems from the fact that the presence of warm horizontal branch stars (which cannot be excluded) strengthens the Balmer absorption and can mimic a younger stellar population age \\citep{FB95,MT00,Leeetal00}.  We will additionally discuss this problem in the present paper.  The principal aim of this study is to constrain the formation epochs of the stellar populations in early-type galaxies as a function of their type (elliptical and lenticular), mass and environmental density.  For this purpose we analyze a homogeneous, high-quality data sample of 124 early-type galaxies in various environmental densities. With our new stellar population models (TMB) we determine ages, total metallicities, and \\aFe\\ ratios from the absorption line indices \\Hb, \\Mgb, and \\Fe, and seek for possible correlations of these parameters with velocity dispersion.  First tentative results have already been discussed in \\citet*{TMB02b}. Here we present the final and more comprehensive analysis. In particular, we compare the observed data set with mock samples produced through Monte Carlo simulations taking the typical average observational errors into account, in order to eliminate the confusion caused by correlated errors \\citep{Traetal00b,Kunetal01,TF02}.  The \\aFe\\ element ratio plays a key role for the accomplishment of our main goal, namely the derivation of formation epochs, While the so-called $\\alpha$-elements O, Ne, Mg, Si, S, Ar, Ca, Ti (particles that are build up with $\\alpha$-particle nuclei) plus the elements N and Na are delivered mainly by Type~II supernova explosions of massive progenitor stars, a substantial fraction of the Fe-peak elements Fe and Cr comes from the delayed exploding Type~Ia supernovae \\citep*[e.g.][]{NTY84,WW95,TNH96}. Hence, the \\aFe\\ ratio quantifies the relative importance of Type~II and Type~Ia supernovae \\citep*{GR83,MG86,PT95,TGB98}, and therefore carries information about the timescale over which star formation occurs.  Thus, the \\aFe\\ ratio can be considered as an additional measure of late star formation, and we will use it both to constrain formation timescales and to lift the degeneracy between age and horizontal branch morphology.  The paper is organized as follows. In Sections~\\ref{sec:datasample} and~\\ref{sec:montecarlo} we present the methodology of the paper, i.e.\\ we introduce the data sample, we briefly summarize the stellar population model of TMB, and we explain the construction of the mock galaxy sample through Monte Carlo simulations. The results of this analysis are presented in Section~\\ref{sec:results}. We use these results to explore the origin of the Mg-$\\sigma$ relation (Section~\\ref{sec:mgsigma}) and to derive star formation histories (Section~\\ref{sec:sfh}) with the aim to constrain the epoch of early-type galaxy formation. The main results of the paper are discussed and summarized in Sections~\\ref{sec:discussion} and~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We study the stellar population properties of 124 early-type galaxies \\citep{G93,Beuetal02,Mehetal03} with velocity dispersions between 50 and 350$\\;$km/s in both high and low density environments. From the absorption line indices \\Hb, \\Mgb, and \\Fe\\ we derive ages, metallicities \\ZH, and \\aFe\\ element ratios using our abundance ratio sensitive stellar population models (TMB).  The \\aFe\\ ratio is the key to determine the formation timescales of the objects, as it quantifies the relative importance of the chemical enrichment from Type~II and the delayed Type~Ia supernovae. Together with average ages, we use this information to set constraints on the epochs, when early-type galaxies form the bulk of their stars. To get a handle on the effect of correlated errors, we compare the observed sample with mock galaxy samples with properties as close as possible to the observational data in terms of sample size, velocity dispersion distribution, and measurement errors.  For the bulk of the sample we find that all three stellar population parameters age, total metallicity, and \\aFe\\ ratio correlate with velocity dispersion, hence galaxy mass.  We show that these results are robust against recent revisions of the local abundance pattern at high metallicities.  To recover the observed scatter we need to assume an intrinsic scatter of about 20 per cent in age, $0.08\\;$dex in [\\ZH] and $0.05\\;$dex in [\\aFe].  Both zero-point and slope of the \\aFe-$\\sigma$ relation are independent of the environmental density, while the ages of objects in low density environments appear systematically lower accompanied by slightly higher metallicities.  We identify three classes of 'outliers' from these relationships.  1) The low-mass objects with $M_*\\la 10^{10}\\;\\Msun$ ($\\sigma<130\\;$km/s) in our sample show clear evidence for the presence of intermediate-age populations. Their line indices are best reproduced with a 2-component model in which a with a 1.5--2$\\;$Gyr population with low \\aFe\\ ratios is added to a base population defined by the above relationships. 2) About 20 per cent of the intermediate-mass objects with $10^{10}\\la M_*/\\Msun\\la 10^{11}$ ($110\\la\\sigma/{\\rm km/s}\\la 230$) must have either a young subpopulation or a blue horizontal branch. Compared to the low-mass objects, the case for the intermediate-age population is less convincing because of relatively high \\aFe\\ ratios. 3) Most intriguing is a population of massive ($M_*\\ga 10^{11}\\;\\Msun$ or $\\sigma\\ga 200\\;$km/s) S0 galaxies, only present in the high density sample, that require a major fraction of their stars ($\\sim 70$ per cent) to have formed relatively recently at lookback times of about $2$--$3\\;$Gyr, corresponding to $z\\sim 0.2$. Note that these objects, even though being interesting in themselves, represent only a minor fraction of the total stellar mass budget in early-type galaxies.  \\medskip The increase of total metallicity, indicating the completeness of chemical processing, mainly reflects the inability of massive galaxies with deeper potential wells to develop galactic winds. The increase of age and \\aFe\\ ratio with $\\sigma$ shows that more massive galaxies form their stellar populations both earlier and faster. Our results suggest that the well-known Mg-$\\sigma$ relation is mainly driven by metallicity ($\\sim 60$ per cent contribution) with almost equal share from \\aFe\\ ratio ($\\sim 23$ per cent) and age ($\\sim 17$ per cent). The increase of \\aFe\\ with velocity dispersion is responsible for the much weaker correlation of \\Fe\\ with galaxy mass.  With a simple chemical evolution model that takes the delayed enrichment from Type~Ia supernovae into account, we translate the ages and \\aFe\\ ratios of the above relationships into star formation histories as a functions of galaxy's velocity dispersions and environmental densities. We show that the higher ages and \\aFe\\ ratios of more massive early-type galaxies imply faster formation timescales at earlier epochs. The lower the mass of a galaxy is, the more it suffers from episodes of late star formation extending to redshifts below $z\\sim 1$ for galaxies with $M_*\\la 10^{10}\\;\\Msun$. Massive early-type galaxies ($M_*\\ga 10^{11}\\;\\Msun$) in high density environments form their stellar populations between redshifts 2 and~5. Galaxies in low density environments appear on average $1$--$2\\;$Gyr younger than their counterparts in high environmental densities, but obey the same relationship between \\aFe\\ ratio and velocity dispersion. This implies that they form on the same timescales, but the epochs of their formation are delayed to lower redshifts between $z\\sim 1$ and $z\\sim 2$.  \\medskip The 'archaeology approach' pursued in this paper and observations of the high-redshift universe are now converging to a consistent picture. The results of this work imply that at least 50 per cent of the total stellar mass density must have already been formed by $z\\sim 1$, which is in good agreement with observational estimates of the total stellar mass density as a function of redshift. We conclude that significant mass growth in the early-type galaxy population below $z\\sim 1$ must be restricted to less massive objects, and a significant increase of the stellar mass density between redshifts 1 and 2 should be present caused mainly by the field galaxy population.  As shown here, the properties of local galaxies predict the presence of massive objects with very high star formation rates of the order $1000\\;$\\Msun yr$^{-1}$ around redshifts 2--5, and the existence of old evolved galaxies at redshifts between 1 and 2, both now confirmed observationally.  Current models of hierarchical galaxy formation are consistent with global properties like the stellar mass function, but remain challenged by the ages of local massive galaxies , the number densities of luminous red objects at intermediate redshift, and the high star formation rates in massive objects at high redshifts. The problem seems tightly linked to the (uncertain) treatment of baryonic matter in dark halos, and likely solutions lie in the inclusion of AGN feedback and refinement of gas heating/cooling processes."
    },
    "0410/astro-ph0410515_arXiv.txt": {
        "abstract": "We present observations of the singly and doubly deuterated species of thioformaldehyde, HDCS and D$_2$CS, towards the dark cloud Barnard 1. This is the first detection of D$_2$CS in Space and in dense and cold prestellar regions. Column densities obtained using rotational diagrams and a Large Velocity Gradient model show an extremely high D-enhancement in thioformaldehyde in Barnard 1. Although the column density of H$_2$CS is smaller than that of H$_2$CO, both species show similar D-enhancements in their singly and doubly deuterated species. A chemical model -including multiply deuterated species- has been used in order to interpret the observations.  Predicted rotational frequencies from laboratory data for HDCS and D$_2$CS are significantly in error when compared to the observed frequencies in Space. Consequently, we have derived new rotational constants for these two species and for H$_2$CS and H$_2$C$^{34}$S using the observed frequencies in Barnard 1. The new rotational constants allow to predict the rotational transitions of these species with the accuracy needed for the narrow line emerging from dark clouds. Rotational constants for HDCS and D$_2$CS have been obtained from the observed transitions in the laboratory and in Space. ",
        "introduction": "For a long time, singly deuterated interstellar molecules are known to be present in cold molecular clouds with an extremely high abundance enhancement due to chemical fractionation \\citep{van95,Min97,Sai00,Tin00,Tur01,Par02}. The formation of deuterated molecules is preferred at low temperatures ($\\leq$80 K) and leads to a high degree of fractionation in cold and dense dark clouds without embedded luminous sources. These prestellar cores offer the opportunity to study the cold and dense evolutionary stage in which the deuteration of molecules supposedly occurred. Although this is a relatively well understood chemical process, the observation in the last few years of multiply deuterated molecules has renewed interest in deuterium fractionation in dense clouds : ND$_2$H \\citep{Rou00}; ND$_2$H and D$_2$CO \\citep{Loi01}; ND$_3$ \\citep{Lis02}; D$_2$S \\citep{Vas03}; D$_2$CO \\citep{Cec02,Bac03}; CHD$_2$OH and CD$_3$OH \\citep{Par02,Par04}. Among these deuterium isotopic species, there are, so far, only three sulfur bearing deuterated molecules, HDS and D$_2$S, detected towards several protostars and dense cores \\citep{Vas03,van95}; and HDCS, which has been detected toward TMC-1 \\citep{Min97}.  Our new data and, in particular the high enhancement in the fractionation of multiply deuterated species, are requiring a review of chemical models. The presence of deuterated molecular species in star forming regions like Orion-IRc2 and IRAS 16293-2422 \\citep{Tur90,van95,Cec98,Loi00} has been taken as signs for an active grain surface chemistry followed by desorption when the physical processes associated to recent star formation activity increases the temperature of the gas and of dust grains. However, this is not the case for dark clouds, where the kinetic temperature is too low for efficient evaporation of ice grain mantles. \\cite{Rob00} found that gas phase chemistry can account for the large abundances of both singly and doubly deuterated molecules when the gas is very depleted of molecules and atoms. The reason is that both CO and O efficiently destroy H$_3^+$. Hence when CO and O are sufficiently depleted, H$_3^+$ is converted efficiently into H$_2$D$^+$ in the reaction with HD. Subsequent reactions of H$_2$D$^+$ with HD leads to D$_2$H$^+$ and finally to the formation of D$_3^+$. These ions propagate the deuterium to other molecules \\citep{Lis02,Rob03}. In their model, \\cite{Rob03}, conclude that in regions where CO is higly depleted, the abundance of D$_2$H$^+$ will be simillar than that of H$_2$D$^+$ and that D$_3^+$ could be the most abundant form of deuterated H$_3^+$. The first prediction has been confirmed recently with the detection of D$_2$H$^+$ toward the prestellar core 16293E \\citep{Vas04}. For an early evolution time ($\\sim$2-4$\\times$10$^4$ yr) the double and triple deuterated species could have large enhancements \\citep{Rob03}. In the case of Barnard 1, \\cite{Lis02}, based on their C$^{18}$O and C$^{17}$O (2-1) observations found that C$^{18}$O/H$_2$=5.4$\\times$10$^{-8}$, which is $\\sim$3 times lower than the canonical value, C$^{18}$O/H$_2$=1.7$\\times$10$^{-7}$ \\citep{Fre82, Cer87}, indicating that an appreciable amount of CO is depleted onto grains. Modelling with low temperature conditions and high depletion factors allows a satisfactory comparison with the observations of doubly and triply deuterated ammonia -see e.g. NHD$_2$ towards L134N \\citep{Rou00}- and also helps to explain the high abundances observed in star forming regions in terms of grain chemistry models. The recent detection of doubly deuterated H$_2$S provides a stimulating impetus to consider other sulfur molecules and extend the previously elaborated chemical schemes.  During a 3mm line survey on dark clouds (Barnard 1, TMC-1, L1544, L183 -Marcelino et al., in prep.), we found several transitions that initially were considered as unidentified. Taking into account the reduced number of possible candidates, the fact that the lines were strongest in Barnard 1 (see Figure 1), and the strong deuterium enhancement found in these sources by several authors \\citep[see, e.g.,][and references therein]{Rou00,Ger01}, we concluded that these lines were arising from D$_2$CS and HDCS. While single deuterated thioformaldehyde had been detected at lower frequencies in TMC-1 \\citep{Min97}; this is now the first detection of D$_2$CS in Space. Once the carriers were identified, and given the extremely large deuterium enhacement indicated by the observations, we concentrated our efforts and available observing time in the source showing the strongest lines, i.e., Barnard 1. In Section 2 we present the observing procedures and in Section 3 the main results. In this section we discuss the rotational constants of thioformaldehyde and its isotopomers. We also analyze the physical and chemical conditions of Barnard 1 and derive the column densities for the different thioformaldehyde isotopic species. Section 4 is devoted to the discussion and to the chemical analysis of the results. ",
        "conclusions": "We have carried out a comprehensive observational study in Barnard 1 of the deuterated species of thioformaldehyde, together with the observation of several lines from the main isotope and from the $^{34}$S isotopic substituted species. In order to compare the deuterium enhancement we have also observed formaldehyde and its deuterated species. The main results are:  1) The systematic velocity of the dense gas in Barnard 1 has been determined to be 6.70$\\pm$0.05 km s$^{-1}$ from the observation of near 20 optically thin molecular lines with well known frequencies.  2) The observed frequencies for HDCS, D$_2$CS and H$_2$C$^{34}$S as well as the available laboratory data have been fit to obtain a set of molecular rotational constants allowing precise frequency determinations for these species in the millimeter domain.  3) Rotational temperatures and column densities for H$_2$CS, H$_2$CO and their isotopes have been derived from rotational diagrams or through the use of a LVG code. The derived H$_2$CS/HDCS and H$_2$CS/D$_2$CS abundance ratios are $\\sim$3 and 9 respectively, while for formaldehyde are $\\sim$7 and 13.  4) Steady state gas phase chemical models including multiply deuterated species and with appropriate depletion factors, such as those found in different dense cores, allow to reproduce the deuterium fractionation approximately. In this context, surface processes have not to be included specifically , as they are poorly known, and their role is essentially passive as the abundant oxygen, CO, H$_2$O ...  components are frozen on the grains and are not available for gas phase chemical processes to proceed. Our results involve also some specific chemical assumptions which contribute to the high deuterium fractionation, namely (1) specific branching ratios in dissociative recombination of deuterated molecular ions with a 1:2 ratio for the ejection of H compared to D as observed in several experiments. And (2) fractionation reactions of H$_2$CO$^+$, H$_3$CO$^+$, H$_3$S$^+$, H$_2$CS$^+$ with HD. These processes have not been studied in the laboratory nor theoretically and such studies would be very desirable. However, their influence on our results is not determinant except for the deuterium fractionation of H$_2$S."
    },
    "0410/astro-ph0410273_arXiv.txt": {
        "abstract": "Extensive radio data of two Algol systems and two RS~CVn binaries were re-analyzed. We found evidence for a new periodicity that we interpret as a superhump in $\\algol$=Algol, in which it may have been expected according to its semi-detached nature and low binary mass ratio. The concluded presence of an accretion disk (or an annulus) is consistent with previous studies of optical data and numerical simulations of $\\algol$. In our model, the 50-d period, previously found in the radio data of $\\algol$, is explained as the apsidal precession of the elliptical accretion disk / annulus. If our interpretation is correct, this is the first detection of the superhump phenomenon in the radio and the first observation of superhumps in Algol systems. According to our result, the accretion disk / annulus in $\\algol$ precesses in spite of its non-Keplerian nature and therefore this phenomenon is not restricted to the classical Keplerian accretion disks in compact binaries.   We propose that in Algol systems with short orbital periods, which have accretion disks / annuli, the disk is magnetically active as well as the cool secondary star. The magnetic field in the disk originates from amplification of the seed field in the magnetized material transferred from the secondary. The disk and stellar fields interact with each other, with reconnection of the field lines causing flares and particle acceleration. Relativistic particles are trapped in the field and directed toward the polar regions of the secondary star because of the dipole structure of its magnetic field. These systems are, therefore, somewhat analogous to the RS~CVn systems, which have two magnetically interacting stars. Our proposed model for the magnetic activity in Algol systems provides a simple explanation to the observed properties of $\\algol$ in the radio wavelengths, and in particular, to the presence of quiescent gyrosynchrotron emission near the polar region of the secondary star, where electrons are difficult to be confined if the field lines are open as in normal single magnetic stars. It further explains the recent discovery that the Doppler shifts of the emission lines in the X-ray data of $\\algol$ are somewhat lower than what is expected from the orbital motion of the secondary star. We propose that the superhump variation in the radio is generated by enhanced reconnection when the elongated side of the elliptic accretion disk is the closest to the cool star. This leads to flares and enhancement in particle acceleration and is manifested as stronger gyrosynchrotron radiation.  The observed superhump period, at 3.037$\\pm$0.013 d, $\\sim$6\\% longer than the orbital period, was used to deduce a binary mass ratio of $\\sim$0.23 in $\\algol$, which is consistent with previous studies. Our finding opens the possibility to extend the superhump phenomenon to Algol systems and to test the theories of precessing accretion disks in various types of interacting binaries. It provides a new method to estimate the mass ratios in these binaries. It also offers new insights and improves our understanding of the complicated magnetic interaction and feedback between mass transfer dynamics, time-dependent disk accretion, and feedback and induced magnetic activity in the Algols and related systems. ",
        "introduction": "In semi-detached systems, one star fills its Roche-Lobe and transfers mass to the other. There are several different groups of interacting binaries that are classified according to the properties of their component stars. In cataclysmic variables (CVs), a white dwarf accretes mass from a red dwarf \\citep{W1995}. In low mass X-ray binaries (LMXBs), the mass-receiver star is a black hole or a neutron star and the donor is typically a low mass solar-like star \\citep{LPH1995}. Algol systems contain a hot main-sequence star and a cool subgiant or giant companion \\citep[e.g.,][]{PP1984}. The mass transfer in interacting binaries is usually sustained through an accretion disk which is formed around the mass-receiver star. In CVs and LMXBs very often the disk is the dominant light source in the optical regime.   Many CVs show in their optical light curves quasi-periodicities a few percent different than the orbital periods \\citep{W1995, P1999, RN2000, PTK2003}. For historical reasons, these periods are known as superhumps. These humps were initially observed in the bright outbursts (superoutbursts) of SU~UMa systems, a subclass of dwarf novae, which are a subgroup of CVs. \\citet{OC1996} tested and confirmed the optical detections of superhumps in three LMXBs. \\citet{HKM2001} suggested that many more LMXBs share this property. There is, however, only one clear detection (in the LMXB V1405~Aql=X1916-053) of superhumps in the X-ray regime \\citep{RCB2002}. When the superhump periods are longer than the orbital periods they are termed `positive superhumps'. Similarly, periods shorter than the orbital periods are called `negative superhumps'. Hereafter, we refer to superhumps as positive superhumps, which are the more common phenomenon. The superhump is understood as the beat periodicity between the orbital period and the apsidal precession of an elliptical accretion disk \\citep{O1985, W1988}. It is argued that superhumps can only occur in systems with small mass ratios: M$_{donor}$/M$_{receiver}\\leq$ 0.33 \\citep{WK1991, M2000}, although \\citet{RHA2003} raised some questions about this limit by finding superhumps in TV~Col, a CV with a relatively large orbital period, which probably implies a mass ratio above the theoretical limit.   Algol systems are named after the proto-type Algol ($\\algol$) and may be progenitors of CVs \\citep{IT1985}. Algol binaries typically contain a hot B--A main-sequence star and a cool F--K subgiant or giant which show a substantial magnetic activity \\citep[e.g.,][]{PP1984, RWG2003}. As in CVs and LMXBs, the mass transfer in the Algol systems with orbital periods longer than about 5 days is sustained through an accretion disk. In systems with shorter periods the disk may be transient, and in some cases it does not follow the Keplerian velocity field (v $\\propto r^{-1/2}$) as in compact binaries, so it was termed `an annulus' \\citep{PP1984, R1992, R1993, RA1999}. Hydrodynamic simulations demonstrate that the mass transfer in the Algol binaries can indeed result in the formation of an accretion disk, a transient disk or an annulus around the mass-receiver star \\citep{BRM1995, RR1998} and Doppler tomography supports this suggestion \\citep{RAB1995, RJS1996}. We note that the numerical simulations show that the accretion material around a $\\algol$-type system can develop an asymmetrical structure that looks elliptical in shape \\citep[e.g., see fig.~1 of][]{BRM1995}.  Eccentric accretion disks may, therefore, be expected in Algol systems whose mass ratios are below the critical value as well. However, in these objects the contribution of the disk in the optical band is very weak compared with the bright binary stars \\citep{R1992}. Thus, to look for superhumps it may be preferred to study other wavebands where the disk may have a larger impact on the light curve. The radio regime may, therefore, offers an opportunity to search for superhumps in Algol binaries.   In this work, extensive radio observations presented by \\citet{RWG2003} were re-analyzed. The data cover four binary systems, two of which are Algol systems, and the other two belong to the RS~CVn class. The RS~CVn systems consist of two detached magnetically active stars, and thus should not have an accretion disk nor superhumps. The mass ratios of the Algol systems studied, $\\algol$ and $\\delta$~Lib, are 0.22 and 0.35 respectively \\citep{RWG2003}. Therefore, $\\algol$ has a mass ratio well below the theoretical limit of 0.33 and it should be the only system among the four that may show superhumps, although the mass ratio of $\\delta$~Lib is just above the critical limit.    $\\algol$ (Algol, HD 19356) at V=2.1-3.4 and d=28.76 pc is the brightest eclipsing system and the brightest and closest semi-detached binary \\citep{L1979, MM1998, RA1999, RWG2003}. It has strong emission across the wavelength spectrum including the X-ray, ultraviolet, optical, infrared and radio \\citep[e.g.,][]{LMP1988, R1990, SLS1995, NSB2002, RWG2003, SNF2003}. $\\algol$ is actually a hierarchical triple system with a B8 V main-sequence star, a K2 IV subgiant companion and a binary orbital period of 2.8673 days. The tertiary, which is an F1 IV star, orbits around the binary system with a period of 680.08 days \\citep{HBH1971, RMB1988, R1992, R1993, RWG2003}. The binary mass ratio of $\\algol$ was estimated as 0.217$\\pm0.005$ \\citep{HBH1971, TL1978} and 0.22$\\pm0.03$ \\citep{RMB1988}. These values locate the system well inside the permitted range for superhumps, which would be predicted to be about 4--8\\% longer than the orbital period \\citep{O1985, W1994, M2000}. In this work, we present evidence for such a signal. ",
        "conclusions": "We found evidence for a new periodicity about 6\\% longer than the orbital period in extensive existing radio photometry of $\\algol$. The period is seen in the flaring data but not in quiescence. This peak is interpreted as a positive superhump period. The detection of superhumps in the radio light curve of $\\algol$ implies that it had some sort of an asymmetrical accretion structure that maintained its eccentricity for nearly six years. We propose that it is an accretion disk or a ring. This suggestion is consistent with interpretation of optical data, Doppler tomography and numerical simulations that imply that $\\algol$ has an accretion annulus around its primary B star. Therefore, Algol systems with short orbital periods may have accretion disks / rings / annuli that may be more stable than what was previously believed. In our model, the 50-d period in the radio data is explained as the apsidal precession of the accretion disk / annulus.  Our results extend the superhump phenomenon, which has been observed in Keplerian accretion disks in compact binaries, to accretion structures with sub-Keplerian velocities. It is also the first detection of superhumps in an Algol system and the first time this phenomenon is seen in the radio.  A new model is proposed for the magnetic activity of Algol-like systems. According to our scenario, in addition to the strong magnetic field of the secondary star, the magnetic field carried in the accreting material is dragged, twisted and amplified in the accretion disk / annulus around the primary star. The reconnection of the field lines accelerates electrons to relativistic energies and trapping occurs near the magnetic poles of the cool star. The binary system is thus reminiscent of RS~CVn systems, which consist of two magnetically active stars. This model explains several peculiarities in the radio emission of $\\algol$. In fact, it may further explain the low flaring activity of $\\delta$~Lib \\citep{RWG2003} by the fact that it has a binary mass ratio above the theoretical limit for superhumps (Sections 1--3).  Our model for the formation of the superhump signal in the radio light curves of Algol systems, namely that the reconnection of the field lines varies with the phase of the eccentricity, yields a prediction that Algol binaries with intermediate orbital periods (thus having a transient disk or an accretion annulus) and mass ratios below the critical value for superhumps, 0.33 (Section 1), should be somewhat stronger and harder radio sources than the other objects. Similarly, the radiation in other wavebands may be enhanced compared with systems with larger mass ratios.  Our scenario for $\\algol$ can further explain the fact that the X-ray Doppler shifts in emission lines of its secondary star are slightly smaller than what would be predicted from its orbital rotation. The accretion disk annulus is proposed to contribute about 10-15\\% of the total X-ray light. Therefore, we may predict that the X-ray light curve of Algol will be modulated with the superhump period as well.  The binary mass ratios in non-eclipsing Algol systems might be estimated using the superhump phenomenon. Our results should also help to explain some of the observed peculiarities of these systems in the infrared, radio, optical, ultraviolet and X-ray bands and shed new light on the magnetic activity in binary systems and its interaction with the accretion structures.  We hope that future radio and X-ray observations of $\\algol$ and other Algol-type systems with low mass ratios that should have superhumps would be able to confirm or refute our suggestions."
    },
    "0410/astro-ph0410045_arXiv.txt": {
        "abstract": "Transient radio sources are necessarily compact and usually are the locations of explosive or dynamic events, therefore offering unique opportunities for probing fundamental physics and astrophysics. In addition, short-duration transients are powerful probes of intervening media owing to dispersion, scattering and Faraday rotation that modify the signals. While radio astronomy has an impressive record obtaining high time resolution, usually it is achieved in quite narrow fields of view. Consequently, the dynamic radio sky is poorly sampled, in contrast to the situation in the X-ray and $\\gamma$-ray bands. The SKA has the potential to change this situation, opening up new parameter space in the search for radio transients. We summarize the wide variety of known and hypothesized radio transients and demonstrate that the SKA offers  considerable power in exploring this parameter space. Requirements on the SKA to search the parameter space include the abilities to (1)~Make targeted searches using beamforming capability; (2)~Conduct blind, all-sky surveys with dense sampling of the frequency-time plane in wide fields; (3)~Sample the sky with multiple fields of view from spatially well-separated sites in order to discriminate celestial and terrestrial signals; (4)~Utilize as much of the SKA's aggregate collecting area as possible in blind surveys, thus requiring a centrally condensed configuration; and (5)~Localize repeating transient sources to high angular precision, requiring a configuration with long baselines, thus requiring collecting area in both a centrally condensed ``core'' array and sufficient area on long baselines. ",
        "introduction": "\\label{sec:drs.intro}  Transient emission---bursts, flares, and pulses on time scales of order 1~month---marks compact sources or the locations of explosive or dynamic events.  As such, radio transient sources offer insight into a variety of fundamental physical and astrophysical questions including \\begin{itemize} \\item The mechanisms of efficient particle acceleration; \\item Possible physics beyond the Standard Model; \\item The nature of strong field gravity; \\item The nuclear equation of state; \\item The cosmological star formation history; \\item Detecting and probing the intervening medium(a); and \\item The possibility of extraterrestrial civilizations. \\end{itemize}  A figure of merit for transient detection is $A\\Omega(T/\\Delta t)$, where $A$ is the collecting area of the telescope, $\\Omega$ is the solid angle coverage in the search for transients, $T$ is the total duration of observation, and~$\\Delta t$ is the time resolution. Effective detection of transients requires \\begin{equation} A\\Omega\\left(\\frac{T}{\\Delta t}\\right) \\to \\mathrm{``large\\char'42} \\label{eqn:drs.fom} \\end{equation} At high energies (X- and $\\gamma$-rays), detectors with large solid angle coverage and high time resolution have had great success in finding classes of transient objects.  At optical wavelengths there has been recent progress in constructing wide-field detectors with high time resolution.  Historically, radio telescopes have been able to obtain high time resolution (with some modern telescopes achieving nanosecond time resolution) or large fields of view, but rarely have both high time resolution and large field of view been obtained simultaneously.  Nonetheless, there are a number of indications that the radio sky may be quite dynamic.  Radio observations of sources triggered by high-energy observations (e.g., radio observations of gamma-ray burst afterglows), monitoring programs of known high-energy transients (e.g., radio monitoring of X-ray binaries), giant pulses from the Crab pulsar, a small number of dedicated radio transient surveys, and the serendipitous discovery of transient radio sources (e.g., near the Galactic center) suggest that the radio sky is likely to be quite active on short time scales.  There also may be unknown classes of sources.  In this chapter, we discuss radio transients from the standpoint of the available phase space that the SKA can probe and how the design of the SKA can be optimized to improve its study of both known and suspected classes of transient radio sources.  Several specific classes of radio transients---including pulsars, supernovae and gamma-ray bursts, ultra-high energy cosmic rays, and scintillation-induced variability---are discussed in more depth in additional chapters.  We summarize the wide variety of both known and hypothesized radio transients in \\S\\ref{sec:drs.sources}, the available phase space that the SKA can probe in \\S\\ref{sec:drs.phasespace}, and the operational modes and aspects of transient detection for the SKA in \\S\\ref{sec:drs.survey}. ",
        "conclusions": "\\label{sec:drs.conclude}  Transient radio sources offer insight into dynamic or explosive events as well as being powerful probes of intervening media. Historically, the radio sky also has been searched only poorly for radio transients, meaning that the SKA has the potential to explore previously-unexplored parameter space at radio wavelengths.  In order to be effective at searching for transient radio sources, the SKA must achive a ``large'' value for the figure of merit $A\\Omega(T/\\Delta t)$, equation~(\\ref{eqn:drs.fom}).  By its very nature, $A$ is anticipated to be large for the \\hbox{SKA}.  Table~\\ref{tab:drs.survey} summarizes the other requirements on the \\hbox{SKA}.  These requirements flow either from the desire to achieve a large figure of merit or are based on experience from known radio transients.  \\begin{table} \\caption{SKA Requirements for Transient\tSurveys\\label{tab:drs.survey}} \\begin{center} \\begin{tabular}{lc} \\noalign{\\hrule} Parameter & Requirement \\\\ \\noalign{\\hrule\\hrule} Field of View (at~1~GHz) & 1~deg.${}^2$ \\\\ Time Sampling            & $\\sim 1$~ns to 10~$\\mu$s \\\\ Frequency Coverage       & $\\sim 0.1$--10~GHz \\\\ Configuration            & ``large'' core filling factor \\\\ & long baselines \\\\ \\noalign{\\hrule} \\end{tabular} \\end{center} \\end{table}  In order to conduct successfully the Key Science Project, ``Strong Field Tests of Gravity Using Pulsars and Black Holes,'' the SKA is required to obtain high time resolution.  The challenge for the \\hbox{SKA}, and historically the difficulty with radio telescopes, is that they have not provided a large field of view simultaneously with high sensitivity.  A related requirement is on the configuration of the array, as the filling factor of the array affects the signal-to-noise ratio in a search as $\\sqrt{f}$.  Thus, the SKA requires a core, containing a significant fraction of the collecting area with a reasonablely high filling factor, as well as long baselines, for localizing sources and imaging searches such as for $\\gamma$-ray bursts and extrasolar planets.  From known classes of sources alone, this search promises to be fruitful---one can explore particle acceleration by observing radio pulses from ultra-high energy particles impacting the Earth's atmosphere as well as from flares from the Sun, brown dwarfs, flare stars and X-ray binaries while using afterglows from $\\gamma$-ray bursts one can explore the cosmological star formation history.  Likely, though not yet detected, classes of sources include giant pulses from extragalactic pulsars, which would allow a direct probe of the local intergalactic medium, and extrasolar planets, which would be a direct detection of these objects.  Most exciting would be the \\emph{discovery of new classes of sources}."
    },
    "0410/astro-ph0410335_arXiv.txt": {
        "abstract": "In 2003, Kuan, Charnley, and co-workers reported the detection of interstellar glycine (NH$_2$CH$_2$COOH) based on observations of 27 lines in 19 different spectral bands in one or more of the sources Sgr B2(N-LMH), Orion KL, and W51 e1/e2.  They supported their detection report with rotational temperature diagrams for all three sources.  In this paper, we present essential criteria which can be used in a straightforward analysis technique to confirm the identity of an interstellar asymmetric rotor such as glycine. We use new laboratory measurements of glycine as a basis for applying this analysis technique, both to our previously unpublished 12 m telescope data and to the previously published SEST data of Nummelin and colleagues. We conclude that key lines necessary for an interstellar glycine identification have not yet been found. We identify some common molecular candidates that should be examined further as more likely carriers of several of the lines reported as glycine. Finally, we illustrate that a rotational temperature diagram used without the support of correct spectroscopic assignments is not a reliable tool for the identification of interstellar molecules. ",
        "introduction": "Interstellar glycine (NH$_2$CH$_2$COOH) has been unsuccessfully sought using both single-element telescopes (Brown et al.  1979; Hollis et al.  1980; Snyder et al.  1983; Beralis et al.  1985; Guelin \\& Cernicharo 1989; Combes, Rieu, \\& Wlodarczak 1996; Ceccarelli et al.  2000) and interferometric arrays (Snyder 1997; Hollis et al.  2003a,b).  Recently, Kuan et al.  (2003) reported the identification of interstellar glycine (NH$_2$CH$_2$COOH) in the direction of the hot molecular cores Sgr B2(N-LMH), Orion KL, and W51 e1/e2.  We note several unusual aspects about the Kuan et al. (2003) reported glycine results.  First, their glycine report was based on the detection of 27 spectral line features in the sources Orion~KL, W51 e1/e2, and Sgr~B2(N-LMH). They used the NRAO\\footnotemark[7]\\footnotetext[7]{The National Radio Astronomy Observatory (NRAO) is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} 12~m radio telescope operating in the range of 130 to 242 GHz. The 27 independent spectral line features detected in 3 different sources might seem to suggest reasonable evidence. Certainly the detection of the same ensemble of spectral lines in multiple sources is a characteristic of interstellar molecules with structural components similar to glycine, such as formic acid (HCOOH) and acetic acid (CH$_{3}$COOH)(see, for example, Liu, Mehringer, \\& Snyder 2001; Remijan et al.  2002). However, each of the three glycine sources was observed to contain no more than 13 to 16 of the 27 independent spectral lines features and only 3 of the 27 spectral line features were found to be common to all 3 sources. Second, 16 of the 27 reported glycine features are not the four-fold degenerate rotational transitions that should be the most likely to be detected for R-branch asymmetric rotor transitions. Third, the glycine rest frequencies used in their observations were based on a Hamiltonian calculation which was extrapolated over 100 GHz.  This method can yield potentially large uncertainties in the frequency predictions.  Fourth, the glycine column densities derived from rotational temperature diagrams were surprisingly high. In fact, Orion KL was reported to have the highest glycine column density, but the interstellar molecule with structure closest to glycine, acetic acid, has never been detected there (Remijan et al.  2003).  In this paper, we investigate the above concerns by starting from new laboratory measurements of glycine that improve the Hamiltonian fit and thereby result in highly reliable interpolated frequency predictions which provide the foundation for our discussion.  We list and explain rigorous and essential criteria for identifying a new interstellar molecule such as glycine. Then we use these criteria on key glycine spectral lines reported by Kuan et al. (2003) to interpret independent observational data for Orion KL, W51 e1/e2, and Sgr~B2(N-LMH), and show that their key lines fail the tests required for a correct glycine identification.  In investigating alternative identifications, we find that the carriers of several important spectral lines which have been assigned to glycine are more logically assigned to other, more common species.  Finally, we find that random choices of unidentified spectral lines can be used to generate a reasonable glycine rotational temperature diagram similar to those published by Kuan et al.  (2003). ",
        "conclusions": "As we discussed in \\S 1, there are several pecularities if the lines reported by Kuan et al. (2003) in Sgr B2 (N-LMH), Orion KL, and W51el/e2 belong to interstellar glycine.  Therefore, we examined their assignments of key glycine transitions to the spectral lines listed in Table 1.  We noted that their lack of correction for beam dilution amounted to the tacit assumption that their glycine sources have an extended source size with $\\Theta_s$ $\\geq$ 45$\"$. In Orion KL, we found that none of the 15 glycine line assignments could be verified because the reported intensities of all 15 lines predict that a major glycine quartet (J=19-18, K$_{-1}$ = 0 or 1) should have been detected in our Orion data from the NRAO 12~m telescope (Figures 1 and 2).  So we can conclude that the Orion spectral lines reported by Kuan et al. (2003) (column 8, Table 1) are not glycine.  Previously, the VLA data of Hollis et al. (2003b) established upper limits in Orion for compact glycine sources.  Therefore, glycine has not been verified in either extended or in compact sources in Orion. In W51 el/e2, 16 lines were assigned to glycine.  Fifteen of the 16 line intensities lend to the predictions that the J=19-18, K$_{-1}$ = 0 or 1 glycine quartets should have been detected in our W51 data from the NRAO 12~m telescope (Figures 3 and 4). Since our spectra show a negative result, we can conclude that the reported W51 spectral lines are also not glycine (column 9, Table 1). In the direction of Sgr~ B2(N-LMH), only 13 lines were assigned to glycine but the situation is slightly more complicated by the high spectral line density.  However, it can be simplified by recognizing that, based on both line strength values and intensity calculations, the two strongest glycine lines should be those assigned to lines 26 and 27 in column 7 of Table 1. Line 26 appears with a peak intensity of 300~mK in the Sgr~B2(N) band scan data of Nummelin et al. (1998), but the important connected lines, J=42-41 and 41-40 in the K$_{-1}$ =l or 2 series, are missing altogether (Table 3).  Furthermore, line 27 should be stronger than line 26, but it too is missing, along with the connected J=43-42 and 42-41 lines in the K$_{-1}$ = 0 or 1 series (Table 4).  All of these key glycine lines strongly connect major K$_{-1}$ = 0, 1, or 2 rotational levels, so their absence is not possible if the other Sgr B2(N-LMH) glycine lines are assigned correctly. Therefore, we conclude that the Sgr B2(N-LMH) spectral lines reported by Kuan et al.  (2003) (column 7, Table 1) are not glycine.  In addition, the VLA data of Hollis et al. (2003a) established upper limits for compact glycine sources in Sgr B2(N-LM).  Consequently, glycine has not been verified in Sgr B2(N-LMH) either in extended or in compact sources.  If the 27 lines reported by Kuan et al.  (2003) in Orion KL, W51el/e2, and Sgr~B2(N-LMH) are not glycine, what are they?  Lines 21, 24, and 26 in Table 1 are the 3 lines reported to be common to all 3 sources.  We have argued that vibrationally excited ethyl cyanide (CH$_3$CH$_2$CN) is a logical assignment for line 21, and vibrationally excited vinyl cyanide (CH$_2$CHCN) for line 24.  The intensity of line 26 increases as the observational beam width narrows, so it could also be a constituent of a somewhat compact source. Four other reported glycine lines could be possibly assigned to more common species listed in Table 5.  Beyond this, it is difficult to comment on other possible line carrier candidates because Kuan et al. (2003) only published sample spectra for 11 of their 27 lines. Finally, we have demonstrated that rotational temperature diagrams can not by themselves be used to identify interstellar molecules.  \\bigskip"
    },
    "0410/astro-ph0410429_arXiv.txt": {
        "abstract": "The identification of the astrophysical site and the specific conditions in which r-process nucleosynthesis takes place remain unsolved mysteries of astrophysics. The present paper emphasizes some important future challenges faced by nuclear physics in this problem, particularly in the determination of the radiative neutron capture rates by exotic nuclei close to the neutron drip line and the fission probabilities of heavy neutron-rich nuclei. These quantities are particularly relevant to determine the composition of the matter resulting from the decompression of initially cold neutron star matter. New detailed r-process calculations are performed and the final composition of ejected inner and outer neutron star crust material is estimated. We discuss the impact of the many uncertainties in the astrophysics and nuclear physics on the final composition of the ejected matter. The similarity between the predicted and the solar abundance pattern for $A\\ge 140$ nuclei as well as the robustness of the prediction with varied input parameters makes this scenario one of the most promising that deserves further exploration. ",
        "introduction": "The rapid neutron-capture process, or r-process, is known to be of fundamental importance for explaining the origin of approximately half of the $A>60$ stable nuclei observed in nature. In recent years nuclear astrophysicists have developed more and more sophisticated r-process models, eagerly trying to add new astrophysical or nuclear physics ingredients to explain the solar system composition in a satisfactory way. The r-process remains the most complex nucleosynthetic process to model from the astrophysics as well as nuclear-physics points of view. The site(s) of the r-process is (are) not identified yet,  all the proposed scenarios facing serious problems. Complex---and often exotic---sites have been considered in the hope of discovering  astrophysical environments in which the production of neutrons is large enough to give rise to a successful r-process. Progress in the modelling of type-II supernovae and $\\gamma$-ray bursts has raised a lot of excitement about the so-called neutrino-driven wind model. However, until now no r-process can be simulated {\\it ab initio} without having to call for an arbitrary modification of the model parameters, leading quite often to physically unrealistic scenarios.  On top of the astrophysics uncertainties, the nuclear physics of relevance for the r-process is far from being under control. The nuclear properties of thousands of nuclei located between the  valley of $\\beta$-stability and the neutron drip line are required. These include the ($n,\\gamma$) and ($\\gamma, n$) rates, $\\alpha$- and $\\beta$-decay half-lives, rates of $\\beta$-delayed single and multiple neutron emission, and the probabilities of neutron-induced, spontaneous, and $\\beta$-delayed fission.  When considering complex astrophysics sites like the neutrino-driven wind, proton-, $\\alpha$-, and neutrino-capture rates need to be estimated, too.  New developments of both nuclear physics and astrophysics aspects of the r-process are discussed in this paper. Section 2 is devoted to the estimate of the  neutron-capture rates and the fission probabilities of exotic neutron-rich nuclei. These properties are of particular relevance for determining the still poorly-known composition that results from the decompression of neutron-star (NS) crust matter, as shown in Sect.~3. ",
        "conclusions": "Most of the problems faced in understanding the origin of r-process elements and observed r-abundances are related to our ignorance of the astrophysical site that is capable of providing the required  large neutron flux. In this respect, understanding the r-process nucleosynthesis is essentially an astrophysics issue that will require improved hydrodynamic models to shed light on possible scenarios. Different sites have been proposed, the currently most favoured ones being related to neutrino-driven outflow during supernova or $\\gamma$-ray burst explosions. Nevertheless, we have shown here that the decompression of initially cold NS matter is also extremely promising. In particular, it provides suitable conditions for a robust r-processing. The resulting r-abundance distribution is very similar to the solar one, at least for $A>140$ nuclei. The underlying nuclear mechanisms, however, differ significantly from  those acting in previous scenarios, in particular photoreactions do not play a major role. The similarity between the predicted and solar abundance patterns as well as the robustness of the prediction against variations of input parameters make this site one of the most promising that deserves further exploration with respect to various aspects such as nucleosynthesis, hydrodynamics and galactic chemical evolution. From the nuclear point of view, this site also represents a new challenge, since it will be necessary to determine $\\beta$-decay, neutron-capture and fission rates for nuclei which, in contrast to the more conventional sites, are situated right at the neutron drip line.  {\\small"
    },
    "0410/astro-ph0410103_arXiv.txt": {
        "abstract": "Chameleons are scalar fields whose mass depends on the environment, specifically on the ambient matter density. While nearly massless in the cosmos, where the matter density is tiny, their mass is of order of an inverse millimeter on Earth, where the density is high. In this note, we review how chameleons can satisfy current experimental constraints on deviations from General Relativity (GR). Moreover, we study the cosmological evolution with a chameleon field and show the existence of an attractor solution, akin to the tracker solution in quintessence models. We discuss how chameleons can naturally drive the observed acceleration of the universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410158_arXiv.txt": {
        "abstract": "We describe a technique to identify Be stars in open clusters using Str\\\"omgren $b$, $y$, and narrow-band H$\\alpha$ photometry.  We first identify the B-type stars of the cluster using a theoretical isochrone fit to the $(b-y, y)$ color-magnitude diagram.  The strongest Be stars are easily identified in a $(b-y, y-\\rm H\\alpha)$ color-color diagram, but those with weaker H$\\alpha$ emission (classified as possible Be star detections) may be confused with evolved or foreground stars.  Here we present such photometry plus H$\\alpha$ spectroscopy of members of the cluster NGC 3766 to demonstrate the accuracy of our technique. Statistical results on the relative numbers of Be and B-type stars in additional clusters will be presented in a future paper. ",
        "introduction": "Be stars are broadly defined as non-supergiant (luminosity class III$-$V) B-type stars that show or have shown Balmer emission \\citep{porter2003}, although a few supergiant Be stars are also known \\citep{negueruela2004}.  Their emission is due to a circumstellar disk that is often highly variable on both short and long time scales. The emission lines change intensity, and the relative intensity of their double-peaked profiles changes, over short time scales of days and even hours.  In addition, the disks themselves have been observed to disappear and reappear over years or decades \\citep{slettebak1988}. Be stars are well known to be a class of rapidly rotating stars \\citep*{slettebak1949, slettebak1966, slettebak1992}, and there are three possible reasons for their rapid rotation.  They may have been born as rapid rotators \\citep{zorec1997}, spun up by binary mass transfer \\citep{pols1991}, or spun up during the main-sequence (MS) evolution of B stars \\citep{meynet2000, fabregat2000}.  Many observational surveys have searched for Be stars in open clusters using multi-band photometry.  \\citet{shobbrook1985, shobbrook1987} used plots of $c_1$ vs. $\\beta$, $c_1$ vs. $V$ magnitude, and $\\beta$ vs. $V$ to identify Be stars in NGC 3766. Many recent surveys have relied on color-color diagrams, rather than Shobbrook's color-magnitude diagrams, to detect Be stars more efficiently. \\citet*{grebel1992} used a technique somewhat similar to ours (\\S3) to identify Be stars in NGC 330.  They used Str\\\"omgren $b, y$, and a narrow-band H$\\alpha$ filter for their observations, and they plotted a color-color diagram for the cluster using $b-y$ vs H$\\alpha-y$.  They were able to verify almost all of their detections with low resolution spectra, demonstrating the accuracy of this technique.  Several other surveys (\\citealt*{grebel1997, dieball1998, keller1999, keller2000, grebel2000, pigulski2001}) have used variations of the color-color diagram technique to identify Be stars in open clusters.  \\citet{grebel1997} emphasizes that the number of Be stars detected in this type of surveys actually represents a lower limit on the true number of Be stars for several reasons.  Most importantly, since the Be state is not permanent, only those stars with current disk outbursts will be detectable as Be stars. In addition, the number of late type Be stars is likely to be incomplete since they are fainter, and they also tend to have weaker H$\\alpha$ emission than early Be stars.  In this paper, we present an improved photometric technique based on $(b-y,y)$ color-magnitude and $(b-y, y-\\rm H\\alpha)$ color-color diagrams to detect Be stars in open clusters.  While our method is very similar to that of \\citet{grebel1992}, we have enhanced it by using cluster isochrone fitting to identify the B-type stars in each cluster.  We have also used a fit of the color-color curve to identify H$\\alpha$ emitters among the B star population.  To demonstrate the accuracy of our technique, we present spectra of 20 stars in the cluster NGC 3766 and correlate their H$\\alpha$ equivalent widths with their $y-\\rm H\\alpha$ colors.  Our photometry of additional clusters will be presented in a future paper, and at that time we will use our statistics to discuss the implications on the possible Be star formation scenarios. ",
        "conclusions": "Our photometric technique for identifying Be stars is very efficient for finding those stars with strong H$\\alpha$ emission, as we have shown for NGC 3766.  While we correctly identify 63\\% of the previously known, currently active Be stars in NGC 3766, we do not confirm the active Be status of 7 stars claimed as Be stars in the past.  As our results show, Be stars with weaker emission are more difficult to identify conclusively with the color-color diagram method.  Even if the errors in the photometry were negligible, these stars may be confused with unreddened foreground stars or supergiants in a color-color diagram.  In some cases it may be possible to eliminate such foreground stars from the sample based on their positions relative to the MS in a color-magnitude diagram; alternatively, proper motions or radial velocities may be used to separate non-members from the cluster sample.  We do not think that B-type supergiants are a significant source of contamination since supergiants within the cluster will be rejected by the $y_{Bmin}$ constraint in the color-magnitude diagram, and the chances of having a background B supergiant in the sample are very small unless two distant clusters fall along the same line of sight (this is the case with some of the clusters in our forthcoming paper).  However, B-type supergiants can be removed from the sample by including IR photometry in the observational program since Be star disks will create an IR excess that is correlated to their luminosity in H$\\alpha$ \\citep{stee2001}.  If such measurements are not available, H$\\alpha$ spectroscopy may be required to conclusively identify weak H$\\alpha$ emitters.  As Figure \\ref{colormag} shows, our photometry of NGC 3766 does not cover the entire range of B star magnitudes.  Unless the observational program is designed to reach A-type stars, the total number of B, and perhaps Be, stars will be underestimated in any search for Be stars.  This is especially a concern when measuring the fraction of Be stars relative to normal B-type stars because the number of B-type stars increases dramatically with later spectral types.  \\citet{keller2001} illustrate this problem with their deep photometric survey of the clusters $h$ and $\\chi$ Persei.  They find that more than 30\\% of the brightest B stars are Be stars, but the Be star fraction decreases to 2\\% or less for B5--B9 spectral types.  Because many other studies of Be stars (e.g. \\citealt{mermilliod1982, shobbrook1985, grebel1997, fabregat2000}) rely upon magnitude limited samples that do not accurately reflect the true Be fractions, this source of incompleteness is worth discussion.  However, in this situation a more accurate Be/B star fraction can be determined by assuming a typical IMF distribution for the cluster and extrapolating to estimate the number of late stars omitted.  As \\citet{keller2001} show, the contribution of late-type Be stars is small and may be neglected in such an extrapolation.  Both \\citet{abt1987} and \\citet{grebel1997} have discussed another problem with short-term photometric studies of Be stars, namely that they only detect those Be stars with active emission at the time of the observations, and thus only a lower limit of the true Be population is found.  This effect can be significant in observational surveys; \\citet{cote1993} found that 14\\% of the known Be stars in their sample of B stars from the Bright Star Catalogue did not show active emission.  From our own spectra of NGC 3766, we found that only 9 of 21 Be stars discussed in the literature had active disks in 2003 March, although we believe that at least some of these Be stars have been misidentified in the past.  From these two studies, the true fractions of Be stars in a sample could be anywhere between $1.2-2.3$ times the number detected.  Only a long term, multi-year investigation will approach the complete population of Be stars in a given sample.  While our $(b-y, y-\\rm H\\alpha)$ color-color technique is limited in its ability to detect weakly emitting Be stars, this is still an efficient method for identifying Be stars in open clusters.  Applied consistently to a large sample of open clusters, the variability of the Be phenomenon can be neglected and this technique can provide important statistical results about these stars, especially regarding their evolutionary status and the cause of their rapid rotation.  In an upcoming paper, we will apply our photometric method to more than 50 additional clusters and discuss the results."
    },
    "0410/astro-ph0410191_arXiv.txt": {
        "abstract": "We present \\textit{Infrared Space Observatory} (ISO) observations of several OH, CH and H$_2$O rotational lines toward the bright infrared galaxies NGC~253 and NGC~1068. As found in the Galactic clouds in Sgr~B2 and Orion, the extragalactic far--IR OH lines change from absorption to emission depending on the  physical conditions and distribution of gas and dust along the line of sight. As a result, most of the OH rotational lines that appear in absorption toward NGC~253 are observed in emission toward NGC~1068. We  show that  the far--IR spectrum of OH can be used as a powerful diagnostic to derive the physical conditions of extragalactic neutral gas. In particular, we find  that a warm (T$_k$$\\sim$150~K, $n$(H$_2$)$\\lesssim$5$\\times$10$^4$~cm$^{-3}$) component of molecular gas with an OH abundance of $\\simeq$10$^{-7}$ from the inner $\\lesssim$15$''$ can qualitatively reproduce the OH lines toward NGC~253. Similar temperatures but higher densities ($\\sim$5$\\times$10$^5$~cm$^{-3}$) are required to explain the OH emission in NGC~1068. ",
        "introduction": "Galactic nuclei play a key role for our understanding of galactic evolution. The use of molecular diagnostics to extract their physical conditions helps to understand  the nature of the prevailing scenarios, e.g., starbursts vs. active galactic nuclei (AGN). Space observations at mid-- and far--IR wavelengths offer an unique opportunity to observe the pure rotational lines of light hydrides, the simplest chemical building blocks. In particular, H$_2$ observations have revealed the presence of a considerable amount of \\textit{``warm''} (excitation temperatures, T$_{ex}$, around 100-200~K) neutral gas in the Galactic Center (GC; Rodr\\'{\\i}guez--Fernandez et al. 2001) and also in galaxies (Rigopoulou et al. 2002).  \\textit{ISO} has given the opportunity to observe polar species that are  sensitive to a broad range of densities and kinetic temperatures. The far--IR spectrum, and thus the molecular content, of molecular clouds such as Sgr~B2 in the GC (Goicoechea et al. 2004) is surprisingly similar to that observed in ultra-luminous IR galaxies such as Arp~220 (Gonz\\'alez--Alfonso et al. 2004). One of the  species with the largest abundance is the OH radical.  Bright IR galaxies containing OH usually show megamaser emission at 18~cm. These lines arise from the hyperfine splitting of the OH $^2\\Pi_{3/2}$ $J$=3/2 ground rotational state. Among the possible excitation mechanisms of these masers; near-- and far--IR pumping, collisions and UV radiation, the far--IR excitation seems to be the dominant mechanism inverting the population of the hyperfine levels (e.g. Skinner et al. 1997). Furthermore, the relative population of the OH  levels is governed by the far--IR pure rotational lines and detailed treatment of the rotational excitation is required to analyze the OH radio lines at 18~cm but also the $^2\\Pi_{3/2}$ $J$=5/2 and $^2\\Pi_{1/2}$ $J$=1/2 lines at 5 and 6~cm (Randell et al. 1995).   The OH ground--state rotational line at $\\sim$119~$\\mu$m was first detected in absorption toward Sgr~B2 (Storey et al. 1981), while several excited rotational lines have been analyzed by Goicoechea \\& Cernicharo (2002). A detailed analysis of the far--IR  spectrum of OH has been also carried out in the Orion region (Watson et al. 1985; Vicuso et al. 1985; Melnick et al. 1987, 1990; Betz \\& Boreiko 1989). The observed lines appear in emission and some of them even show more complex P--Cygni profiles. Different excitation conditions determine the distinct OH spectral signatures observed toward Orion and Sgr~B2.  These studies have shown that OH can be a powerful prove of the warm neutral gas in molecular clouds. Scaled to extragalactic studies, OH can thus provide significant information about star forming regions and the heating processes in the inner arcsecs of a galaxy, in addition to the understanding of the megamaser excitation mechanism. Moreover, OH is a key molecule in the gas--phase chemistry networks. The presence of OH in the warmer regions may enhance and even dominate the formation of other O--bearing species, so that the determination of its abundance is also  relevant to understand the prevailing chemistry.  In this letter we present and analyze several OH far--IR lines, involving levels  up to $\\sim$420~K, toward the Starburst galaxy NGC~253 and the Seyfert 2 galaxy NGC~1068. We compare these  extragalactic OH lines with our far--IR observations of Sgr~B2(M) and Orion~IRc2, and run simple non--LTE radiative transfer models to qualitatively understand the behavior of the detected lines.   \\begin{figure*} [ht] % \\centering \\includegraphics[angle=-90,width=16.5cm]{f1.eps} \\caption{ISO/LWS grating observations of NGC~253 and NGC~1068. The ordinate corresponds to the flux and the abscissa is the rest wavelength in $\\mu$m. Atomic fine structure lines are labeled. Zooms around selected wavelengths are also shown in the smaller panels with models of the ground--state rotational lines of OH, CH and H$_2$O (see text). The ordinate of these panels is the continuum normalized flux.} \\label{gratings} \\end{figure*}% ",
        "conclusions": "Irrespective of the specific nuclei scenario and of the density of the dominant molecular clouds, OH observations show that a similar component of warm neutral gas (T$_k$$\\sim$150~K) exists in NGC~253 and NGC~1068. This component can influence the chemistry as neutral--neutral reactions involving OH start to be efficient for the production of other oxygen molecules such as H$_2$O and SiO. The contribution of PDRs, shocks, X--rays and cosmic rays to the heating of such component in Seyfert vs. Starburst galaxies is still far from settled. Future works using diagnostics of the warm gas will be extremely useful. In particular, several rotational lines of light hydrides such as OH, CH and H$_2$O will be observed at larger angular resolution by the HIFI and PACS instruments, on board Herschel in many galactic and extragalactic sources."
    },
    "0410/astro-ph0410472_arXiv.txt": {
        "abstract": "We examine outstanding issues in the analysis and interpretation of the halo Li plateau. We show that the majority of very Li-poor halo Li-plateau stars (5 out of 8) have high projected rotation velocities $v$sin$i$ between 4.7 and 10.4~km~s$^{-1}$. Such stars have very different evolutionary histories to Li-normal plateau stars, and hence cannot be included in studies of Li depletion by normal halo dwarfs. Uncertainties in the effective temperature scale for metal-poor stars continue to challenge the analysis of Li. ",
        "introduction": "From observations of 11 main-sequence stars belonging to the Galactic halo, Spite \\& Spite \\cite{ryan:ss82} concluded that the lithium abundance was essentially independent of metallicity for halo stars hotter than 5600~K, and inferred that the Li abundance was ``hardly altered'' from the Big Bang. Two decades of work has followed, increasing the number of stars observed and the range of metallicity that they span, in an effort to establish the primordial Li abundance more securely.  The primordial Li abundance was sought primarily because of its ability to constrain the baryon to photon ratio in the Universe, or equivalently the baryon contribution to the critical density. In this way, Li was able to complement estimates from $^4$He, the primordial abundance of which varied only slightly with baryon density. Li also made up for the fact that the other primordial isotopes, $^2$H (i.e. D) and  $^3$He, were at that time difficult to observe and/or interpret. During the late 1990's, however, measurements of D in damped Lyman alpha systems (high column-density gas believed to be related to galaxy discs) provided more reliable constraints on the baryon density than Li could do (e.g. \\cite{ryan:otksplw01}). Even more recently, the baryon density has been inferred from the angular power spectrum of the cosmic microwave background radiation,  for example from the WMAP measurements \\cite{ryan:s03}. We consider the role of Li plateau observations post WMAP. ",
        "conclusions": ""
    },
    "0410/astro-ph0410644_arXiv.txt": {
        "abstract": "We propose a novel multivariate Monte Carlo method as an efficient and flexible approach to analyzing extended X-ray sources with the Reflection Grating Spectrometer (RGS) on XMM Newton.  A multi-dimensional interpolation method is used to efficiently calculate the response function for the RGS in conjunction with an arbitrary spatially-varying spectral model. Several methods of event comparison that effectively compare the multivariate RGS data are discussed.  The use of a multi-dimensional instrument Monte Carlo also creates many opportunities for the use of complex astrophysical Monte Carlo calculations in diffuse X-ray spectroscopy.  The methods presented here could be generalized to other X-ray instruments as well. ",
        "introduction": "Diffuse sources, such as clusters of galaxies, supernova remnants, and the interstellar hot haloes of elliptical galaxies, comprise some of the most interesting targets for astrophysical X-ray spectroscopy.  Photons from these sources are focused and dispersed by optics and recorded in detectors designed to measure three interesting quantities:  two related to the position on the detector ($x$, $y$) and one related to an intrinsic photon energy measurement ($p$). These quantities are indirectly related to the sky coordinates ($\\theta, \\phi$) and the energy of the incident photon ($e$) that can be predicted from an astrophysical model with parameters, \\{$\\eta_1$, $\\eta_2$, $\\eta_3$, \\ldots\\} $\\equiv$ $\\eta_i$.  The fundamental goal of data analysis for diffuse X-ray spectroscopy is to calculate the probability of a given distribution of photons in the ($x$, $y$, $p$) data-space given the parameters of a particular astrophysical model.  Prior to the launch of the {\\it Chandra} (\\citealt{weisskopf}) and {\\it XMM-Newton} (\\citealt{jansen}) observatories, available instruments were characterized by sufficiently poor spatial and/or spectral resolution that simple one-dimensional spectral fitting techniques applied to data extracted from the image were sufficient for most applications.  However, the grating experiments and non-dispersive CCD experiments on these two new observatories both have significant imaging and spectral capabilities, making it warranted to develop new techniques that utilize the full dimensionality of the data. In particular, the Reflection Grating Spectrometers (RGS) on XMM-Newton have a number of unique characteristics that make them powerful X-ray spectrometers for arcminute-size X-ray sources.  In this paper, we demonstrate that the complex nature of the response function for the RGS requires new analysis methods that utilize the full dimensionality of the data.  We also demonstrate how these methods might be useful for other X-ray instruments.  The techniques for analyzing an X-ray spectrum of an unresolved source are well-developed and have been essentially unchanged since the analysis of early X-ray spectra (\\citealt{gorenstein}).  For that case, there is only one interesting measured quantity that is related to the energy of the photon.  For non-dispersive spectrometers (e.g. CCDs, calorimeters, and proportional counters), the photon energy measurement ($p$) is the useful quantity.  For dispersive spectrometers (e.g. gratings and crystals), the dispersion coordinate, $x$, or the position along the detector parallel to the dispersion axis, is the useful measured value.  In either case, the detection probability, $D$, of finding a photon at a given $x$ or $p$ is then given by \\begin{eqnarray} D(x~\\mbox{or}~p~\\vert~\\eta_i ) = \\int de~R(x~\\mbox{or}~p ~\\vert~ e)~S(e~\\vert~\\eta_i ) \\end{eqnarray} \\noindent which is an integral over all energies, $e$.  Here, R is the response kernel and S is the input spectral model, which is a function of both the energy and the input parameters.  The response kernel describes the probability of obtaining a distribution of measured values given a photon with incident energy, $e$.  The integral is computed numerically by converting the expression into a sum of discrete energies in well-developed software packages, such as XSPEC (\\citealt{arnaud}) or SPEX (\\citealt{kaastraspex}).  For an extended source, the probability distribution is described by a three-dimensional function and its calculation requires an integration over both the sky coordinates ($\\theta,~\\phi$) and the intrinsic energy ($e$).  The three-dimensional detection probability, D, is then given by \\begin{equation} D(x,y,p~\\vert~\\eta_i ) = \\int de~d\\theta~d\\phi~R(x,y,p ~\\vert~ \\theta,\\phi,e)~S(\\theta,\\phi,e~\\vert~\\eta_i ) \\end{equation} \\noindent Computing this entire integral directly is often impractical.  Nevertheless, some approximations can be made that are useful in restricted situations.  In particular, one can assume: 1) the source spectrum is independent of spatial position, and 2) the response does not vary as function of the off-axis angle.  The problem can then be reduced to a one-dimensional integral, as for a point source.  The former assumption is sometimes justified, but has become untenable in many recent analyses. For example, modeling an X-ray cluster of galaxies as a thermal plasma whose temperature varies spatially already violates this assumption. The second assumption depends on the nature of the X-ray instrument as well as the angular scale of the problem being investigated.  Attempts to circumvent this problem by weighting the response matrix for Chandra ACIS-S observations can be found in \\cite{markevitch} and in \\cite{arnaud2} for XMM-Newton EPIC data. We demonstrate in \\S2 that the complex nature of the response function, $R$, for the Reflection Grating Spectrometer makes this approximation impossible to implement without assuming that there is no spectral variation over the entire field of view. Instead, we are forced to consider new methods that consider the full multi-dimensionality of Equation 2.  Equation 2 can be evaluated directly through a Monte Carlo calculation as we will demonstrate in \\S3. It has been known for some time that Monte Carlo methods are efficient for the evaluation of multi-dimensional integrals (\\citealt{metropolis}).  Few photons are typically detected in X-ray observations, so the exact calculation of the full integral is superfluous. After constructing an efficient Monte Carlo algorithm, we outline a generic approach for the analysis of diffuse X-ray sources with Monte Carlo calculations in \\S4.  In \\S5, we discuss several methods of event comparison that can be used with the photons simulated with the RGS Monte Carlo and real data.  In \\S6, we discuss several straightforward extensions to these methods. ",
        "conclusions": ""
    },
    "0410/gr-qc0410023_arXiv.txt": {
        "abstract": "We have developed a theoretical model and a numerical code for stationary rotating superfluid neutron stars in full general relativity. The underlying two-fluid model is based on Carter's covariant multi-fluid hydrodynamic formalism. The two fluids, representing the superfluid neutrons on one hand, and the protons and electrons on the other, are restricted to uniform rotation around a common axis, but are allowed to have different rotation rates. We have performed extensive tests of the numerical code, including quantitative comparisons to previous approximative results for these models. The results presented here are the first ``exact'' calculations of such models in the sense that no approximations (other than that inherent in a discretized numerical treatment) are used. Using this code we reconfirm the existence of prolate-oblate shaped configurations. We studied the dependency of the Kepler rotation limit and of the mass-density relation on the relative rotation rate. We further demonstrate how one can simulate a (albeit fluid) neutron-star ``crust'' by letting one fluid extend further outwards than the other, which results in interesting cases where the Kepler limit is actually determined by the outermost but \\emph{slower} fluid. ",
        "introduction": "The aim of this work is to calculate fully relativistic stationary models of superfluid neutron stars including all non-dissipative couplings between the two fluids induced by the equation of state (EOS), in particular the entrainment effect. In addition to studying the stationary properties of relativistic superfluid neutron stars, these models can serve as the unperturbed initial state in a dynamical study of neutron star oscillations, neutron star collapse to a black hole, or as a starting point in studying pulsar glitch-models.  Neutron stars are fascinating astrophysical objects: on one hand they represent a formidable ``laboratory'' of fundamental physics, as the composition and equation of state of their inner core still lies beyond the reach of experimental and theoretical physics. On the other hand, the advent of increasingly sensitive gravitational wave detectors promises to open a new observational window on neutron stars, which will allow us to gain new insights into these still rather poorly understood objects. Gravitational wave astronomy could represent the first opportunity to observe neutron star oscillations, providing a new view on their inner dynamics. Considering the success of classical terrestrial seismology and astero-seismology of the sun and of main-sequence stars, one could expect this to result in substantial progress in our understanding of the dynamics and composition of neutron stars.  Additionally, observing quasi-permanent quadrupolar deformations (``mountains'') on neutron stars via gravitational waves\\footnote{Note that this search has already begun, see \\cite{lsc:_settin_upper_limit} for a discussion and first results} will give valuable complementary information about their rotational behavior, which is currently only observable via their electromagnetic pulses.  Most theoretical studies of neutron star dynamics have relied on rather simplistic single-fluid models. In this work we attempt a more realistic description of neutron stars by taking their superfluidity into account via the use of a two-fluid model. Neutrons and protons in neutron stars are predicted to be superfluid (e.g.~see~\\cite{baldo92:_superfl_neutr_star_matter,sjoberg76:_effect_mass}), and this feature forms a fundamental ingredient in the current (albeit rudimentary) understanding of the glitch phenomenon observed in pulsars (e.g.~see~\\cite{link00:_probing_NS_glitches,carter00:_centr_buoyancy,acp03:_twostream_prl}). Due to the superfluidity and therefore lack of viscosity of the neutrons in the crust and in the outer core, they can flow freely through the other components. The remaining constituents (i.e. crust-nuclei, electrons, muons and protons) are assumed to be ``locked'' together on short timescales by viscosity and the magnetic field. Thereby they form another fluid, which in the following will be referred to as ``protons'' for simplicity. These assumptions characterize the so-called two-fluid model of neutron stars. These two fluids are strongly coupled by the strong nuclear force acting between protons and neutrons, and therefore a hydrodynamic two-fluid framework incorporating these couplings is required for their description. This framework will be presented in the next section. Recently it was pointed out that such a two-fluid system can be subject to a two-stream instability if the relative velocity of the two fluids exceeds a critical velocity \\citep{acp04:_twostream}. This could therefore be relevant in neutron stars and might be related to the glitch phenomenon\\citep{acp03:_twostream_prl}, which provides another motivation for studying the properties of such two-fluid systems.  In this paper we study the stationary structure of such two-fluid models, in which the two fluids are restricted to uniform rotation around a common axis, but allowing for two different rotation rates. This neutron star model was first studied quantitatively by \\citet{prix99:_slowl_rotat_two_fl_ns} in the Newtonian context using a generalized Chandrasekhar-Milne slow-rotation approximation, and neglecting the direct interactions between the two fluids. \\citet{andersson01:_slowl_rotat_GR_superfl_NS} used Hartle's variant of the slow-rotation approximation to study this model in general relativity. \\citet{prix02:_slow_rot_ns_entrain} further extended the Newtonian study to fully include all (non-dissipative) couplings via entrainment and the nuclear ``symmetry-energy'', and they found an analytic solution for a subclass of two-fluid equations of state (which generalizes the $P\\propto \\rho^2$-type polytropes). More recently, \\citet{yoshida04:_rapid_newton} have devised an alternative approach in the Newtonian case, by treating only the relative rotation between the two fluids as small, while allowing for fast rotation of the neutron star as a whole. Furthermore, \\citet{comer04:_slowl} has recently used the relativistic slow-rotation approximation to study the properties of the first available fully relativistic two-fluid EOS incorporating entrainment, which was derived by \\citet{comer03:_relat_mean_field}.  Here we present a generally relativistic numerical code for solving the full two-fluid model without approximations. A preliminary progress-report on the development of this code, and some early results were presented in \\citep{prix02:_relativistic_sf_ns}.  While our model and code allow in principle for any given two-fluid equation of state (EOS), for the sake of simplicity and a better numerical convergence we restrict ourselves in this paper to the use of a (rather general) class of two-fluid ``polytropes''. This choice is also motivated by the lack of a useful two-fluid neutron star equation of state in the literature, especially concerning the aspect of entrainment.  Even though \\citet{comer03:_relat_mean_field} have a fully relativistic model that includes entrainment, it has not yet been developed to the point that it will produce a tabular equation of state that could be used in our code. We expect the qualitative features of our model to be well represented by the analytic EOS used in this work.  The plan of this paper is as follows: In section~\\ref{sec:canonical-two-fluid} we introduce the formalism and notation of  covariant two-fluid hydrodynamics. In section~\\ref{sec:stat-axisymm-conf} we discuss the specialization to an axisymmetric and stationary system, and we introduce the $3+1$ framework for Einstein's equations. In section~\\ref{sec:numerical-procedure} we describe the numerical procedure for solving the resulting elliptical system of equations. The tests performed on the numerical code are discussed in section~\\ref{sec:tests-numerical-code}, and our numerical results are presented in section~\\ref{sec:numerical-results}. A discussion of this work is given in section~\\ref{sec:summary}. In appendix~\\ref{sec:newt-analyt-slow} we derive a new analytic Newtonian slow-rotation solution, which is used for comparison to our numerical results. ",
        "conclusions": "\\label{sec:summary}  We have developed a theoretical framework and a numerical code for computing stationary, fully relativistic superfluid neutron star models.   Using this code we have reconfirmed the existence of oblate-prolate shaped two-fluid configurations, previously shown in \\cite{prix02:_slow_rot_ns_entrain,prix02:_relativistic_sf_ns}. We have studied the dependency of the Kepler rate of a two-fluid star on the relative rotation rate $\\R$. We have compared this to the Kepler-rate predicted by a Newtonian slow-rotation approximation, which is found to typically overestimate the Kepler-rate by about $15\\%$, but which can also sometimes underestimate it, as seen in the case of EOS~I and $\\R\\lesssim -0.25$ (cf. Fig.~\\ref{fig:Kepler}).  The relative rotation rate can also have a large influence (at fixed central density) on the mass-density relation, as shown in Fig.~\\ref{fig:MassDens}.  Another interesting aspect of this model is that we are not restricted to configurations in chemical equilibrium at the center. Choosing the central chemical potentials to be different allows one to emulate a neutron star ``crust'' (albeit a fluid one), as one fluid will now extend further outwards than the other, as seen in Fig.~\\ref{fig:msPulsar} and Fig.~\\ref{fig:Kepler_II}. One interesting observation from such configurations is that the Kepler-limit will be determined by the outer fluid (forming a cusp), while this can actually be rotating \\emph{slower} then the inner fluid.  We currently use a (quite general) EOS class of two-fluid polytropes, but this can be extended straightforwardly to more ``realistic'' nuclear-physics equations of state. In particular it might be interesting in the next step to use the first relativistic two-fluid EOS incorporating entrainment by \\citet{comer03:_relat_mean_field}. Furthermore it would be important to add the presence of a solid crust and to allow for differential rotation in the superfluid neutrons (differential rotation in single-fluid stars has been implemented and used in LORENE already, cf. \\cite{villain04:_evolut,goussard98:_rapid}).  The astrophysically most interesting future extension of this work would probably consist in studying the oscillation modes of such models, which would be directly related to the emission of gravitational waves. In these non-stationary situations, however, dissipative mechanisms like viscosity and mutual friction would also start to play a role and should be included in the model.  \\appendix"
    },
    "0410/astro-ph0410587_arXiv.txt": {
        "abstract": "{ We determine the Strouhal number (hereafter St), which is essentially a nondimensional measure of the correlation time, from numerical calculations of convection. We use two independent methods to estimate St. Firstly, we apply the minimal tau-approximation (MTA) on the equation of the time derivative of the Reynolds stress. A relaxation time is obtained from which St can be estimated by normalising with a typical turnover time. Secondly, we calculate the correlation and turnover times separately, the former from the autocorrelation of velocity and the latter by following test particles embedded in the flow. We find that the Strouhal number is in general of the order of 0.1 to 1, i.e. rather large in comparison to the typical assumption in the mean-field theories that ${\\rm St} \\ll 1$. However, there is a clear decreasing trend as function of the Rayleigh number and increasing rotation. Furthermore, for the present range of parameters the decrease of St does not show signs of saturation, indicating that in stellar convection zones, where the Rayleigh numbers are much larger, the Strouhal number may indeed be significantly smaller. ",
        "introduction": "The mean-field theories of hydromagnetic dynamos and angular momentum transport need knowledge of turbulent correlations, namely the electromotive force and Reynolds stresses, respectively. The explicit calculation of these is virtually impossible for any astrophysical conditions due to the lack of necessary computational cababilities. Therefore the small-scale effects are usually parametrised by transport coefficients which relate the turbulent correlations to the mean quantities (e.g. Steenbeck, Krause \\& R\\\"adler \\cite{SteenKrauRaed1966}; R\\\"udiger \\cite{Ruediger1989}). However, in order to calculate the transport coefficients, knowledge of the small-scale velocities and magnetic fields are still needed. Thus, simplifying assumptions such as the first order smoothing approximation (hereafter FOSA) have been used (Steenbeck et al. \\cite{SteenKrauRaed1966}). Although the results obtained with FOSA are in agreement with observations in many cases, the basic assumptions of the theory have rarely been thoroughly studied (however, see Petrovay \\& Zsarg$\\acute{\\rm o}$ \\cite{PetrZsar1998}; K\\\"apyl\\\"a et al. \\cite{Kaepylaeea2004b}).  The validity condition for FOSA is that either the Reynolds or the Strouhal number is small. The former condition is clearly never met, but of the latter in the context of convection and dynamo theory no systematic study seems to exist. A further notion associated with the Strouhal number is that even though if St is not small in the sense that ${\\rm St} \\ll 1$, it is still possible to calculate the transport coefficients from a cumulative series expansion if St $<$ 1 (e.g. Knobloch \\cite{Knobloch1978}; Nicklaus \\& Stix \\cite{NickStix1988}).  Recently it has been shown numerically that St can \\emph{exceed} unity for forced turbulence (Brandenburg, K\\\"apyl\\\"a \\& Mohammed \\cite{Brandea04}; Brandenburg \\& Subramanian \\cite{BrandSubr04}) in the contexts of passive scalar diffusion and mean-field dynamos, respectively. In these studies, instead of FOSA, the minimal tau-approximation (hereafter MTA) is found to be in better agreement with the calculations. In MTA, instead of the turbulent correlation itself, the time derivative is investigated, and the higher order terms are parameterised by a term which is just the original turbulent correlation divided by a relaxation time (Blackman \\& Field \\cite{BlackField2003}). Furthermore, interpreting the relaxation time as the correlation time of the turbulence the Strouhal number can be estimated.  The forced turbulence results raise the question of the value of St for convection, which seems to be rather badly known. For example, solar surface observations indicate that the lifetime and turnover time of granules is approximately the same, yielding ${\\rm St} \\approx 1$ (e.g. Stix \\cite{Stix2002}). However, this result may not be relevant for the solar dynamo which is working in the deeper layers. In the present study we estimate the Strouhal number for numerical convection by two independent methods. Firstly we apply the MTA on the equation of the Reynolds stresses, and secondly we calculate the correlation and turnover times separately directly from the flow. The correlation time is estimated from the autocorrelation of velocity and the turnover time by following test particles embedded in the flow. The computational model is a Cartesian box situated at a latitude $\\Theta$ on a star. The model is described in detail in K\\\"apyl\\\"a, Korpi \\& Tuominen (\\cite{Kaepylaeea2004a}).  The remainder of the paper is organised as follows: in Sect.~(\\ref{sec:strouhal}) the two methods used to estimate the Strouhal number are discussed. Sects.~(\\ref{sec:results}) and (\\ref{sec:conclusions}) give the results and conclusions, respectively. ",
        "conclusions": "\\label{sec:conclusions} We estimate the Strouhal number from numerical models of convection with two independent methods. Firstly, we apply the minimal tau-approximation of the equation of the Reynolds stress. Secondly, we calculate the correlation and turnover times directly from the flow without any recourse to a mean-field theory.  We find that the Strouhal number from the MTA reaches values of maximally $\\approx 0.5$. St has a maximum for intermediate rotation, where the Reynolds stress itself also peaks. The Strouhal number follows closely the same trend as the vertical $\\Lambda$-effect (see e.g. K\\\"apyl\\\"a et al. \\cite{Kaepylaeea2004a}) as function of rotation due to the similar functional form.  The correlation time is seen to decrease consistent with power law $\\tau_{\\rm c} \\propto {\\rm Ra}^{-0.45}$ as function of the Rayleigh number and approximately with $\\tau_{\\rm c} \\propto {\\rm Co}^{-0.45}$ for moderate and rapid rotation. These results indicate that although the values of St in this study are generally of the order of 0.1 to 1, the Strouhal number in stellar convection zones, where Ra is most certainly, and Co at least probably, much larger than in the present study, may be significantly smaller."
    },
    "0410/astro-ph0410064_arXiv.txt": {
        "abstract": "We investigate the possibility to detect giant planets that are still embedded in young circumstellar disks. Based on models with different stellar, planetary, and disk masses, and different radial positions of the planet we analyze the resulting submillimeter appearance of these systems. We find that the influence of the planet on the spectral energy distribution could not be distinguished from that of other disk parameters. However, dust reemission {\\em images} of the disks show that the hot region in the proximity of a young planet, along with the gap, could indeed be detected and mapped with the Atacama Large Millimeter Array in the case of nearby circumstellar disks (d$<$100\\,pc) in approximate face-on orientation. ",
        "introduction": "Planets are expected to form in circumstellar disks, which are considered as the natural outcome of the protostellar evolution, at least in the case of low and medium mass stars (e.g., Adams, Lada, \\& Shu~1987; Lissauer~1993). While a detailed picture of the evolution of the circumstellar environment, in particular of circumstellar disks, has been developed already, the planet formation process is mostly still under discussion. There exist two main different scenarios for planet formation. The first is characterized by three major phases: (1)~dust grain growth from submicron-sized particles to centimeter / decimeter-sized bodies via coagulation is followed by (2)~an agglomeration process that leads to the formation of (sub)kilometer-sized planetesimals, which (3)~form terrestrial (rocky) and Uranian (icy) planets or the cores of Jovian (gaseous) planets by further accretion of solid material (e.g., Pollack et al. 1996, Weidenschilling 1997). Alternatively, planetesimals may form via the gravitational instability of solids that have settled to the midplane of a circumstellar disk (Goldreich \\& Ward~1973, Youdin \\& Shu~2002). Adequate constraints from observations are required in order to either verify or rule out existing hypotheses about these planet formation scenarios.  During recent years, numerical simulations studying planet-disk interactions have shown that planets may cause characteristic large-scale signatures in the disk density distributions. The most important of these signatures are gaps and spiral density waves in young circumstellar disks (e.g., Bryden et al.~1999; Kley~1999; Lubow, Seibert, \\& Artymowicz~1999; Kley, D'Angelo, \\& Henning~2001; Bate et al.~2003; Winters, Balbus, \\& Hawley~2003; Nelson \\& Papaloizou~2003) and resonance structures in evolved systems, so-called Debris disks (e.g., Liou \\& Zook~1999, Ozernoy et al.~2000, Moro-Mart\\'{\\i}n \\& Malhotra~2002). The importance of investigating these signatures lies in the possibility that they can be used to search for embedded (i.e., young) planets. Therefore, disk features can provide constraints on the processes and timescales of planet formation. Indeed, several debris disks around nearby main-sequence stars show structures and asymmetries that are considered to result from planetary perturbations (Holland et al.~1998, 2003; Schneider et al.~1999; Koerner, Sargent, \\& Ostroff~2001).  In our present study we consider young disks with a density structure that is dominated by gas dynamics. While it was shown before, that the planned Atacama Large Millimeter Array (ALMA) will be able to map the gap caused by a massive planet in such disks (Wolf et al.~2002), we are now using more detailed simulations in order to investigate whether the planet itself and/or its surrounding environment could be detected. The detection of a gap would already represent a strong indication of the existence of a planet, thus giving information such as planetary mass, viscosity, and pressure scale-height of the disk. The detection / non-detection of warm dust close to the planet, however, would additionally provide valuable indications on the temperature and luminosity of the planet and on the density structure of the surrounding medium.  We base our investigation on hydrodynamical simulations of circumstellar disks with an embedded planet and subsequent radiative transfer simulations, with the aim of deriving observable quantities of these systems. Technical details on the performed simulations together with an explanation and justification of the chosen model setup are given in \\S~\\ref{basics}. The results are discussed in \\S~\\ref{results} and our conclusions are presented in \\S~\\ref{summary}. ",
        "conclusions": "\\label{summary}  We performed simulations with the goal to answer the question of whether the circumplanetary environment of a giant planet, embedded in a young circumstellar disk, can be detected. Indeed, we could show that a detection will be possible via mapping with the Atacama Large Millimeter Array, even if under particular circumstances. We found two major constraints that limit the applicability of this approach to study planet formation. First, only for nearby objects at distances of typically less than $\\approx 100$\\,pc the spatial resolution provided by ALMA will be high enough to allow to spatially separate the circumstellar disk and circumplanetary material from one another. Second, the emission from the circumplanetary environment is significantly strong only in the case of massive and thus young circumstellar disks. Furthermore, the signal of the radiation from the planet / circumplanetary material will be best distinguishable if a high planet-to-star mass ratio is targeted. We want to remark that in contrast to the most successful planet detection technique so far, based on radial velocity measurements (e.g., Mayor \\& Queloz~1995, Marcy \\& Butler~2000), the likelihood of detection does not increase with decreasing distance of the planet to the star. In particular, in none of the model configurations with a planet at a distance of only 1\\,AU from the central star (as opposed to 5\\,AU in models with positive detections) the circumplanetary environment could be distinguished from the circumstellar disk.  As our study shows, high-resolution (sub)millimeter observations will provide a valuable tool in the near future to challenge planet formation theories. For this reason it is worth to continue this investigation with the goal to achieve more detailed predictions. There are three major issues which could be significantly improved: First, the simulation of the radiative transfer (or the energy transfer in general) should be coupled with or embedded in the hydrodynamical calculations. According to the existing computational resources this goal can be achieved in the near future (see, e.g., D'Angelo et al.~2003a). Second, the hydrodynamical simulations that allow to gain a higher resolution in the region of the planet should be employed (e.g., D'Angelo et al.~2003b; Bate et al.~2003). In this manner the dust reemission could be calculated more accurately. Third, the emission spectrum of the protoplanet should be modeled more accurately: Internal heating and heating due to accretion onto the planet have to be investigated separately (taking into account the anisotropy of different radiative sources). Furthermore, objects with temperatures $\\sim 10^3$\\,K are expected to have strong absorption bands from molecular transitions and even dust absorption solid state features in their spectra (e.g., Allard et al.~2001), which should be considered in the heating process of the local planetary environment."
    },
    "0410/astro-ph0410122_arXiv.txt": {
        "abstract": "The study of sterile neutrinos has recently acquired a different flavor: being now excluded as the dominant solution for the solar or atmospheric conversions, sterile neutrinos, still attractive for many other reasons, have thus become even more elusive. The present relevant questions are: which subdominant role can they have? Where (and how) can they show up? Cosmology and supernov\\ae\\ turn out to be powerful tools to address these issues. With the most general mixing scenarios in mind, I present the analysis of many possible effects on BBN, CMB, LSS, and in SN physics due to sterile neutrinos. I discuss the computational techniques, present the state-of-the-art bounds, identify the still allowed regions and study some of the most promising future probes. I show how the region of the LSND sterile neutrino is excluded by the constraints of standard cosmology. ",
        "introduction": "The study of sterile neutrinos (namely: additional light fermionic particles that are neutral under all Standard Model gauge forces, but can be a non-negligible ingredient of our world through their mixing with $e, \\mu, \\tau$ neutrinos) has recently acquired a different flavor: the established solar and atmospheric anomalies seem produced by oscillations among the three active neutrinos~\\cite{sun nus, atmo nus} so that their explanation in terms of oscillation into a $\\nus$ state (which was viable and fairly popular up to a few years ago) is believed to be now ruled out as the dominant mechanism. This means that the {\\bf relevant questions concerning sterile neutrinos nowadays} have become: $\\bullet$which is the {\\it subdominant} role still possible for sterile neutrinos in solar and atmospheric neutrinos? $\\bullet$where can we detect the effects of a sterile neutrino? i.e. which are the most sensitive experiments (in astrophysics, cosmology or man-made set-ups) in which sterile neutrinos can be discovered? $\\bullet$how can we detect the effects of a sterile neutrino? i.e. which are the signatures of its presence?  To answer these questions, the investigation on sterile neutrinos requires a more extensive and deep approach.\\\\ Indeed, for instance, most of the previous analysis used to consider only the peculiar oscillation pattern that gives the simplest physics: the initial active neutrino $|\\nu_a \\rangle$ ($\\nue$ in the case of solar neutrinos, $\\num$ in the case of the atmospheric ones) oscillates into an energy-independent mixed neutrino $\\cos\\theta_s |\\nu'_a \\rangle + \\sin\\theta_s |\\nu_s \\rangle$. This of course leaves unexplored the largest part of the parameter space. Sometimes, moreover, previous analysis used to neglect for simplicity the mixing among active neutrinos (this is the case of most studies on the sterile effects in cosmology or in supernov\\ae). Such a mixing is now established and its parameters are reasonably pinned down.   These considerations motivate the analysis performed in~\\cite{CMSV} and presented here, which  considers and includes: \\begin{itemize} \\item[{\\bf a.}] any possible $\\nu_{e,\\mu,\\tau}-\\nus$ mixing pattern; \\item[{\\bf b.}] the established $\\nue-\\nu_{\\mu,\\tau}$, $\\num-\\nut$ mixings; \\item[{\\bf c.}] all possible neutrino sources and contexts (cosmology (BigBang Nucleosynthesis-BBN, Cosmic Microwave Background-CMB, Large Scale Structures-LSS), astrophysics (the Sun, supernov\\ae-SNe...), atmospheric neutrinos, reactor and accelerator experiments...), \\end{itemize} The purpose is to set the state-of-the-art bounds on the active-sterile mixing parameters (in each context separately and then in a combined way) and to identify the most promising future probes. Since the different cosmological quantities and the observables of SN physics turn out to be very important to this aim, investigating complementary regions and allowing the application of techniques that extend to other fields, these are the contexts on which I focus in these Proceedings.  \\comment{ The main {\\bf purposes} of such a study are evident: \\begin{enumerate} \\item we {\\bf look for evidences} of sterile neutrino effects in present data from all contexts.\\\\ I anticipate we do find {\\bf no evidence} with any statistical significance (the LSND result is addressed separately). Therefore... \\item we {\\bf set the present bounds} on the active-sterile mixing parameters in each context separately, and then in a combined way.\\footnote{We do not perform a global fit of all data (cosmological, astrophysical, solar, reactor, atmospheric...) together because of the different reliability of the different probes, the diversity of the assumptions and so on. Not to mention the computational effort that such a procedure would require. The plot with superimposed bounds presented in the end is in our opinion a better way to present the combined analysis.} \\item we {\\bf compare the LSND hint} against those bounds. \\item we {\\bf identify the most promising future probes} of sterile neutrino effects in all contexts. \\end{enumerate} } \\comment{ Before going to the details, a crucial issue that concerns the {\\bf motivations} of such a study is still to be addressed. Now that the need of sterile neutrinos for the solar and atmospheric anomalies is died, can sterile neutrinos still be interesting ingredient of our world at all? The answer is yes, for at least two different sets of reasons.\\\\ } Before going to the details, let's stress that now that the hunger of sterile neutrinos for the solar and atmospheric anomalies is over, nevertheless {\\bf sterile neutrinos are even more attractive} for at least two sets of reasons.\\\\ First, from a top-down point of view, sterile fermions that are naturally light or cleverly lightened arise in many theories that try to figure out what is beyond the Standard Model. To begin, slightly beyond the context of the SM, the right handed neutrino is an obvious candidate, to complete the lepton sector in similarity (symmetry?) with the quark one. In this case, actually, three states (one per family) would be natural. More broadly, several (SuSy-/GUT-/string-/ED- inspired) SM gauge singlets line up awaiting for consideration (axino, branino, dilatino, familino, Goldstino, Majorino, mirror fermion, modulino, radino...)~\\cite{sterile models}. The origin and the load of information of any of these particles can be very different, but from an effective point of view we simply need to parameterize their mixing with the active neutrinos in terms of  mixing angles and $\\Delta m^2$ (see below) to include them all in the analysis. Independently from the specific model, the discovery of a new light particle would be of fundamental importance and deserves to be investigated per se.\\\\ Second, from a more phenomenological bottom-up perspective, sterile neutrinos are repeatedly pointed as a possible explanation for several puzzling situations in particle physics, astronomy and cosmology. For instance, they have been invoked~\\cite{sterile puzzles} to account for the origin of the pulsar kicks, to constitute a Dark Matter candidate, to explain (via their decay) the diffuse ionization of the Milky Way, to help the r-process nucleosynthesis in the environment of exploding stars, to interpret the slightly too low Argon production in the Homestake experiment... and notoriously to explain the LSND claimed evidence of oscillations~\\cite{LSND}. Any one of these puzzles, in general, points to specific sterile neutrinos (i.e. with a specific mixing pattern with the active ones) and it is therefore worthwhile to explore them in an extensive way.  \\comment{ \\medskip  In Sec.\\ref{4nu} I present the non-standard four neutrino formalism that we use. In Sec.\\ref{cosmology} I address the sterile effects in cosmology, comparing the relative sensitivities of two BBN probes (the $^4$He and Deuterium abundances), of CMB and of LSS. In Sec.\\ref{LSND section} I single out the LSND hint for a sterile neutrino and I discuss the issue of its compatibility with the cosmological constraints. In Sec.\\ref{SN} I study sterile oscillations in SN1987A and future supernov\\ae. The study of the sterile effects in all other contexts (solar, atmospheric neutrinos; reactor, accelerator SBL, LBL neutrinos; high energy astrophysical neutrinos) can be found in~\\cite{CMSV}; in Sec.\\ref{conclusions} I collect all the combined bounds and draw some conclusions. } ",
        "conclusions": "\\label{conclusions}  We have performed in~\\cite{CMSV} a systematical study of the effect of an extra sterile neutrino in all possible contexts, for any choice of its mixing parameters with the active neutrinos and fully including the active/active mixings now established by the results on solar and atmospheric oscillations. We considered cosmology (BBN, CMB, LSS), astrophysics (the Sun, SN...) and terrestrial neutrino experiments (atmospheric neutrinos, reactors, short- and long-base line beams). In these Proceedings I presented in detail the analysis relative to the effects in cosmology and in supernov\\ae\\, focussing on the principles of the computational techniques and on the understanding of the results.   We find no evidence for a sterile neutrino in the present data; figure\\fig{all} collects the {\\bf presents constraints}, in the six limiting cases of mixing with a mass eigenstate $\\nu_{1,2,3}$ or a flavor eigenstate $\\nu_{e, \\mu, \\tau}$. Every specific model of a sterile neutrino identifies a point or a region in one of these spaces or in a suitable combination of them, and should therefore be compared with the reported bounds. In particular, I showed in Sec.\\ref{LSND section} that the region of the LSND sterile neutrino is excluded by the cosmological constraints (unless they are relaxed by some non-standard modification to cosmology).  Finally, I discussed promising {\\bf future probes} of sterile effects. In the context of the cosmological observations, it looks important to improve the measurements of the primordial $^4$He and Deuterium abundances, overcoming the systematic uncertainties. It looks also especially important to exploit the complementarity of the two probes, that are differently sensitive to sterile neutrinos in several regions of the parameter space, also in order to constrain the non-standard cosmological modifications. Moreover, future measurements of the CMB and LSS power spectra will expand the tested range of sterile parameters. Concerning supernov\\ae\\, the best improvement would come from \\ldots the occurring of a new explosion, which would allow to probe sterile effects through the predictions on the neutrino flux and spectrum, in regions that are not easily accessible to other tools (thanks to the long base-line and the extreme matter effects). Although progress must also come on the overall theoretical uncertainties of SN models, most of the results could be neat enough.\\\\ Other important probes for sterile neutrino effects are discussed in~\\cite{CMSV}: among them, solar neutrino experiments at sub-MeV energies and proposed reactor and LBL experiments look most interesting."
    },
    "0410/astro-ph0410408_arXiv.txt": {
        "abstract": "We confront Cardassian models with recent observational data. These models can be viewed either as purely phenomenological modifications of the Friedmann equation, or as arising from cosmic fluids with non-standard properties. In the first case, we find that the models are consistent with the data for a wide range of parameters but no significant preference over the \\lcdm model is found. In the latter case we find that the fact that the sound speed is non-zero in these models makes them inconsistent with the galaxy power spectrum from the Sloan Digital Sky Survey. ",
        "introduction": "The Universe is expanding at an accelerating rate.  We have direct evidence for this from the observed magnitude-redshift relationship of supernovae type Ia (SNIa) \\cite{riess, perlmutter}. Cosmic acceleration requires a contribution to the energy density with negative pressure, the simplest possibility being a cosmological constant. Independent evidence for a non-standard contribution to the energy budget of the universe comes from e.g. the combination of the power spectrum of the cosmic microwave background (CMB) temperature anisotropies and large-scale structure (LSS): \u00a0on the one hand, the position of the first peak in the CMB is consistent with the universe having zero spatial curvature, which means that the energy density is equal to the critical density. On the other hand, observations of large-scale structure, e.g. the clustering of galaxies, show that the contribution of standard sources of energy density, whether luminous or dark, is only a fraction of the critical density.  Thus, an extra, unknown component is needed to explain the observations \\cite{efstathiou, tegmark}.  The model resulting from reintroducing Einstein's cosmological constant, $\\Lambda$, in a universe with baryons, radiation, and cold dark matter (CDM) is consistent with all large-scale cosmological observations like the anisotropies in the CMB radiation and the power spectrum of galaxies \\cite{sdss}. However, the modern interpretation of $\\Lambda$ is that it represents the energy of the vacuum, and from the viewpoint of quantum field theory, the tiny value inferred for it from observations is highly unnatural \\cite{weinberg}.  Faced with this problem, the popular choice is to set $\\Lambda$ to zero and invoke a new component to explain the acceleration.  This does not solve the problem of the smallness of $\\Lambda$; if it is indeed equal to zero, one still needs to understand the physical mechanism behind this. However, one may hope that $\\Lambda=0$ may be easier to explain than a small, but nonzero $\\Lambda$. The question is then what the unknown component driving the accelerated phase of expansion is.  One needs to introduce a component with negative pressure, and this can be done e.g. by invoking a slowly evolving scalar field \\cite{wetterich, peebles, ratra} , or a negative-pressure fluid, e.g. a Chaplygin gas \\cite{kamenshchik, bilic} A negative-pressure fluid, however, can be problematic due to its fluctuations. Fluctuations of the unknown fluid can grow very rapidly and hence LSS surveys can place strict constraints on such models (see e.g. \\cite{bean, sandvik}). In order to circumvent this, it is sometimes assumed that the new fluid does not fluctuate on the scales of interest and its existence is visible only through modified background evolution.  As an alternative to fluid models of dark energy, modifications of the Friedmann equation have been considered by several authors \\cite{dim1, dim2,dim3, dvaliturner, carroll, card, card2}. In an earlier paper \\cite{elgar1}, we investigated a model \\cite{dim1,dvaliturner} where the left-hand side of the Friedmann equation is modified.  In the present paper, we will look at the so-called Cardassian models \\cite{card,card2}, where the right-hand side of the Friedmann equation is modified.  The structure of this paper is as follows.  In section 2 we summarize the basic properties of the Cardassian models.  In section 3 we consider fits to SNIa and WMAP data.  In section 4 we extend the analysis to include probes of the large-scale distribution of matter, and in section 5 we summarize our constraints.  Finally, section 6 contains our conclusions. ",
        "conclusions": "We have considered the observational constraints arising from current CMB, SNIa and LSS data on commonly considered models within the Cardassian framework, the original Cardassian model and the Modified Polytropic Cardassian model. Both models were introduced as an alternative explanation to the acceleration of the universe.  Considering that the models only have an effect on the background evolution, CMB and SNIa observations constrain the allowed parameter space. We find that, within the approximation that only background evolution is modified, using the full CMB spectrum constrains the original Cardassian model significantly and that the model offers no advantage over the \\lcdm model. Adding a new degree of freedom by extending the consideration to the Modified Polytropic Cardassian model relaxes the constraints somewhat but still the parameter space is more constrained than what one finds using the supernovae data only. The constraints indicate that a somewhat non-standard model is slightly preferred over the \\lcdm model but not significantly.  Viewing the Cardassian models as arising from a interacting fluid constrains the scenario critically. Considering the growth of fluctuations along with the SDSS data, we find that any deviation from the \\lcdm model is strongly disfavoured, with or without including baryons in the calculation. The fluid interpretation of the Cardassian model is hence effectively excluded as a viable alternative to the \\lcdm model. This is further supported by recent work \\cite{koivisto}, where the CMB TT power spectrum in the Cardassian model with fluctuations is demonstrated to be modified strongly on large scales.  In short, even small deviations from the \\lcdm model within the fluid interpretation of the Cardassian model are strongly disfavoured by current data. If on the other hand, only background evolution is modified, the data indicates a slight preference over the \\lcdm model but the difference is not significant.  \\label{conclusions}   \\ack \\O E thanks NORDITA, where parts of this work were carried out, for their hospitality. MA and \\O E acknowledge support from the Research Council of Norway through grant 159637/V30."
    },
    "0410/astro-ph0410314_arXiv.txt": {
        "abstract": "We have conducted a three band ($BVI$) variability survey of the globular cluster M3. This is the first three band survey of the cluster using modern image subtraction techniques. Observations were made over 9 nights in 1998 on the 1.2~m.~telescope at the F.~L.~Whipple Observatory in Arizona. We present photometry for 180 variable stars in the M3 field, of which 12 are newly discovered. New discoveries include six SX Phe type variables which all lie in the blue straggler region of the color magnitude diagram, two new first overtone RR Lyrae, a candidate multi-mode RR Lyrae, a detached eclipsing binary, and two unclassified variables. We also provide revised periods for 52 of the 168 previously known variables that we observe. The catalog and photometry for the variable stars are available via anonymous ftp at ftp://cfa-ftp.harvard.edu/pub/kstanek/M3/. ",
        "introduction": "The globular cluster M3 (NGC 5272) is one of the most studied clusters in the Galaxy. Since the work of Bailey\\ (1913) and Shapley\\ (1914), there have been numerous surveys for photometric variability in the cluster. In 1973, Sawyer-Hogg compiled the most complete catalog of variables in the cluster prior to the modern era. Recently Bakos, Benko, \\& Jurcsik\\ (2000) compiled a catalog of 274 variable stars in the cluster, and published improved identification and astrometry for them. M3 is particularly known for containing the largest number of RR Lyrae variables, with more than 182 such stars (Clement et al.\\ 2001). Many groups have focused on studying the population of RR Lyrae variables. Corwin et al.\\ (2001) and Clementini et al.\\ (2004) have published a catalog of 222 confirmed or suspected RR Lyrae variables, and have provided a detailed analysis of the 8 known double-mode RR Lyraes (RRd). Other surveys have focused specifically on searches for variability in blue straggler stars (BSS), notably Kaluzny et al.\\ (1998) discovered only one SX Phe star out of 25 monitored BSS. They also discovered a contact binary star near the base of the red giant branch.  In this contribution we present the data from a $BVI$ variability survey of the cluster performed in the spring of 1998. This is the first three band survey of the cluster using modern image subtraction techniques. New results include the discovery of 6 new SX Phe variables, 2 new first-overtone RR Lyraes, a candidate multi-mode RR Lyrae, a detached eclipsing binary, and 2 new variables that we do not classify. We also provide revised periods for 52 of the 168 previously known variables that we observe.  We describe our observations and data reduction in the following section. In \\S 3 we present the catalog of variables, in \\S 4 we analyze the light curve of the eclipsing binary NV296, in \\S 5 we briefly describe the population of variables using color magnitude diagrams, and in \\S 6 we discuss the results. ",
        "conclusions": "There has been a great deal of interest in studying the variability of BSS in globular clusters. A number of globular clusters show a substantial population of variable blue stragglers. Gilliland et al.\\ (1998) found that 9 out of 47 monitored BSS in 47 Tucanae showed variability above 0.02 mag. Recently Kaluzny et al\\ (2004) identified 35 new SX Phe variables in $\\omega$ Centauri. However, despite having a sizeable population of BSS, Kaluzny et al.\\ (1998) found only a single SX Phe star in M3. We confirm the relative under-abundance of SX Phe stars in M3, finding only 7 variables out of 122 monitored BSS. We note that with image subtraction we would have been able to detect single-mode, short-period variability in the monitored BSS with amplitudes greater than 0.02 magnitudes (see for example NV289 in Figure~\\ref{lc}).  Eclipsing binaries are another interesting target for globular cluster variability surveys. Besides the possibility that binaries play an important role in the dynamical evolution of globular clusters (Hut et al.\\ 1992) and the formation of the BSS (Leonard 1989; Leonard \\& Fahlman 1991), eclipsing binaries can also be used as a tool to determine accurate ages and distances to the clusters (Paczy\\'{n}ski 1997; see Kaluzny et al.\\ 2004 for a discussion of a systematic program to search for detached eclipsing binaries in nearby globular clusters). The location of the components of NV296 on the CMD suggests that the binary may indeed be a member of the cluster. If it is a member, then there must be some mechanism that has allowed the components to avoid evolution into a giant. One hypothesis may be mass transfer onto the primary, however the fact that the light curve seems well fit by a detached eclipsing binary model suggests that mass transfer is not currently occurring at a significant rate. Moreover, the secondary star itself appears to be younger than other F3-4V stars in the cluster. It is necessary to obtain spectra for this system before anything more can be said about its membership. If it is a member, then spectra will be useful in determining the masses of the components.  As mentioned in the introduction, many of the recent studies of M3 have focused on the RR Lyrae population. Most recently Cacciari, Corwin \\& Carney\\ (2004) used data from previous $BVI$ surveys to conduct a detailed analysis of the RR Lyrae population in M3. Other authors have pointed to inconsistencies between existing models for horizontal branch evolution and the observed population (Catelan 2004; Clementini et al.\\ 2004 discuss the observation of rapid evolution among the double mode RR Lyrae). Although we do not perform a systematic analysis of the RR Lyrae population, the data that we present should be useful for further investigations of these stars."
    },
    "0410/hep-ph0410129_arXiv.txt": {
        "abstract": "The showers of cosmic rays entering the Earth's atmosphere are main sources of information on cosmic rays and are also believed to provide information on elementary interactions at energies not accessible to accelerators. In this context we would like first to remind the role of inelasticity $K$ and elementary cross section $\\sigma$ and then argue that similar in importance are fluctuations of different observables. The later will be illustrated by multiplicity fluctuations in hadronic and nuclear collisions. \\vspace{1pc} ",
        "introduction": "The EAS (Extensive Air Showers) initiated by cosmic rays (CR) entering the Earth's atmosphere are our main tool to investigate the CR as such and they are believed to remain our only possible window to look at the elementary interactions at energies not accessible to accelerators \\cite{CR}. Agreeing fully with the first part of this statement we would like to remind here that EAS are themselves macroscopic stochastic processes of variable length and this fact sometimes leads to results, which are unexpected from the elementary particle interactions point of view, the apparent intermittency seen there and connected solely with stochasticity of EAS only \\cite{INTERM} being the best example. Other example worth to mention at this point is the apparent self-organized character of EAS depending in a visible way on their length and clearly seen in data \\cite{SOCEAS}. Both phenomena do not depend on details of elementary interactions used in describing EAS. Actually, as it was widely appreciated some time ago \\cite{SWWW} and is being rediscovered also at present \\cite{Know}, out of numerous parameters of models of elementary interactions used in MC codes describing production and development of EAS the most important are: inelasticity $K$ and elementary cross section $\\sigma$ (actually, the hadron - air nucleus cross section). Out of these two, the inelasticity $K$, i.e., the fraction of energy used for production of secondaries and therefore not available for the subsequent interactions in developing EAS, is nowadays not really a parameter as it used to be before (cf. \\cite{SWWW}) but rather a number which must be reproduced by combination of all parameters used in the present models \\cite{Know}\\footnote{Actually $K$ appears usually together with $\\sigma$, therefore, as we have shown in \\cite{SWWW}, one needs at least two independent different experiments (measuring quantities in which $K$ enters in different way) to be able to estimate $K(s)$ from experiment in a model independent way. For other possibilities see \\cite{KWW} and discussion below.}. ",
        "conclusions": "To summarize: phenomenon of EAS is macroscopic and so complicated that it depends only on a small number of characteristics of elementary interactions. It seems that they are: inelasticity $K$, cross section $\\sigma$ and its fluctuations. We have argued that although we do not measure directly the later we measure fluctuations of multiplicity distributions, which influence such observables as rapidity distributions. It is argued that they should also provide us with some estimations of the fluctuations of the cross section, albeit we do not yet know the respective formulas. \\vspace{-2mm} \\begin{figure}[h] \\setlength{\\unitlength}{1cm} \\begin{picture}(8.0,5) \\special{psfile=figvarq.eps voffset=-15 hoffset=0 hscale=55 vscale=57 angle=0} \\end{picture} \\vspace{-10mm} \\caption{Results on the nonmonoticity of fluctuations represented by $Var(n_{neg})/\\langle n_{neg}\\rangle$ as function of number of participating nucleons (cf. \\cite{Nonmon}).} \\end{figure} \\vspace{-6mm} However, there are already some data, which warn us that to get such connection could be very difficult task. Namely, in Fig. 5 we see clearly that such fluctuations measured in nuclear collisions are nonmonotonic function of the number of participating nucleons, which in turn is proportional to the volume of interaction \\cite{Nonmon}\\footnote{As shown in \\cite{RWW} for nuclear collisions fluctuations in $n$ are directly connected with the equation of state parameter $c=p/\\varepsilon$ and $q$: $Var(n) = (q-1)\\cdot \\langle n\\rangle^2/(1+c)^2 \\simeq 0.28\\cdot \\langle n\\rangle^2$ for $c=1/3$ and limiting value of $q=1.5$.}. This could suggest that monotonicity observed in behaviour of first moments (cf. Fig. 4) could be lost when going to higher moments."
    },
    "0410/astro-ph0410300_arXiv.txt": {
        "abstract": "We report on $VI$ CCD photometry of two fields centered in the region of the open clusters NGC~6404 and NGC~6583 down to $V=22.0$. These clusters have never been studied insofar, and we provide for the first time estimates of their fundamental parameters, namely, radial extent, age, distance and reddening. We find that NGC~6404 radius is 2.0 arcmin, as previously proposed, while NGC~6583 radius is 1.0 arcmin, significantly lower than  previous estimates.\\\\ Both clusters turn out to be of intermediate age (0.5-1.0 Gyr old), and located inside the solar ring, at a Galactocentric distance of about 6.5 kpc. These results make these objects very interesting targets for spectroscopic follow-up to measure their metallicity. In fact they might allow us to enlarge by more than 1 kpc the baseline of the radial abundance gradient in the Galactic disk toward the Galactic Center direction. This baseline is currently rather narrow especially for clusters of this age. ",
        "introduction": "Intermediate age and old open clusters (older than half a Gyr) are widely used to probe the chemical evolution of the Galactic disk (Friel \\& Janes 1993, Carraro \\& Chiosi 1994, Carraro et al. 1998, Friel et al. 2002, Carraro et al. 2004), since they cover the entire life of the disk and are evenly distributed across the disk itself.\\\\ With these objects it is possible to derive the age-metallicity  relationship and  the present and past radial abundance gradients in the Galactic disk: these relations are routinely used to constrain Galactic chemical evolution models (Tosi 1996).\\\\ One of the major limitation of the samples commonly in use is the range in Galactocentric distances: a few clusters are known to be located beyond 12 kpc from the Galactic center, and none is currently known to lie closer than 7.5 kpc from the Galactic center (see Friel et al. 2002, Fig.~3). This is basically due to selection effects; star clusters inside the solar ring do not survive for enough time due to encounters with molecular clouds and in general the higher density environment (Wielen 1971). On the other hand, in the antic-enter direction we expect quite a few clusters due to the low efficiency of clusters formation in the Galaxy periphery.\\\\ \\noindent In an effort to enlarge the distance baseline of intermediate age and old open clusters, we searched for candidates towards the Galactic Bulge, by using criteria similar to those adopted by Phelps et al. (1994), i.e. the presence of a number of similar brightness red stars. This in fact would imply the existence of a red clump, typical of intermediate age-old open clusters.\\\\ This search is complementary to our survey of the open cluster remnants (Villanova et al. 2004) designed to seek for old open clusters in an advanced stage of dynamical evolution, close to their final dissolution and merging with the general Galactic disk field.\\\\ In this paper we report on NGC~6404 and NGC~6583, two clusters located low in the Galactic plane, not very far from the Galactic Center direction (see Table~1) and which fulfill our searching criteria.\\\\ \\noindent The layout of the paper is as follows. Sect.~2 illustrates the observation and reduction strategies. An analysis of  the geometrical structure and star counts in the field of the two clusters are presented in Sect.~3, whereas a discussion of the Color-Magnitude Diagrams (CMD) is performed in Sect.~4. Sect.~5 deals with the determination of clusters reddening, distance and age and, finally, Sect.~6 summarizes our findings.   \\begin{table} \\caption{Basic parameters of the observed objects. Coordinates are for J2000.0 equinox} \\begin{tabular}{ccccc} \\hline \\hline \\multicolumn{1}{c}{Name} & \\multicolumn{1}{c}{$RA$}  & \\multicolumn{1}{c}{$DEC$}  & \\multicolumn{1}{c}{$L$} & \\multicolumn{1}{c}{$B$} \\\\ \\hline & {\\rm $hh:mm:ss$} & {\\rm $^{o}$~:~$^{\\prime}$~:~$^{\\prime\\prime}$} & [deg] & [deg]\\\\ \\hline NGC~6404        & 17:39:37 & -33:14:48 & 355.66 & -1.18\\\\ NGC~6853        & 18:15:49 & -22:08:12 &   9.28 & -2.53\\\\ \\hline\\hline \\end{tabular} \\end{table} ",
        "conclusions": "We have presented the first CCD $VI$ photometric study of the open clusters NGC~6404 and NGC~6583. The CMDs we derive allow us to infer estimates of the cluster basic parameters, which are summarized in Table~4.\\\\ \\noindent In detail, we find that:  \\begin{description} \\item $\\bullet$ Both  clusters are of intermediate age; NGC~6404 is 0.5 Gyr old, NGC~6583 1.0 Gyrs old; \\item $\\bullet$ The reddening $E_{B-V}$ turns out to be 0.92$\\pm$0.05 for NGC~6404 and $0.51\\pm0.05$ for NGC~6583; for both clusters solar metallicity isochrones provide a reasonable fit across the whole CMDs; by the way this metal abundance is not unexpected at the clusters position ; in fact for this age range the radial abundance gradient is almost flat around the solar metallicity; \\item $\\bullet$ We place NGC~6404  and NGC~6583 at about 1.7 and 2.1 kpc from the Sun toward the Galactic Center direction; \\item $\\bullet$ This way they both turn out to be intermediate-age open clusters located inside the solar ring, in a region from the Galactic Center where clusters of this age were never found insofar. \\end{description}   \\noindent Future work should concentrate on obtaining an estimate of the cluster metal abundance through spectra of the RGB stars. The knowledge of the cluster metallicity, which we could not constrain very well, is of paramount importance to better probe the trend of  metallicity across the whole Galactic disk (Friel et al. 2002)."
    },
    "0410/astro-ph0410011.txt": {
        "abstract": "We constructed the composite luminosity function (LF) of clusters of galaxies in the five SDSS photometric bands u,g,r,i and z from the RASS-SDSS galaxy cluster catalog. Background and foreground galaxies are subtracted using both a local and a global background correction to take in account the presence of large scale structures and variations from field to field, respectively. The composite LF clearly shows two components: a bright-end LF with a classical slope of -1.25 in each photometric band, and a faint-end LF much steeper ($-2.1 \\le \\alpha \\le  -1.6$) in the dwarf  galaxy region. The observed upturn of the faint galaxies has a location ranging from -16 +5log(h) in the g band to -18.5 +5log(h) in the z band.  To study the universality of the cluster LF we compare the individual cluster LFs with the composite luminosity function. We notice that, in agreement with the composite LF, a single Schechter component is not a good fit for the majority of the clusters.  We fit a Schechter function to the bright-end of the individual clusters LFs in the magnitude region brighter than the observed upturn of the dwarf galaxies. We observe that the distributions of the derived parameters is close to a Gaussian around the value of the composite bright-end LF parameters with a dispersion compatible with the statistical errors. We conclude that the bright-end of the galaxy clusters is universal.  To study the behavior of the individual faint-end LF we define the Dwarf to Giant galaxy Ratio (DGR) of the single clusters. We notice that the distribution of DGR has a spread much larger than the statistical errors. Our conclusion is that the cluster luminosity function is not universal since the cluster faint-end, differently from the bright-end, varies from cluster to cluster. ",
        "introduction": "The galaxy luminosity function (LF) is one of the most direct observational test of theories of galaxy formation and evolution. Clusters of galaxies are ideal systems within which to measure the galaxy LF for the large number of galaxies at the same distance. There are two main purposes in the study of the cluster LF: the comparison of the galaxy LF in clusters and field and thus the study of the influence of the environment on the global statistical properties of galaxies, and the search for differences in the LF of different clusters as indicators of differences in the galaxy formation due to environmental effects or dynamical processes.  The cluster galaxy over-density with respect to the surrounding field is sufficiently high to efficiently identify members either photometrically through the statistical removal of foreground and background galaxies or spectroscopically. These techniques have been used to measure LFs for individual clusters or to form a composite LF, in order to eliminate the peculiarity of the individual LFs and enhance the underlying possibly universal LF (Dressler 1978; Lugger 1986; Colless 1989; Lugger 1989; Lumsden et al.  1997; Valotto et al. 1997; Rauzy et al.  1998; Garilli et al.  1999; Paolillo et al. 2001; Goto et al.  2002; Yagi et al. 2002).  Many of these studies do not agree on the exact form of the LF.  Several authors (Dressler 1978; Lumdsen 1997; Valotto et al.  1997; Garilli et al. 1999; Goto et al. 2002) found differences between the LFs of different clusters and between cluster and field, while others (Lugger 1986; Colless 1989; Lugger 1989; Rauzy 1998; Trentham 1998; Paolillo et al.  2001) concluded that the galaxy LF is universal in all environments. However, all these works used different techniques to check the universality of the cluster LF. Therefore, it is difficult to understand if their conclusions depend on the different tests beeing applied or to actual physical distinctions. Table 1 summarizes the variations between previous studies in the same color and their $\\sigma$ error limits for the Schechter parameters $M^*$ and $\\alpha$.  We have transformed magnitudes to $H_0=100$ km $\\rm{s}^{-1}$ $\\rm{Mpc} ^{-1}$ without changing their cosmology.  So far, the majority of the studies on the cluster composite LF has concentrated on the slope at the relatively bright end of the cluster LF ($M_g \\le -17$) without taking into account the behavior of the dwarf galaxy population in clusters. Instead much work has been done in the recent years in measuring the faint end ($-18 \\le M_g \\le -10$) of the galaxy LF in several nearby clusters (e.g.  Driver 1994; Smith et al.  1997; Phillipps et al.  1998; Boyce et al.  2001; Beijersbergen et al.  2001; Sabatini et al.  2002; Trentham 2003; Cortese et al.  2004).  The LF of these clusters typically steepens faintward of about $M_g \\sim -18$ by showing the debated upturn of the dwarf galaxies .  The faint end slope $\\alpha$ of the LF in this range of magnitudes typically lies in the range -1.4 to -2.2.  Phillipps et al. 1998 noted that the steepness of the faint end slope appears to depend on the cluster density, with dwarfs being more common in lower density environments.  This is possibly because the various dynamical processes which can destroy dwarf galaxies act preferentially in dense environments.  In this paper we present the analysis of the cluster composite LF based on the second release of the Sloan Digital Sky survey (SDSS DR2, Abazajian et al.  2004). The excellence of the SDSS DR2 in terms of its size, depth and sky coverage and the accurate photometry in 5 different optical wavebands gives unprecedented advantages in comparison to the previous studies.  Firstly, the sky coverage (3324 $\\rm{deg}^2$) gives us the possibility to overcome the well known problem of the statistical subtraction of the galaxy background. We used large areas of the survey to define a mean global galaxy background and a region close to the clusters to determine the local galaxy background in order to check for systematics in the field subtraction. Secondly, the apparent magnitude limit of the SDSS DR2 in all the five bands is sufficiently deep (e.g.  $r_{lim}=22.2$, 95\\% completeness) that, at the mean redshift of our cluster sample ($z \\sim 0.15$), the cluster LF can extend and cover a significant part of the dwarf region, going deeper than in all previous studies of the composite luminosity function (more than 6 magnitudes fainter than M*). Thirdly, the high accuracy of the SDSS photometry in all bands gives us the possibility to measure in a statistically significant way the individual cluster LF with the consequent opportunity to check directly the universality of the LF. Furthermore, the accurate multi-color photometry allows us to use several objectively-measured galaxy properties like galaxy morphology. Finally, our comparison of the cluster and field LFs can done  within the SDSS data.  To calculate the cluster composite LF we used the RASS-SDSS galaxy cluster sample (Popesso et al. 2004), which includes 130 systems observed in X-rays.  The use of the RASS-SDSS galaxy cluster catalog ensures that none of the systems is a simple projection of large scale structure along the line of sight.  In this paper we focus on the bright end of the cluster LF while a more detailed analysis of the faint end will appear in a fortcoming paper. The paper is organized as follows: in sect.  2 we describe the properties of the cluster sample and the optical galaxy photometry; in sect. 3 we explain the methods used in constructing the individual cluster LFs and the methods of the background subtraction, in sect. 4 we describe the methods used for building the Composite LF, in sect. 5 we describe in details our results and finally sect. 6 contains our conclusions.  Throughout the paper we use $\\rm{H}_0=100$ $\\rm{km}$$\\rm{s}^{-1}$$\\rm{Mpc}^{-1}$, $\\Omega_{m}=0.3$ and $\\Omega_{\\lambda}=0.7$.   \\begin{table*} \\begin{center} \\begin{tabular}[b]{cccccc} \\hline Reference & $M^*$ & $\\alpha$ & Band & Ncluster& Luminosity range \\\\ \\hline Goto et al. 2002 & -20.84 $\\pm$0.26 & -1.40$\\pm$0.11 & u& 204 & $-24 \\le M_u \\le -18$ \\\\ \\hline \\hline Schechter (1976) &-19.9$\\pm$0.50 & -1.24 & $b_j$ & 13 & $-22.5 \\le M_J \\le -18.5$ \\\\ Dressler (1978)& -19.7$\\pm$0.50 & -1.25 & F & 12 & $-23.5 \\le M_F \\le -18.5$ \\\\ Colless (1989) & -20.10$\\pm$0.07 & -1.25 & $b_j$ & 14 & $-22.5 \\le M_J \\le -17$ \\\\ Lumsden et al. (1997) & -20.16$\\pm$0.02 & -1.22$\\pm$0.04 & $b_j$ & 46 & $-21 \\le M_b \\le -18$ \\\\ Valotto et al. (1997) & -20.00$\\pm$0.10 & -1.40$\\pm$0.10& $b_j$ & 55 & $-21 \\le M_b \\le -17$ \\\\ Rauzy et al. (1998) & -20.91$\\pm$0.21 & -1.50$\\pm$0.11 & $b_j$ & 28 & $-21 \\le M_b \\le -17$ \\\\ Garilli et al. (1999) & -20.30$\\pm$0.10 & -0.94$\\pm$0.07 & g & 65 & $-22.5 \\le M_g \\le -15.5$ \\\\ Paolillo et al. (2001) & -20.22$\\pm$0.15 & -1.07$\\pm$0.08 & g & 39 & $-24.5 \\le M_g \\le -16.5$ \\\\ Goto et al. (2002) & -21.24$\\pm$0.11 &  -1.00$\\pm$0.06 & g& 204 & $-24 \\le M_g \\lg -18$ \\\\ De Propris et al. (2003) & -20.07$\\pm$0.07& -1.28$\\pm$0.03 & $b_j$& 60 & $-22.5 \\le M_b \\le -16$ \\\\ \\hline \\hline Lugger et al (1989) & -21.31$\\pm$0.13 & -1.21$\\pm$0.09 & R & 9& $-23 \\le M_R \\le -18.5$ \\\\ Garilli et al. (1999) & -20.66$\\pm$0.16 & -0.95$\\pm$0.07 & r & 65 & $-22.5 \\le M_r \\le -15.5$ \\\\ Paolillo et al. (2001) & -20.67$\\pm$0.16 & -1.11$\\pm$0.08 & r & 39 & $-24.5 \\le M_r \\le -16.5$ \\\\ Yagi et al. (2002) & -21.30$\\pm$0.20 & -1.31$\\pm$0.05 & $R_C$ & 10 & $-23.5 \\le M_Rc \\le -16$ \\\\ Goto et al. (2002) & -21.44$\\pm$0.05 & -0.85$\\pm$0.03 & r & 204 & $-24 \\le M_r \\le -18$ \\\\ \\hline \\hline Paolillo et al. (2002) & -20.85$\\pm$0.20& -1.09$\\pm$0.11& i & 39 & $-24 \\le M_i \\le -17$ \\\\ Goto et al. (2002) & -21.54$\\pm$0.08 & -0.70$\\pm$0.05 & i & 204 & $-24 \\le M_i \\le -18$ \\\\ \\hline \\hline Goto et al. (2002) & -21.59$\\pm$0.06 & -0.58$\\pm$0.04 & z & 204 & $-24 \\le M_z \\le -18$ \\\\ \\hline \\hline \\end{tabular} \\caption{ Schechter parameters fitted to the Composite LF retrieved in the literature.} \\end{center} \\end{table*} ",
        "conclusions": "The  main conclusion of  our analysis  are as follows: \\begin{itemize} \\item we determine the   composite LF of galaxies in   clusters from the SDSS data.  The LF clearly shows a bimodal behavior with an upturn and a evident steepening in the faint magnitude range in any SDSS band.  The LF is well fitted by the sum of two Schechter functions.  The results are well confirmed by different methods of background subtraction. The observed upturn of the faint galaxies has a location ranging from -16 +5log(h) in the g band to -18.5 +5log(h) in the z band. \\item   The  bright end LF shows   the classical slope  of -1.25 in each photometric band, while $M^*$ is brighter in the red bands than in the blue bands.  The distribution of the Schechter parameters obtained fitting only the bright end of the individual cluster LF is close to a Gaussian around the corresponding value of the composite bright-end LF.  We check the dependence of the Schecter parameters of the composite LF on the clustercentric distance calculating the LF within different cluster apertures.  We do not find any significant variation of the results with different apertures.  Therefore, we conclude that the bright-end of the galaxy clusters is universal in different cluster environments, both in different systems and in different locations within the same cluster. \\item The faint end  LF is much steeper than the bright  end  LF with slope $-2.5 \\le \\alpha \\le -1.6$.  We apply different tests to check whether the observed faint end in the single clusters is due to the presence of background large scale structures or a second cluster on the line of sight.  To check the first possibility we measure the individual cluster LF with a color cut method to identify the cluster members. We obtain the same slope observed with the statistical background subtraction.  Moreover, we observe that the faint population is dominated by galaxies with colors compatible with late type galaxies at the cluster redshift.  We, then, conclude that the observed steepening of the cluster LF is due to the presence of a real population of faint cluster galaxies.  \\item We defined the Dwarf  to Giant galaxy Ratio DGR as the ratio between the number of galaxies in the magnitude range $ -18 \\le M \\le -16.5$ and the number of galaxies brither than -20 mag.  In each waveband the DGR seems to slightly increase from the very center 0.3 Mpc $\\rm{h}^{-1}$ to 1.0 Mpc $\\rm{h}^{-1}$.  The distribution of DGR of the single clusters has a peak around the value predicted by the composite bright-end and and a large spread at larger values. We, then, conclude that the faint end of the cluster LF is not universal and that the fraction of dwarf galaxies varies from cluster to cluster. We check the relation between the DGR and the cluster richness and between DGR and $M^*$ through the Spearman's rank coefficient and we do not find any correlation between the parameters. \\end{itemize} We compare the above results with the field LF calculated in the Sloan Digital Sky Survey and in other surveys.  The magnitude range covered by the SDSS field LF allows us to compare only the bright end of the cluster luminosity function with the field LF.  Unfortunately there is no good agreement between the results retrieved in the literature.  Therefore we cannot perform a conclusive comparison between the LF of the different environments.  Moreover, several surveys find evidence for the presence of an upturn at the faint end of the field luminosity function in agreement with our results for the cluster LF.  In particular Zucca et al. (1997) find in the ESP field LF evidence for the presence of a late type galaxy population dominating the faint end of the field luminosity function.  However, it is important to stress that this is only a qualitative comparison and does not allow us to any conclusion about the nature of the faint population in clusters and in the field.  We can only conclude that the faint end of the cluster LF is systematically steeper than the field LF, although the field LF seems to show some evidences for an excess of galaxies in the faint magnitude range relative to a Schechter function.  Hierarchical clustering theories of galaxy formation generically predict a steep mass function of galactic halos (Kauffmann, White \\& Guideroni 1993; Cole et al. 1994). This is in conflict with the flat galaxy LF measured in the field and in diffuse local groups, but not with the steep LF measured in many clusters.  However in the hierarchical universe, clusters form relatively recently from the accretion of smaller systems. The dynamical processes that operate in clusters are destructive. Ram pressure stripping (e.g. Moore \\& Bauer 1999) and gravitational tides/galaxy harassment (e.g.  Moore et al. 1996, 1998) will both tend to fade galaxies by removing gas or stripping stars. These processes are most effective for less massive, less bound systems. Hence, we might expect to see a flattening of the faint end slope in clusters compared to the field, rather than the observed steepening.  Understanding the nature of the observed faint galaxy population requires a more detailed study of the galaxy population in cluster through the analysis of the morphological type, the colors and the study of the relation between the fraction of dwarf galaxies and the cluster parameters such as the cluster mass, velocity dispersion or the X-ray luminosity.  Moreover, the origin of this faint population can be understood only if a conclusive comparison between cluster and field is possible. At the moment, as we discussed above, the SDSS field LF based on the Sloan spectroscopic galaxy sample does not allow to an exhaustive comparison and analysis of the different environments.  It is clear from our results and from all the existing works in literature that the composite bright end of the cluster LF can give useful information on the global cluster properties (such as the total optical luminosity, which is dominated by the very bright cluster galaxies), but it does not provide useful information on the cluster galaxy population as a whole. Equally it is clear that a Schechter function is a good fit of the cluster LF only in a very restricted magnitude range (the bright end). The photometric data available now should make it possible to consider non-parametric comparisons between the individual and the composite cluster LFs using the full range of the available data. Our results on the dwarf galaxy fraction are the first step in this direction, but it must be possible to devise a test that does not require a split in bright and faint galaxies but consider the cluster galaxy population as a whole.  \\vspace{2cm}  Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P.  Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S.  Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions.  The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0410/astro-ph0410136_arXiv.txt": {
        "abstract": "We present a description of the data reduction methods and the derived catalog of more than 1600 X-ray point sources from the exceptionally deep January 2003 Chandra X-ray Observatory ($Chandra$) observation of the Orion Nebula Cluster and embedded populations around OMC-1. The observation was obtained with $Chandra$'s Advanced CCD Imaging Spectrometer (ACIS) and has been nicknamed the $Chandra$ Orion Ultradeep Project (COUP). With an 838~ks exposure made over a continuous period of 13.2 days, the COUP observation provides the most uniform and comprehensive dataset on the X-ray emission of normal stars ever obtained in the history of X-ray astronomy. ",
        "introduction": "\\label{introduction_section}  The processes of star formation and early stellar evolution have proved to be complex in many respects. Gravitational collapse can be hindered by magnetic and turbulent pressures and is accompanied by collimated jets and outflows during the protostellar phase. The dynamics and evolution of circumstellar disks where planets form are still under study.  The birth of stars rarely occurs in isolation. Most stars form in a hierarchy ranging from binaries to rich clusters. These processes, which take place in thermodynamically cold and neutral media with characteristic energies of $\\ll1$ eV per particle, paradoxically produce and are subject to violent high energy processes with characteristic energies of $\\gtrsim 10^3$~eV.  The principal evidence for this is X-ray emission (and sometimes nonthermal radio emission) from stars throughout their pre-main sequence (PMS) evolution \\citep{Feigelson99}.  PMS X-ray emission is elevated $10^1$ to $\\gtrsim 10^4$ times above typical main sequence levels for stars with masses $0.1 M_\\odot \\la M \\la 2 M_\\odot$. High-amplitude X-ray variability and hard spectra, supported by multiwavelength studies such as photospheric Zeeman measurements and starspot mapping, indicate that the emission is primarily attributable to solar-type magnetic flares where plasma is heated to high temperatures by violent reconnection events in magnetic loops. However, the empirical relationships between PMS X-ray emission and stellar age, mass, radius, and rotation differ from those seen among solar-type main sequence stars \\citep{Flaccomio03b, Feigelson03}. We have an inadequate understanding of the astrophysical origins of this magnetic activity.  It is also unclear whether or not the high energy radiation from magnetic flaring has important astrophysical effects. For example, it is possible that some planets form in disk regions that are rendered turbulent from MHD instabilities induced by PMS X-rays \\citep[e.g.,][]{Glassgold00, Matsumura03}. X-rays from embedded stars may affect ambipolar diffusion, and thus future star formation, in the surrounding molecular cloud. X-ray images are also particularly effective in locating low mass companions around luminous young stars. There are thus a variety of astrophysical reasons to investigate the magnetic activity of young stars.  Due to the wide range of activity levels and other properties, it is advantageous to study large samples of stars.  The nearest rich and concentrated sample of PMS stars is the Orion Nebula Cluster (ONC), also known as the Trapezium Cluster or Ori Id OB association. The OB members of the ONC illuminate the Orion Nebula (= Messier 42), a blister H~$_{\\rm{II}}$ region at the near edge of Orion A, the nearest giant molecular cloud ($D \\simeq 450$ pc). The ONC has $\\simeq 2000$ members within a 1 pc ($8^{\\prime}$) radius sphere with 80\\% of the stars younger than 1 Myr \\citep{Hillenbrand97}. It was the first young stellar cluster to be detected in the X-ray band \\citep{Giacconi72} and non-imaging studies soon found that the X-ray emission is extended on scales of a parsec or larger \\citep{denBoggende78,Bradt79}. Early explanations for the Orion X-rays included winds from the massive Trapezium stars colliding with each other or the molecular cloud, and hot coronae or magnetic activity in lower mass T Tauri stars.  The $Einstein$ \\citep{Ku79}, $ROSAT$ \\citep{Gagne95, Geier95, Alcala96} and $ASCA$ \\citep{Yamauchi96} imaging X-ray observatories established that both the Trapezium stars and many lower-mass T Tauri stars contribute to the X-ray emission.  But these early studies could detect only a modest fraction of the ONC stars due to the high stellar densities and heavy absorption by molecular material along the line of sight. These limitations are alleviated by NASA's Chandra X-ray Observatory ($Chandra$). Its mirrors provide an unprecedented combination of capabilities - a wide spectral bandpass and $\\lesssim 1$\\arcsec\\/ spatial resolution - and are particularly well-suited to the study of crowded and absorbed clusters of faint X-ray sources.  The ONC was thus intensively studied during the first year of the $Chandra$ mission with several instrumental setups: the Advanced CCD Imaging Spectrometer (ACIS) in imaging mode \\citep{Garmire00, Feigelson02a, Feigelson02b, Feigelson03} and as detector for the High Energy Transmission Grating Spectrometer \\citep{Schulz00, Schulz01, Schulz03}, and with the High Resolution Imager \\citep[HRI;][]{Flaccomio03a, Flaccomio03b}.  While many valuable results emerged from these early $Chandra$ studies, it was recognized that more would be learned from a much deeper and longer exposure of the Orion Nebula region.  We present here a unique $\\sim 10$ day, nearly continuous observation of the Orion Nebula obtained during the fourth year of the mission, nicknamed the $Chandra$ Orion Ultradeep Project (COUP). This paper describes the COUP observations (\\S \\ref{observations_section}), the data reduction (\\S \\ref{data_selection_section}), source detection and characterization (\\S \\ref{source_detection_section}-\\ref{luminosities_section}), and identification with previously known PMS stars (\\S \\ref{stellar_counterparts_section}-\\ref{nondetections_section}). We provide extensive machine-readable tables of the X-ray sources and properties, and an atlas of their location in the sky (\\S \\ref{atlas_section}).  Other COUP papers will discuss various astronomical and astrophysical issues emerging from these observations. ",
        "conclusions": "\\label{summary_section}  This paper presents the observations, data analysis methodology, and tabulated results for the $Chandra$ Orion Ultradeep Project.  We describe techniques used in the COUP data reduction, derivation of time-averaged source properties, and production of the source catalog. The results of this study are the tables of the 1616 COUP source properties (Tables \\ref{COUP_src_table}, \\ref{COUP_pho_table}, \\ref{COUP_spe_table}, \\ref{COUP_var_table}, \\ref{COUP_lum_table}, \\ref{COUP_opt_table}, and \\ref{COUP_nir_table}), tables of undetected sources (Tables \\ref{COUP_lim opt_table} and \\ref{COUP_lim 2m_table}), map of the field (Figure \\ref{expanded_view_fig}), and the 1616 page atlas detailing the analysis and findings for each source.  Our data reduction procedure seeks to use the most well-developed methods available for the identification and extraction of point sources in $Chandra$ ACIS images.  The source detection procedure is based on two wavelet-based search algorithms, run on datasets constructed in several bands, supplemented by visual inspection of the image. The dataset is optimized for maximum reduction of background and maximum spatial resolution with subpixel event repositioning. We identify 1616 X-ray sources in the $\\sim 17\\arcmin \\times 17\\arcmin$ ACIS-I field.  For each source, we perform data extraction, pile-up analysis, spectral analysis, variability analysis, and broad-band luminosity determinations using the sophisticated semi-automated IDL-based {\\rm{\\it acis\\_extract}} package.  The algorithms used in {\\rm{\\it acis\\_extract}} successfully address many difficulties related to the COUP observation such as the position-dependent point spread function, multiple exposures with different pointings and time gaps, source crowding, piled-up sources, nonuniform local background and exposure map, and so on. Our treatment of photon pile-up for strong sources using annular extraction regions is first used here in a formal procedure. The {\\rm{\\it acis\\_extract}} package, which is available to the public, can be used to create point source catalogs for many $Chandra$ studies.  These images, atlases and tables serve as the foundation for various studies of the COUP observations. These will include investigation of X-ray emission from different classes of pre-main sequence stars: OB, intermediate-mass, T Tauri and brown dwarfs; deeply embedded stars which are often younger protostars; and binary and multiple star systems.  We will also investigate the astrophysical issues such as the magnetic field generation and geometries in pre-main sequence systems, plasma energetics in powerful magnetic reconnection flaring events, and X-ray ionization of circumstellar disks and cloud material. Altogether, the COUP data products presented here offer the most detailed and comprehensive information to date concerning the X-ray emission and magnetic activity in pre-main sequence stars.  We thank George Chartas (Penn State) for providing the ARF correction tool, $\\rm {\\it xpsf}$, and acknowledge the unusual efforts of the {\\it Chandra} Mission Operations group in scheduling the COUP observations. We also thank the referee, M. Muno, for his time and many useful comments that improved this work. COUP is supported by {\\it Chandra} guest observer grant SAO GO3-4009A (E.\\ D.\\ Feigelson, PI). This work was also supported by the ACIS Team contract NAS8-38252, HRC Team contract NAS8-39073. EF, FD, GM and SS acknowledge financial contribution from Italian MIUR (COFIN-PRIN 2001) and from INAF. MJM and G. Meeus acknowledge support from the European Commission Research Training Network ``The Formation and Evolution of Young Stellar Clusters'' (HPRN-CT-2000-00155). 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."
    },
    "0410/astro-ph0410246_arXiv.txt": {
        "abstract": "Effects of the anisotropy of Cooper pairs in spin-triplet ferromagnetic superconductors are investigated on the basis of the Ginzburg-Landau theory. A special attention is paid to the triggering of the superconducting state by the ferromagnetic order. The ground states of these superconductors are outlined and discussed. The idea about a possible coexistence of ferromagnetism and spin-triplet superconductivity in neutron stars is introduced. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0410/astro-ph0410593_arXiv.txt": {
        "abstract": "We explore the effect of contamination of intermediate redshift Type~Ia supernova samples by Type~Ibc supernovae. Simulating observed samples of Ia and mixed Ibc/Ia populations at a range of redshifts for an underlying cosmological concordance model ($\\Omega_{m}=0.27$, $\\Omega_{\\Lambda}=0.73$), we find that even small contamination levels, $2-5$\\% may bias the derived $\\Omega_{\\Lambda}$ and $\\Omega_{m}$ towards larger values. We thus emphasize the need for clean samples of Type~Ia SNe for accurate measurements of the cosmological parameters. We also simulate a SN sample similar to the fiducial SNAP detected distribution (Kim et~al. 2004), but include Ibc contamination. For this distribution we illustrate the effect of Ibc contamination on the distance modulus vs. redshift diagram for low and high precision measurements. ",
        "introduction": "Type~Ia supernovae (SNe) are considered astronomical standard candles and have been used to measure the geometry and dynamics of the Universe \\citep{Riessetal98,Perletal99}. This exceedingly difficult measurement relies on observations of local supernovae samples to calibrate relationships between absolute peak magnitude and more easily obtained observables such as light-curve width, color, etc. (see \\citet{Leib01} for a review). These calibrations are then applied to the observables in higher redshift samples to derive distance moduli. Due to the extreme faintness of intermediate to high redshift SN and the shifting of identifying spectral features outside the optical window, spectroscopic classification becomes difficult, and photometry must often suffice. Fortunately, there is a clear observational difference between hydrogen-rich (Type~II) and hydrogen-poor (Ia, Ib, and Ic) SNe, namely, the UV deficit (UV radiation is blocked in hydrogen-poor SN explosions by metal lines). This property can be used to distinguish between Type~I and II SNe explosions. To date there is no photometric technique to distinguish between Type~Ia and Ibc SN (although see Gal-Yam et~al., 2004), and one must perform a spectroscopic classification. The characteristic spectroscopic difference between Type~Ia and Ibc SNe is the deep absorption trough at 6140\\AA~ in Ia spectra, which is due to the blueshifted \\ion{Si}{2} $\\lambda\\lambda 6347, 6371$\\AA~ feature. Unfortunately this is redshifted out of the optical passband for $z>0.5$, making classification of Type~I SNe difficult. Alternative spectral discriminants between Ia and Ibc SNe are \\ion{Si}{2} $\\lambda 4130$ (e.g. Coil et~al. 2000), \\ion{Fe}{2} $\\lambda 4555$ and \\ion{Mg}{2} $\\lambda 4481$ \\citep{Barrisetal04}. Spectroscopic identifications are made using template-matching, but no statistical information is given about misidentifications with this method \\citep{Tonetal03,Retal04b,Barrisetal04}.    Assuming the same relative Ia and Ibc rates in nearby SNe samples, Riess~et.~al~(2004) and Barris~et~al.~(2004) reasoned that the rate of contamination from Type~Ibc SN in high redshift samples should be less than $\\sim10$\\%. However, the star formation rate increases rapidly at up to $z\\sim 1-2$, and this is likely to have a significant effect on the relative rates of SN types. As Type~Ibc SNe come from massive stars (M~$> 20-30$~M$_{\\odot}$), their rate depends on the star formation rate in the previous 10~Myr. It is well established that the global star formation rate increases at higher redshifts. Hence, assuming a delay for Ia explosions, we expect that the rate of Type~Ibc SNe increases over the short term, and that the contribution from Ibc SNe becomes more important at higher redshift. Motivated by this, we investigate here what effect contamination by Type~Ibc SNe has on cosmological parameters derived from Ia SNe samples. The optimal way to determine the systematic effect in a particular survey is through simulations using the diagnostic machinery associated with said survey. However, the general nature of the effect can be investigated robustly with simulations such as the one we present here. ",
        "conclusions": "We find that even small contamination fractions (5\\%) may bias the measurement of $\\Omega_{M}$ and $\\Omega_{\\Lambda}$. However, we have also illustrated the ease with which Ibc SNe can be identified when the observational error and the dispersion about the corrected Ia peak magnitudes is low. Based on the calculations presented here, the straightforward effect of contamination from Type~Ibc SN in intermediate-to-high redshift SN samples is to reduce the $\\chi^{2}$ goodness-of-fit to the true cosmological parameters. This manifests itself as an increase in the minimum value of the reduced $\\chi^{2}$ grid. For a flat universe with $\\Omega_{M}=0.27, \\Omega_{\\Lambda}=0.73$ even small numbers (5\\%) of Type~Ibc SNe may push both values upwards. Erroneously small ``probability'' curves can be generated by assuming the model is a good fit and that the minimum reduced $\\chi^{2}$ is close enough to 1 for a meaningful conversion to a probability grid. However, if the observational errors are large, it is difficult to distinguish between a minimum reduced $\\chi^{2}$ that is larger due to a mismatch of the model and the observations due to contamination or from simple stochastic effects.  If the observational errors are large enough ($> 0.3$), then a small contamination fraction (5\\%) by Ibc SNe can shift the location of the $\\chi^{2}$ minimum without making $\\chi^{2}$ itself large. In such a case the error contours will be large, so while the contours are not strictly statistically correct, they are still meaningful in the sense that the underlying cosmological parameters have close to a 99\\% chance of being within the $2\\sigma$ contours. If the observational error approaches the scatter in the corrected 'standard candle' Ia magnitudes, the errors bars will shrink. However, with the same amount of Ibc contamination, as the observational error steadily decreases, the reduced $\\chi^{2}$ becomes steadily greater than 1. Then the conversion from a $\\chi^{2}$ grid to probabilities is not valid. One may assume that the random errors are underestimated and renormalize; the error contours may then become large enough to encompass the underlying $\\Omega_{M},\\Omega_{\\Lambda}$ values. But if this precision is achieved, outliers should already be readily discriminated in the $\\Delta (m-M)$ vs. $z$ plane, and a simple $3-5 \\sigma$ clipping routine may suffice. Simply having large numbers of SNe with the precision of current surveys would also allow one to use $\\sigma$-clipping to remove Ibc outliers, however, this will not achieve significantly greater precision in the measurement of cosmological parameters."
    },
    "0410/astro-ph0410520_arXiv.txt": {
        "abstract": "{We present deep VLA and HST observations of the large-scale jet in 3C\\,273 matched to 0\\farcs3 resolution.  The observed spectra show a significant flattening in the infrared-ultraviolet wavelength range.  The jet's emission cannot therefore be assumed to arise from a single electron population and requires the presence of an additional emission component.  The observed smooth variations of the spectral indices along the jet imply that the physical conditions vary correspondingly smoothly.  We determine the maximum particle energy for the optical jet using synchrotron spectral fits.  The slow decline of the maximum energy along the jet implies particle reacceleration acting along the entire jet.  In addition to the already established global anti-correlation between maximum particle energy and surface brightness, we find a weak positive correlation between small-scale variations in maximum particle energy and surface brightness.  The origin of these conflicting global and local correlations is unclear, but they provide tight constraints for reacceleration models. ",
        "introduction": "\\label{s:intro} Jets are collimated outflows, thought to be launched from an accretion disk around a central compact object, which can be a supermassive black hole, a stellar-mass black hole or neutron star, or a young stellar object (YSO). They transport mass, energy, both linear and angular momentum as well as electromagnetic fields outward from the central object.  A detailed understanding of the formation of these jets, their connection to the accretion disk from which they are launched, and the physics governing their internal structure and observable properties has not yet been achieved. Here we consider the synchrotron emission from the kiloparsec-scale jet of the quasar \\object{3C\\,273}, which is one of the brightest and largest and therefore an instructive sample case.  Jets became part of the standard model of extragalactic radio sources \\citep{BBR84} to link the lobes, which emit the bulk of the synchrotron radio luminosity, with the active galactic nucleus (AGN) of the host galaxy. In this model, jets merely transport energy to feed the lobes. In the most powerful sources, a double shock structure terminates the jets, consisting of an outer bow shock and contact discontinuity separating the jet material from the external medium and an internal shock (Mach disk) at which the relativistic flow is decelerated and bulk kinetic energy is channeled into highly relativistic particles through a shock acceleration mechanism. These particles emit the synchrotron radiation observed from the lobes. The radio hot spot is usually assumed to coincide with the Mach disk. The {\\it optical} synchrotron emission observed from some hot spots can also be explained by first-order Fermi acceleration at a jet-terminating shock \\citep{MH86,HM87,magnumopus89,hs_II}.  In this model, the electrons responsible for radio synchrotron emission from the jets themselves are accelerated near the black hole and then simply advected with the jet flow. However, observations of high-energy (optical and X-ray) synchrotron radiation from 3C\\,273 and other jets force us to revise this picture.  Electrons with the highly relativistic kinetic energies required for synchrotron emission in the infrared and optical have a very short radiative lifetime. This lifetime is much less than the light-travel time along the jet in 3C\\,273 \\citep[already noted by ][]{GS64} and other jets. Observations of optical synchrotron emission from such jets \\citep{RM91,M87,SU02} as well as from the ``filament'' near Pictor~A's hot spot \\citep{RM87,PRM97} therefore suggest that in addition to a localized, ``shock-like'' acceleration process operating in hot spots, there is an extended, ``jet-like'' mechanism at work in radio sources in general and 3C\\,273's jet in particular \\citep{hs_II}.  The extended mechanism may also be at work in the lobes of radio galaxies, where the observed maximum particle energies are above the values implied by the losses within the hot spots \\citep{Mei96} and by the dynamical ages of the lobes \\citep{BR00}.  We therefore make a clear distinction between the emission from the hot spot itself, the lobes and the body of the jet. Although all emit synchrotron radiation, the physical processes accelerating the particles emitting in these regions may be quite different.  The lifetime problem has been exacerbated by observations at even higher frequencies: Einstein and ROSAT observations, for example, showed X-ray emission from the jets in M87 \\citep{SGF82,BSH91,NMRF97} and 3C\\,273 \\citep{HS87,Roe00}. More recently, observations with the new X-ray observatory Chandra showed extended X-ray emission from many more jets, like PKS~0637$-752$ \\citep{Sch00} and Pictor\\,A \\citep{Wil01} as well as other jets and hot spots.  Chandra also supplied the first high-resolution X-ray images of the jets in 3C\\,273 and M87 \\citep{Mareta01,Sam01,Mar02}. The X-rays from these objects seem to be of non-thermal origin \\citep[for an overview, see][]{HK02}: they could at least partially be due to synchrotron emission \\citep{Roe00,Mareta01,Mar02}. Alternatively, inverse-Compton scattering could be responsible for the X-rays. The photon seed field can be provided by the synchrotron source itself if it is sufficiently compact, for example in the hot spots of Cygnus A \\citep{HCP94,WilsonCygA}. If the bulk flow of a jet is still highly relativistic on large scales, the boosted energy density of the cosmic microwave background radiation field can lead to the observed X-ray fluxes \\citep{Cel00,Tav00}. In all cases, those electrons producing the radio-optical synchrotron emission suffer additional losses from the inverse-Compton scattering, decreasing their cooling timescale even below the synchrotron cooling scale.  Thus, the fundamental question posed by the observation of optical extragalactic jets is the following: how can we explain high-frequency synchrotron and inverse-Compton emission far from obvious acceleration sites in extragalactic jets? While information on the source's magnetic field structure may be obtained from the polarization structure, the diagnostic tool for the radiating particles is a study of the synchrotron continuum over as broad a range of frequencies as possible, \\emph{i.\\,e.,} from radio to UV wavelengths, and with sufficient resolution to discern morphological details. The shape of the synchrotron spectrum gives direct insight into the shape of the electron energy distribution, thus also constraining the emission by the inverse-Compton process at other wavelengths.  Here, we consider the shape of the synchrotron spectrum of the jet in 3C\\,273.  \\subsection{The  jet in 3C\\,273}  This optical jet was first detected on ground-based images. Like M87, its optical brightness and length are so unusually large that it was detected even before radio jets were known. It appears to consist of a series of bright knots with fainter emission connecting them (see Fig.\\,\\ref{f:obs.nir.IRmap}).  \\citet{GS64} described the jet's optical spectrum as ``weak, bluish continuum'', suspecting that this was synchrotron radiation. This was confirmed by \\citet{RM91} through optical polarimetry.  3C\\,273's radio jet extends continuously from the quasar out to a terminal hot spot at 21\\farcs4 from the core, while optical emission has been observed only from 12\\arcsec\\ outward.\\footnote{For the conversion of angular to physical scales, we assume a flat cosmology with $\\Omega_{\\mathrm{m}}=0.3$ and $H_{0} = h_{70}\\times 70\\,\\mathrm{km\\,s^{-1}Mpc^{-1}}$, leading to a scale of $2.7 h_{70}^{-1} \\mathrm{kpc}$ per second of arc at 3C\\,273's redshift of 0.158.}  We concentrate on this ``outer'' part of the jet here, and will report observations of optical emission from the inner jet with the VLT in a future publication \\citep[see also][]{Marea03}. \\citet{Baheta95} presented the first HST imaging of this jet, noting the jet resembles a helical structure.  Prior to the present work, synchrotron spectra have been derived for the hot spot and the brightest knots using ground-based imaging in the radio \\citep{jetII}, near-infrared $K^{\\prime}$-band \\citep{NMR97} and optical $I, R, B$-bands \\citep{RM91} at a common resolution of 1\\farcs3 \\citep{MNR96,Roe00}. This radio-to-optical continuum can be explained by a single power-law electron population resulting in a constant radio spectral index\\footnote{We define the spectral index $\\alpha$ such that $f_\\nu \\propto \\nu^{\\alpha}$.} of $-0.7$, but with a high-energy cutoff frequency decreasing from $10^{17}\\,$Hz to $10^{15}$\\,Hz outward along the jet.  Here, we present VLA and HST NICMOS observations (\\S\\ref{s:obs}) in addition to the WFPC2 data already published in \\citet{Jes01}. Together, these constitute a unique data set in terms of resolution and wavelength coverage for any extragalactic jet --- only M87 is similarly well-studied \\citep[][and references therein]{MRS96,SBM96,HB97,Pereta99,Per01}. Using these observations at wavelengths 3.6\\cm, 2.0\\cm, 1.3\\cm, 1.6\\micron, 620\\nm\\ and 300\\nm, we derive spatially resolved (at 0\\farcs3) synchrotron spectra for the jet (\\S\\ref{s:results}). By fitting model spectra according to \\citet{HM87}, we derive the maximum particle energy everywhere in the jet in order to identify regions in which particles are either predominantly accelerated, or predominantly lose energy (\\S\\ref{s:ana}).  The model spectra reveal excess near-ultraviolet emission above a synchrotron cutoff spectrum accounting for the emission from radio through optical, which implies that a two-component model is necessary to describe the emission. The radio--optical--X-ray spectral energy distributions (SEDs) suggest a common origin for the UV excess and the X-rays from the jet \\citep[\\S\\ref{s:disc.pop}; see also ][]{Jes02}. By considering just the optical spectral index, we concluded in \\citet{Jes01} that particles must be reaccelerated along the entire jet. Here we confirm this conclusion by using the full spectral information from radio to near-ultraviolet (\\S\\ref{s:disc.acc}). We show that the observed changes of cutoff energy and surface brightness along the jet can be jointly explained as effects of changes in the magnetic field and the Doppler beaming parameter along the jet (\\S\\ref{s:disc.corr}).  We conclude in \\S\\ref{s:summary}. ",
        "conclusions": "\\label{s:disc}  \\subsection{The hot spot}\\label{s:disc.HS} \\citet{MH86} presented a model for particle acceleration in the hot spot of 3C\\,273. They showed that the hot spot spectrum can be understood by modelling the hot spot as a planar shock at which first-order Fermi acceleration of particles takes place (details of the spectral modelling are contained in \\citet{HM87}, while further observations of this and other hot spots have been presented by \\citet{magnumopus89,hs_II}). This model predicts an offset between the peaks of the radio and optical emission of 0\\farcs2, arising from the vastly different synchrotron loss scales of electrons being advected downstream from the shock itself: ``radio'' electrons lose a negligible fraction of their total energy, while ``optical'' electrons quickly lose their entire energy as they move away from the acceleration region.  Recall that the strong compression at the shock leads to an increased magnetic field compared to the jet flow, and hence to stronger losses.  Figure\\,\\ref{f:res.im.morph.HS} clearly shows the predicted offset in the position of the hot spot (H2) between radio and optical wavelengths. This strengthens the identification of H2 with a planar shock with strongly localised particle acceleration.  We note that again the minimum-energy magnetic field determined for the hot spot agrees with the magnetic field determined from the hot spot spectrum by \\citet[][ the only difference arises due to different values of the parameter $k$ in the minimum-energy field determination]{hs_II}.  Thus, at 0\\farcs3 resolution, the \\citet{HM87} hot spot model is consistent with all observations. Moreover, our new observations have confirmed the predicted offset in the hot spot's position.  However, it is remarkable that even the hot spot shows a UV excess, indicating that the hot spot, too, has a more complicated internal structure than previously assumed.  In addition, the presence of emission from the tip of the jet H1 \\citep[which has been called a ``precursor'' in the literature; see][e.g.]{FD85} and its relation to the hot spot must be explained.  The next tests of this model will involve a detailed analysis of the hot spot's radio morphology at the highest available resolution, which we defer to a future paper.  \\subsection{Flatter spectra at higher frequencies: the need for a second emission component}\\label{s:disc.pop}  When describing synchrotron spectra as power laws, the underlying assumption is a power-law electron energy distribution.  As noted above, a non-idealised spectrum with finite maximum particle energy and rapid pitch-angle scattering will furthermore exhibit a quasi-exponential cutoff at some frequency \\nuc, the synchrotron frequency corresponding to the highest electron energies.  This cutoff implies that the spectrum has a convex shape, so that the spectrum steepens progressively towards higher frequencies. Any high-frequency flux must therefore lie below a power-law extrapolation from lower frequencies.  In employing \\emph{broken} power laws, the only physical break is a steepening of the spectrum from $|\\alpha_\\mathrm{low}|$ to $|\\alpha_\\mathrm{high}| = |\\alpha_\\mathrm{low}| + 1/2$ near some break frequency $\\nub < \\nuc$.  This break arises in a model of the synchrotron source as a loss region into which a power-law distribution of electrons with approximately constant number density and maximum energy is continuously injected. The break is produced by adding up the contributions from the electron population observed at different times since acceleration, \\ie with different cutoff frequencies.  The magnitude of the break of $1/2$ is fixed by the cooling mechanism\\footnote{Slightly larger breaks may be possible by assuming certain source conditions, systematically varying magnetic fields, for example \\citep{Wil75}.}. Describing synchrotron spectra with arbitrary breaks is therefore inconsistent with any physically motivated model; in violation of our reasoning above, one could then even invent breaks to flatter spectra.    We used the detailed shape of cutoff spectra as computed by \\citet{HM87}. The spectral shape is the result of observing a mixture of electron populations, and the appearance of the entire source can be understood in terms of the temporal evolution of the initially accelerated population.  As long as the spectral shape is determined only by radiation losses, not even homogeneous source conditions are required.  The resulting spectrum rarely will follow a power law over many decades in frequency, but will be curved and in any case exhibit a convex shape.  In contrast, a more complex electron distribution could be produced if different parts of a source accelerate electrons to different maximum energies, and contain different numbers of electrons.  Such a source could be a jet with an inhomogeneous distribution of particles and magnetic fields. This source must be described by more than one electron population (or equivalently, as the sum of the emission from more than one source).  This corresponds precisely to the situation encountered in 3C\\,273's jet: the observed flattening of the spectrum implies that a description using a single electron population is inadequate.  This implies unambiguously that there are at least two components to the observed spectrum: one is the ``low-frequency'' synchrotron spectrum from radio through optical wavelengths. In addition, there is a ``high-frequency'' component responsible for the UV excess, and hence the observed flattening. We have suggested that this second component is the same as that responsible for the jet's X-rays \\citep{Jes02}, which could be a second synchrotron component \\citep{Roe00,Mareta01} or beamed inverse Compton emission \\citep[\\cf][]{Sam01} . In any case, our spectra prove that 3C\\,273's jet is not a homogeneous synchrotron source.  It is notable that besides 3C\\,273, the spectral energy distributions (SEDs) of M87's jet shows the same upward curvature: \\citet{Mar02} have presented SEDs for that jet from radio through X-rays (their Fig. 3).  They fitted two different single-population models, a Jaffe-Perola model assuming a randomly tangled magnetic field structure and pitch-angle scattering, and a Kardashev-Pacholczyk model, also assuming a tangled field but no pitch angle scattering. Only the latter provides an acceptable fit to the observed SEDs, but it is unclear how pitch-angle scattering could be avoided in a tangled field geometry.  Moreover, the spectra for knots HST-1 and D clearly show a steeper downturn near $10^{15}$\\,Hz \\citep[already hinted at by the ground-based spectra obtained by][ their Figs. 7 and 8]{MRS96} than the Kardashev-Pacholczyk spectrum, which is therefore inadequate. Instead, again a second spectral component must be invoked to account for the X-ray flux lying above the cutoff observed in the near-infrared--optical region.  \\citet{Mar02} viewed this as evidence in favour of the spatially stratified model for M87's jet presented by \\citet{Pereta99}.  \\subsection{Slow decline and smooth changes in \\gmax: the need for an extended acceleration mechanism}\\label{s:disc.acc}  Our determination of the maximum energy of synchrotron-emitting electrons shows a lower limit for the Lorentz factors of $\\gmax>10^7$ at the onset of the optical jet, 12\\arcsec\\ from the quasar core (Fig.\\,\\ref{f:ana.res.run-gmax}).  The maximum Lorentz factor first drops outward along the jet, and then stays nearly constant at a Lorentz factor of a few times $10^6$ at radii $>17\\arcsec$. It drops again by factor of 3 within the hot spot. This drop most likely is caused by the stronger magnetic field there.  In any case, the Lorentz factor of the particles emitting optical synchrotron radiation cannot be much lower than $10^5$ anywhere in the jet. Electrons with such high energies can only be observed close to their acceleration sites: as noted in all previous studies, synchrotron cooling timescales for these particles are of the order of a few hundred or thousand years \\citep[see][ \\emph{e.g.}]{GN75,RM91,MNR96}.  The corresponding loss scales are a few hundred parsec for the case of electrons freely streaming at the speed of light, which is unrealistic since they will have to follow tangled field lines.  The development of the maximum particle energy along the jet is therefore consistent with the absence of any synchrotron cooling. The same conclusion can already be drawn from the smooth changes of just the optical spectral index, which should most strongly reflect the synchrotron losses the particles undergo, and from the mere presence of optical synchrotron radiation along the entire jet \\citep{Jes01}. The apparent absence of cooling in the presence of synchrotron radiation can only be explained by a distributed acceleration mechanism.  The smooth changes of the spectral features along the jet indicate a correspondingly smooth variation of the physical conditions along the jet.  The present detailed high-resolution study of the jet in 3C\\,273 set out with the aim to map those regions in which particles are preferentially accelerated. Since the observed cooling is much less drastic than expected, the conclusion is that the \\emph{entire} jet is the particle acceleration region.  The fact that the jet has a much higher emissivity than the radio cocoon around it (see \\S\\ref{s:res.im.width}) also implies a difference in physical conditions between the jet and the surrounding material.  Furthermore, the jet's X-ray emission is a further sink of energy (whatever the emission process) which needs to be filled by the re-acceleration mechanism. Note that this does not preclude the possibility that the knots in the jet are due to shocks, at which particle acceleration can also take place.  But even in this case, particles must be accelerated between the shocks as well.  This corresponds to the ``jet-like'' acceleration mechanism proposed by \\citet{hs_II}.  The detailed physics of this distributed acceleration mechanisms are unknown, but several theoretical suggestions have been made.  Magnetic reconnection is one way to tap the energy stored in the jet fluid and convert it continuously into relativistic particles \\citep{Lit99}.  Another possible energy source is the velocity shear between the surface of the jet and the surrounding medium \\citep{StaOst02}.  This latter mechanism would be limited to particle acceleration within a thin surface layer. Since it appears that the jet emission is not edge-brightened, as would be expected for emission from a surface layer (\\S\\ref{s:res.im.vol}), surface velocity shear is unlikely to be the acceleration mechanism providing optical electrons along the entire jet.  It may, however, be appropriate to explain the second higher-energy emission component identified above (\\S\\ref{s:disc.pop}), provided this turns out to be edge-brightened.  A similar mechanism has recently been proposed by \\citet{RieMan02}: these authors suggest that velocity shear and centrifugal forces in a rotating jet flow may provide efficient particle acceleration.  The global energy budget of the jet is a related question. If the entire jet operates at the minimum energy condition, and if there is no strong, variable beaming, the increase in the minimum-energy field along the jet (linked to the increase in synchrotron luminosity) corresponds to an increase of the energy density stored in magnetic fields and relativistic particles by about one order of magnitude. This raises interesting questions about the energy budget of the jet: If the total energy flux in magnetic fields and relativistic particles is to be conserved, the jet would need to slow down \\citep[\\cf][]{GK03}. The necessity to account for the kinetic energy and momentum flux imposes further constraints.  On the other hand, it may be possible to convert jet kinetic energy into magnetic and particle energy, as in some of the acceleration mechanisms just mentioned, leading to a different mutual dependence of jet speed and magnetic energy density than in the pure energy conservation case.  A detailed examination of the jet's energy budget addressing these issues is certainly a worthy exercise, but beyond the scope of the present publication.  \\subsection{Correlations between brightness and maximum particle energy}\\label{s:disc.corr}  As described in \\S\\ref{s:ana.gmax.gmax}, there is a local positive correlation between surface brightness and maximum particle energy, although these quantities are \\emph{anti-}correlated globally.  The opposite correlations must be caused by two different physical mechanisms.  Although the cutoff frequency is the direct observable, the underlying physical parameters are the maximum particle energy and the jet's magnetic field.  In 3C\\,273, the global run of the maximum particle energy is practically identical to that of the cutoff frequency, and the local variations in cutoff frequency are much smaller than those in surface brightness (see Figs.\\,\\ref{f:ana.res.run-nuc} and \\ref{f:ana.res.run-gmax}).  We have established above (\\S\\ref{s:disc.acc}) that particles must be accelerated continuously within the jet.  In this case, the value of the maximum particle energy is obtained by equating the acceleration timescale $\\tau_{\\mathrm{acc}}$, which is the time during which a particle's energy is increased by a factor of two, with the synchrotron loss timescale $\\tau_\\mathrm{syn}$ \\citep[\\eg,][]{Longair_eloss}: \\begin{eqnarray} \\tau_{\\mathrm{acc}} &=& \\tau_\\mathrm{syn}(\\gamma_\\mathrm{max}) \\propto \\frac{1}{B^2 \\gamma_\\mathrm{max}}\\nonumber\\\\ \\Leftrightarrow \\gamma_\\mathrm{max} & \\propto &  \\frac{1}{B^2 \\tau_{\\mathrm{acc}}}. \\label{eq:gmax_tacc} \\end{eqnarray} The acceleration timescale is energy-independent both for diffusive shock acceleration \\citep[\\eg,][]{Bell78a} and for acceleration by reconnection \\citep[\\eg,][]{Lit99}. Thus, the behaviour of the maximum particle energy as a function of magnetic field depends critically on the nature of the acceleration mechanism, in particular the scaling of the acceleration time scale with the magnetic field strength.  We have taken the jet Doppler factor to be $\\mathcal{D}=1$ in all calculations up to here, which is justified because the value of the maximum particle energy as inferred from the jet's surface brightness and minimum-energy magnetic field only depends very weakly on the true value of the Doppler factor \\citep[cf. \\S\\ref{s:ana.equi} and][]{NeumDiss,MNR96}.  The observed surface brightness $S$ and cutoff frequency \\nuc\\ have the following scalings with $B$ and $\\mathcal{D}$ for a continuous jet \\citep[see][ e.g.]{BBR84}: \\begin{eqnarray} S & \\propto & B^{1-\\alpha}\\, \\mathcal{D}^{2-\\alpha} \\label{eq:S_BD}\\\\ \\nuc & \\propto & B\\, \\gamma_\\mathrm{max}^{2}\\, \\mathcal{D}. \\label{eq:nuc_BD} \\end{eqnarray} In the following discussion, we will \\emph{not} use the minimum-energy magnetic field, because its computation has introduced correlations between the magnetic field strength, the electron number density, and the surface brightness by construction.  \\subsubsection{Local correlation between brightness and maximum particle energy}  The only other jet which has been studied similarly well is that in M87. There are key differences between the behaviour of the maximum particle energy between these two jets. \\citet{MRS96} showed that M87 shows a strong correlation between surface brightness and cutoff frequency, while the maximum particle energy remains practically unchanged.  This correlation can be understood purely in terms of variations of the magnetic field strength, while the constancy of the maximum particle energy is explained by invoking acceleration in shocks \\citep[\\eg][]{SBM96}, by an extended acceleration mechanism \\citep{MRS96}, or by assuming a sub-equipartition magnetic field and relativistic time dilation which both lower the loss times sufficiently even without particle acceleration \\citep{HB97}. \\citet{MRS96} conclude that re-acceleration is responsible, and that the mechanism in M87's jet obeys the scaling $\\tau_{\\mathrm{acc}} \\propto B^{-2}$ to achieve a constant maximum particle energy (cf. Eqn.\\,\\ref{eq:gmax_tacc}).  It is useful to consider whether this acceleration mechanism proposed for M87 could be identical to the one at work in 3C\\,273, \\ie whether it is universal.  Since the maximum particle energy changes significantly along 3C\\,273's jet, the acceleration mechanism invoked by \\citet{MRS96} for M87 must be different from the one acting in 3C\\,273's jet.  To produce a local correlation between brightness and maximum particle energy via magnetic field variations, the acceleration time scale in 3C\\,273 would need to follow a scaling $\\tau_{\\mathrm{acc}} \\propto B^{s}$ with $s>2$ so that both \\gmax\\ and $S$ correlate with $B$ (Eqns.\\,\\ref{eq:gmax_tacc} and \\ref{eq:S_BD}).  In that case, local magnetic field variations could explain the observed correlations. However, since the brightness changes are larger than the changes in the maximum particle energy and cutoff frequency, it seems more likely that small-scale variations in the Doppler factor $\\mathcal{D}$ are responsible because $S$ varies more strongly with $\\mathcal{D}$ than with $B$.  But variations in $\\mathcal{D}$ by themselves are insufficient for 3C\\,273 because they do not produce a correlation between brightness and maximum particle energy.  To clarify the origin of this correlation, it will be necessary to disentangle the UV excess from the remainder of the jet emission to remove any uncertainty about the true value of the cutoff frequency, and ideally to obtain a measurement of the jet's magnetic field other than the minimum-energy estimate.  \\subsubsection{Global anti-correlation of brightness and maximum particle energy}  The global anti-correlation of brightness and maximum energy appears intuitive for synchrotron radiation: stronger synchrotron emission implies larger energy losses, and hence a lower maximum energy. However, as noted above, the actual behaviour of the maximum particle energy in the presence of particle acceleration depends on the details of the acceleration mechanism.  If the local correlation of brightness and \\gmax\\ can be explained by the nature of the acceleration mechanism, the global anticorrelation must be due to an entirely different effect which reverses the local correlation.  It may be possible that the local correlations are produced by a ``shock-like'' acceleration mechanism acting only in the brighter regions (\\eg if the bright regions \\emph{are} shocks), while electrons elsewhere are accelerated by the ``jet-like'' mechanism \\citep{hs_II}.  However, such a scenario is somewhat unappealing because there would need to be a certain fine-tuning to explain the overall smoothness of the spectral index changes.  Furthermore, although this may explain the different global and local behaviour within 3C\\,273 as well as the presence of two distinct emission components (\\S\\ref{s:disc.pop}), the difference to M87 remains, in which the ``jet-like'' mechanism is assumed to produce the positive correlation.  We therefore need to appeal to further differences.  \\citet{Mei96} suggested that the observations of M87 and 3C\\,273 could be unified if there was a link between the maximum particle energy and the velocity gradient along the jet, and the jet velocity was relativistic, so that changes in the velocity would lead both to changes in the maximum particle energy and in the Doppler factor and hence the apparent brightness.  We will discuss this idea elsewhere.    One of the main unsolved questions in the study of extragalactic jets is posed by the observations of optical synchrotron radiation over scales much larger than typical synchrotron loss scales and far from the ``working surfaces'' of the standard model \\citep{BlaRee74}. Another puzzle has been posed by the recent \\emph{Chandra} observations of X-rays from extragalactic jets, whose emission mechanism remains debated and which form a further sink of energy that has to be filled within the jet. The study of the physical conditions giving rise to the observed emission by keeping up the particle energy against strong synchrotron losses at optical wavelength and supplying the energy observed in X-rays can only be performed via a study of the synchrotron continuum, at the highest possible spatial resolution and covering the largest possible wavelength range.  \\subsection{Observations}  We present new HST and VLA images of the jet at $\\lambda\\lambda$ 3.6\\cm, 2.0\\cm, 1.3\\cm, 1.6\\micron, 620\\nm\\ and 300\\nm\\ matched to a common resolution of 0\\farcs3 (Fig.\\,\\ref{f:res.ima.images}).  We combine the data to obtain spectral energy distributions which we fit with theoretical synchrotron spectra (Fig.\\,\\ref{f:ana.fits.obs1}). Contrary to expectations for synchrotron emission, the observed spectra show a significant flattening in the infrared-ultraviolet wavelength range, implying that the emission cannot be modelled as synchrotron emission due to a single electron population, as has been assumed in previous studies of this jet's emission \\citep{MNR96}. Instead, additional emission must be present which can lead to the observed flattening. The most likely explanation is an additional flat-spectrum component in the ultraviolet. The same component may simultaneously be responsible for the jet's X-ray emission \\citep[for details, see][]{Jes02}.  The optical-ultraviolet spectral index map generated at 0\\farcs3 resolution (Fig.\\,\\ref{f:res.alpha.alpha}a) shows no strong correlation between local spectral index and surface brightness variations. There is, however, such a correlation on the radio-infrared and infrared-optical spectral index maps. Even these correlations are less pronounced than those found in the jet of M87 \\citep{Per01,PMB01}, where they are taken as evidence of localised particle acceleration in the knots of this jet (but see \\citealp{Mei96} for an alternative explanation).  As confirmed by our spectral fits (see below), the spectral index variations reflect variations of the maximum particle energy along the jet.  We find an offset between the radio and optical images of the hotspot that is consistent with the prediction by \\citet{HM87}. This adds further to the confidence that the assumed hot spot model is correct \\citep{hs_II}.  However, the role of the tip of the jet H1 and its relation to the hot spot remains to be clarified, perhaps in a model fully describing the hot spot flow in three dimensions, instead of just one.  If the jet flow is still highly relativistic just upstream of the hot spot, the details of the flow's deceleration in the hot spot may also change its appearance \\citep{GK03}.   \\subsection{Synchrotron spectral fits}  We fitted the observed spectra with model synchrotron spectra to extract physical information, in particular the maximum particle energy and its variation along the jet.  Model synchrotron spectra according to \\citet{HM87} have been used to determine the cutoff frequency \\nuc.  It is mainly determined by the infrared-optical spectral index.  We infer the maximum particle energy from the fitted cutoff frequency by assuming an equipartition magnetic field $B_\\mathrm{min}$.  It is a general result that the value of $B_\\mathrm{min}$ is very robust with respect to assumptions about the details of the spectral shape, as long as the spectrum has a power law index that is steeper than $\\alpha=-0.5$.  The cutoff frequency decreases from above $5\\times10^6$ in region A at $r\\approx 13\\arcsec$ and settles to a plateau of order $10^6$ at $r\\ga 18\\arcsec$. A further decrease to $\\gmax \\approx 3\\times10^5$ is only observed at the position of the radio hot spot, beyond which there is no optical, but still infrared emission.  This global run corresponds to the findings of previous studies at 1\\farcs3 resolution, which showed that the maximum particle energy decreases as the radio jet brightness increases.  As expected from the spectral index variations, there is an additional weak correlation between local variations of maximum particle energy and surface brightness.  An earlier examination of the effect of relativistic beaming on the inferred synchrotron lifetime had shown that the overall slow decrease of \\gmax\\ cannot be explained by an enhancement of the true synchrotron lifetime by beaming effects \\citep{Jes01}. We concluded that particle acceleration must take place distributed across the entire jet. This conclusion is strengthened by the X-ray observations, whatever the X-ray emission mechanism: The radiative lifetimes of X-ray synchrotron electrons in an equipartition field of 20\\,nT are a few tens of years, those of UV-emitting electrons are a few hundreds of years.  In the inverse Compton scenario, the combined effects of relativistic beaming and lower magnetic fields lead to lifetimes of a few times $10^4$ years for UV-emitting electrons.  In both cases, particles must therefore be accelerated continuously within the entire jet \\citep[``jet-like'' acceleration,][]{hs_II}.  The same conclusion is drawn from the fact that the physical conditions within the jet of 3C\\,273 vary very smoothly down to the spatial scales of $0.9h_{70}^{-1}\\kpc$ resolved here.  This does not preclude the possibility that the enhanced-brightness regions are shocks at which particles are accelerated -- but even if they are, re-acceleration between them is necessary to explain the observed spectral features.  The details of the acceleration mechanism determine the scaling of the maximum particle energy with changes in the magnetic field.  The interplay between the ``jet-like'' and ``shock-like'' acceleration mechanisms might explain the conflicting global and local behaviour of the maximum particle energy in 3C\\,273.  However, both mechanisms must be different from the one invoked by \\citet{MRS96} for M87's jet (\\S\\ref{s:disc.corr}).  We stress again the observed flattening of the high-frequency spectrum, which implies that the jet's emission cannot be assumed to arise from a single electron population, but requires the presence of an additional emission component. In particular, the spectrum of knot A cannot be considered as a single power law from radio through X-rays, as concluded from lower-resolution data \\citep{Roe00,Mareta01}.  Instead of considering the jet as a homogeneous synchrotron source filled with a simple electron population, we must realise that jets have an internal energetic structure \\citep[\\emph{cf.}][]{Pereta99,Per01}.  Future work to investigate the acceleration mechanism at work in this and other jets must be based on these results.  \\subsection{Future work}  With the present radio-optical data, it has only been possible to detect the high-frequency hardening of the jet emission, and thus the presence of a second emission component. No statements can be made about the spectral shape of this ``high-energy'' component beyond the approximate match of the X-ray extrapolation and the UV excess.  We have been awarded HST time to extend the wavelength coverage into the far-ultraviolet (150\\,nm). Using these data, we will be able to constrain the run of the optical synchrotron cutoff more accurately, and to characterise the spectrum of the UV excess. This will achieve a separation of the two spectral components from each other.  Particular insight can be expected by considering the high-frequency spectra in those regions of the jet showing the strongest difference between radio, optical and X-ray morphology (\\S\\ref{s:res.im.morph}).  In order to test theories for extended particle acceleration, we require detailed predictions both from theoretical work and from magneto-hydrodynamical simulations of the jet flow, in particular regarding the spatial distribution of relativistic particles.  The most sensitive observational tool is multi-wavelength polarimetry resolving the jet width.  This allows not only a study the particles' energy distribution, but also an assessment whether particles at different energies probe the same underlying magnetic field, that is, whether they actually occupy the same emitting volume. Thus, multi-wavelength polarimetry is the natural followup programme for our detection of multiple emission components. We will use our guaranteed time for the adaptive optics camera CONICA/NAOS at the VLT to obtain near-infrared polarimetry at a resolution comparable to the data presented here. However, optical polarimetry of 3C\\,273's jet making full use of the HST's capabilities is still missing.  Exciting new facts about the jet in 3C\\,273 have been revealed by each advance in optical and radio observing technology. From the surprising observational facts encountered here, more surprises are yet to be expected \\emph{en route} to an understanding of the physics of this object. Current advances in numerical work are also expected to shed further light on the structure and dynamics of jets. As an extreme object of its kind, 3C\\,273 will remain high on the agenda of astrophysicists studying jet phenomena."
    },
    "0410/astro-ph0410699_arXiv.txt": {
        "abstract": "Planetary Nebulae (PNe) are the result of heavy mass loss of the asymptotic giant branch (AGB) stars. They are understood in terms of Kwok's (1978) interacting-winds model as the product of the mass-loss history on the AGB and the central star (CS) evolution. Since the CS are close to the end of nuclear burning and at their hottest stage of evolution then, precise modeling of these pre-white dwarfs is a prerequisite in order to calculate reliable ionizing fluxes which are crucial input for the presently available 3D photoionization codes. In the framework of a systematic study of PNe which show evidence for an interaction with the ISM, we present a NLTE analysis of their CS. ",
        "introduction": "Recent developments in both, observational and deprojection techniques, spectral analysis, and numerical methods facilitate to closely examine and model PNe and their CS consistently. Spectral analysis of CSPN by means of modern NLTE model atmosphere techniques provides information about photospheric parameters. In comparison with evolutionary calculations, evolutionary status, distances, masses, and luminosities can be determined. The spectroscopic distances enable us to deal with precise linear diameters of the PNe. The stellar model fluxes are an important input as realistic ionizing spectra in nowadays available 3D photoionization codes (e.g\\@. Ercolano et al\\@. 2003) for analyses of the PNe. The consistent modeling of PN and CS allows to determine the yields of nuclear processed material which goes back to the ISM via PNe and thus, to investigate the chemical evolution of our Galaxy.  In the last decade, PNe which show evidence for an interaction with the ISM have been found. This process is an indicator for the evolution of their CS which are at their hottest stage of evolution close to the end of nuclear burning where gravitational effects become dominant (i.e\\@. they display directly the formation of white dwarfs). The highly evolved CS are no longer dominating the processes in the PNe (Kerber \\& Rauch 2001); the nebulae display brightness asymmetries that reflect the degree of the interaction process. These complex objects are crucial tests for our models as well as for evolutionary theory.  Once regarded a curiosity, this complex PN -- ISM interaction process appears to be a common phenomenon which allows to determine also properties of the ambient ISM. Recently, we have presented determinations of proper motions and Galactic orbits (Kerber et al\\@. 2004a, 2004b). These unveil thin and thick disk populations and cast light on interaction with the ISM. First hydrodynamical test calculations (M\\\"uller et al\\@. 2004) for these objects were performed with the aim of a more quantitative understanding of processes in PNe in decay. One of the pre-requisites for reliable modeling are spectral analyses of the CS. We present here preliminary results from a study which is based on line-blanketed models including H and He. ",
        "conclusions": ""
    },
    "0410/astro-ph0410650_arXiv.txt": {
        "abstract": "We study the extragalactic protons with universal spectrum, which is independent of mode of propagation, when distance between sources is less than the propagation lengths, such as energy attenuation length or diffusion length (for propagation in magnetic fields). The propagation features in this spectrum, the GZK cutoff, dip and bump, are studied with help of modification factor, which weakly depends on the generation spectrum index $\\gamma_g$. We argue that from the above features the dip is the most model-independent one. For the power-law generation spectrum with $\\gamma_g=2.7$ the dip is very well confirmed by the data of all existing detectors, which gives the strong evidence for extragalactic protons propagating through CMB. We develop the AGN model for origin of UHECR, which successfully explains the observed spectra up to $1\\times 10^{20}$~eV and transition from galactic to extragalactic cosmic rays. The calculated spectrum has the GZK cutoff, and the AGASA excess of events at $E \\gsim 1\\times 10^{20}$~eV needs another component, e.g.\\ from superheavy dark matter. In case of weak extragalactic magnetic fields this model is consistent with small-angle clustering and observed correlation with BL Lacs. ",
        "introduction": "The systematic study of Ultra High Energy Cosmic Rays (UHECR) started in late fifties after construction of Volcano Ranch (USA) and Moscow University (USSR) arrays. During next 50 years of research the origin of UHE particles, which hit the detectors, was not well understood. At present due to the data of the last generation arrays, Haverah Park (UK), Yakutsk (Russia), Akeno and AGASA (Japan), and Fly's Eye and HiRes (USA) \\cite{expCR} we are probably very close to understanding the origin of UHECR, and the data of Auger detector \\cite{auger} will  undoubtedly substantially clarify this problem.  On the theoretical side we have an important clue to understanding the UHECR origin: the interaction of extragalactic protons, nuclei and photons with CMB, which leaves the imprint on UHE proton spectrum, most notably in the form of the Greisen-Zatsepin-Kuzmin (GZK) \\cite{GZK} cutoff.  We shall shortly summarize the basic experimental results and the results of the data analysis, important for understanding of UHECR origin (for a review see \\cite{NaWa}). \\\\*[2mm] {\\em (i)}~~ The spectra of UHECR are measured \\cite{expCR} with good accuracy at 1 - 100 EeV, and these data have a power to reject or confirm some models. The discrepancy between the AGASA and HiRes data at $E > 100$~EeV might have the statistical explanation \\cite{MBO}.\\\\ {\\em (ii)}~~ The mass composition at $E\\gsim 1$~EeV (as well as below) is badly known (for a review see \\cite{Wa}). The different methods give the different mass composition, and the same methods disagree in different experiments. Probably the most reliable method of measuring the mass composition is given by elongation rate (energy dependence of maximum depth of shower $X_{\\rm max}$) measured by the fluorescent method. The data of Fly's Eye in 1994 \\cite{expCR} favored iron nuclei at 1~EeV with a gradual transition to the protons at 10~EeV. The further development of this method by the HiRes detector, which is the extension of Fly's Eye, shows the transition to the proton composition already at 1~EeV \\cite{Sokol}.\\\\*[2mm] {\\em (iii)}~~ The arrival directions of particles with energy $E \\geq 4\\times 10^{19}$~eV show the small-angle clustering within the angular resolution of detectors. AGASA found 3 doublets and one triplet among 47 detected particles \\cite{clust-AGASA} (see the discussion \\cite{FW}). In the combined data of several arrays \\cite{clust-tot} there were found 8 doublets and 2 triplets in 92 events. The stereo  HiRes data \\cite{clust-Hi} do not show small-angle clustering for 27 events at $E \\geq 4\\times 10^{19}$~eV, maybe due to limited statistics.  Small-angle clustering is most naturally explained in case of rectilinear propagation as a random arrival of two (three) particles from a single source \\cite {DTT}. This effect has been calculated in Refs. \\cite{FK,YNS1,YNS2,BlMa,KaSe}. In the last four works the calculations have been performed by MC method and results agree well. According to \\cite{KaSe} the density of the sources, needed to explain the observed number of doublets is $n_s= (1 - 3)\\times 10^{-5}$~Mpc$^{-3}$. In \\cite{BlMa} the best fit is given by $n_s \\sim 1\\times 10^{-5}$~Mpc$^{-3}$ and the large uncertainties (in particular due to ones in observational data) are emphasized.\\\\*[2mm] % {\\em (iv)}~~ There have been recently found the statistically significant correlations between direction of particles with energies $(4 - 8)\\times 10^{19}$~eV and directions to AGN of the special type - ~BL Lacs \\cite{corr} (see also the criticism \\cite{Sar} and the reply \\cite{TT}).  The items {\\em (iii)} and {\\em (iv)} favor rectilinear propagation of primaries from the point-like extragalactic sources, presumably AGN. This statement provokes many questions:\\\\ Is rectilinear propagation the only solution to clustering (e.g.\\ does it appear in propagation of protons in strong magnetic fields)? In case the correlation with AGN is true, what are the primaries (protons in weak magnetic field or neutral particles)?  In this paper we shall analyse the origin of UHECR in two steps. In the first one we will use only most reliable observational data, namely the energy spectra and fluxes under the most conservative assumption that primaries are extragalactic protons. We shall calculate spectra for propagation in weak and reasonably strong magnetic fields, and demonstrate the strong evidence in favor of protons as the primaries at energies 1 - 80~EeV. This part of analysis is almost model independent. In the second step we shall include the data {\\em (iii)} and {\\em (iv)} and formulate our model. In the framework of our model we shall discuss the connection between galactic and extragalactic components of cosmic rays, and the problem with superGZK particles, i.e.\\ ones at $E> 100$~EeV. ",
        "conclusions": "We developed the most conservative scenario for the observed UHECR as extragalactic protons.  There are three signatures of UHE protons propagating through CMB: GZK cutoff, bump and dip. While bump is argued to be absent in the diffuse spectra and presence of the GZK cutoff is questioned by the AGASA data, the dip is confirmed with very good accuracy by the AGASA and HiRes data (see Figs.~\\ref{mfactorA}, \\ref{mfactorH}).  The predictions for the dip in terms of modification factor are very weakly model-dependent: its shape varies but little with $\\gamma_g$ in the interval $2.0 - 2.7$, the assumption of homogeneity in distribution of the sources is well justified, because UHE protons are collected from the large distances of order $l_{\\rm att} \\sim 1000$~Mpc, and its shape is valid for both weak and reasonably strong magnetic fields. For conversion of the observed spectra into modification factor two free parameters are needed: $\\gamma_g$ and normalization. For 19 energy bins of the Akeno-AGASA data and two free parameters, the agreement is characterised by $\\chi^2=19.06$ and $\\chi^2/$d.o.f.=1.12 for d.o.f.=17.  Modification factor must satisfy $\\eta \\leq 1$. At $E \\leq 4\\times 10^{17}$, as Fig.~\\ref{mfactorA} shows, $\\eta > 1$, which signals about appearance of galactic CR component at energy higher, but not much than given above. This conclusion agrees with recent data of HiRes (see Fig.~\\ref{x_max}) which show that proton component becomes dominant at $E \\gsim 1\\times 10^{18}$~eV.   In our model we assume AGN as the sources, with generation spectrum given by Eq.~(\\ref{q-gen}). With emissivity ${\\cal L}_0 \\approx 3.5\\times 10^{46}$~erg/Mpc$^3$yr this spectrum describes well the AGASA, HiRes, Fly's Eye and Yakutsk data. The galactic flux of iron nuclei, calculated in energy range $1\\times 10^{17} - 1\\times 10^{18}$~eV from extragalactic flux of our model and all-particle Akeno spectrum agrees well with iron-nuclei flux measured by KASCADE. The spectrum shape is consistent with the Hall diffusion.  Our spectrum has the GZK cutoff and it is consistent with HiRes, Fly's Eye and Yakutsk data. For the explanation of the AGASA excess at $E \\gsim 1\\times 10^{20}$~eV another component of UHECR is needed. In this paper we discuss UHECR from SHDM as such component. The calculated spectrum agrees well with observations. We argue that with predicted ratio of primary photons to protons $\\gamma/p \\approx 2 - 3$, SHDM does not contradict the AGASA observations \\cite{agasa-gamma}."
    },
    "0410/astro-ph0410466_arXiv.txt": {
        "abstract": "We present {\\it Hubble Space Telescope} Advanced Camera for Surveys multicolor coronagraphic images of the recently discovered edge-on debris disk around the nearby ($\\sim10$ pc) M dwarf AU Microscopii.  The disk is seen between $r = $0\\farcs 75 -- 15'' (7.5 -- 150 AU) from the star.  It has a thin midplane with a projected full-width-at-half-maximum (FWHM) thickness of 2.5 -- 3.5 AU within $r < 50$ AU of the star that increases to 6.5 -- 9 AU at $r \\sim 75$ AU. The disk's radial brightness profile is generally flat for $r < 15$ AU, then decreases gradually ($I \\propto r^{-1.8}$) out to $r \\approx 43$ AU, beyond which it falls rapidly ($I \\propto r^{-4.7}$).  Within 50 AU the midplane is straight and aligned with the star, and beyond that it deviates by $\\sim3^{\\circ}$, resulting in a bowed appearance that was also seen in ground-based images.  Three-dimensional modelling of the disk shows that the inner region ($r < 50$ AU) is inclined to the line-of-sight by $<1^{\\circ}$ and the outer disk by $\\sim3^{\\circ}$.  The inclination of the outer disk and moderate forward scattering ($g \\approx 0.4$) can explain the apparent bow. The intrinsic, deprojected FWHM thickness is 1.5 -- 10 AU, increasing with radius.  The models indicate that the disk is clear of dust within $\\sim12$ AU of the star, in general agreement with the previous prediction of 17 AU based on the infrared spectral energy distribution.  The disk is blue, being 60\\% brighter at $B$ than $I$ relative to the star.  One possible explanation for this is that there is a surplus of very small grains compared to other imaged debris disks that have more neutral or red colors.  This may be due to the low radiation pressure exerted by the late-type star.  Observations at two epochs show that an extended source seen along the midplane is a background galaxy. ",
        "introduction": "Collisions of planetesimals around a star produce dust grains, creating a debris disk that can be detected in thermal emission and scattered light. Without the continual replenishment of these grains, most of the disk will disappear due to radiation and wind pressure, Poynting-Robertson drag, grain coagulation, and tidal interactions with planets.  The lifetime of such disks is not fully understood, and it may depend strongly on the luminosity of the star.  While numerous detections and images of optically-thick accretion disks have been made, only a few debris disks have been resolved.  The current inventory of imaged debris disks is largely derived from {\\it IRAS}-measured infrared excesses of stars.  This list is biased toward stars of spectral types A -- F because their luminosities are high enough to heat a debris disk to a {\\it IRAS}-detectable level.  More luminous stars would blow away such disks while cooler stars cannot heat them enough to be detected.  Most of the resolved debris disks surround A stars ($\\beta$ Pictoris, Vega, Fomalhaut, HD 141569A, and HR 4796), and only a few have been seen in scattered light.  The sensitivity limits of {\\it IRAS} leave uncertain the frequency of debris disks around later-type stars.  The disk around $\\epsilon$ Eridani, a K2V star with one of the largest {\\it IRAS}-measured stellar excesses, was detected by {\\it IRAS} only because of its very close proximity (3.2 pc) to us.  The {\\it Spitzer Space Telescope} should provide a much more comprehensive catalog of disk candidates around later-type stars.  There are only two M dwarfs, AU Microscopii and Hen 3-600 (Song et al. 2002), not associated with molecular clouds that have {\\it IRAS}-measured infrared excesses.  AU Mic (HD 197481, GJ 803) is a M1Ve flare star and a BY Dra-type variable ($V_{max} = 8.59, \\Delta V \\approx 0.15$; Cutispoto, Messina, \\& Rodon\\`{o} 2003). Barrado y Navascu{\\' e}s et al.  (1999) identified AU Mic as part of the $\\beta$ Pictoris moving group, an association of nearby (10 -- 50 pc) young stars.  It is at a {\\it Hipparcos}-measured distance of 9.94 pc (Perryman et al. 1997) and is about 10 Myr old (Zuckerman et al. 2001).  Its excess implies that AU Mic has a substantial amount of circumstellar material. These characteristics led the {\\it Hubble Space Telescope} ({\\it HST}) Advanced Camera for Surveys (ACS) Investigation Definition Team (IDT) to include the star in its circumstellar disk imaging program.  A disk around AU Mic was recently imaged from the ground using coronagraphs (Kalas, Liu, \\& Matthews 2004, hereafter KLM04; Liu 2004, hereafter L04). These images show a disk that at first glance resembles that of $\\beta$ Pic. It is viewed nearly edge-on and has a radius of at least 21'' (210 AU).  The Keck adaptive optics image of L04 also shows small-scale variations in the midplane that might be due to localized density enhancements caused by unseen substellar companions that perturb the disk.  So far, imaging within 1\\farcs 5 of the star has not been possible from the ground.  The high resolution of ACS and its ability to see closer to the star can provide more detailed views of this disk. ",
        "conclusions": "The disk of AU Mic has been imaged in scattered light at high resolution with {\\it HST} in {\\it BVI}-like passbands to within $\\sim0\\farcs 5$ (5 AU) of the star.  The surface brightness profiles of the two sides of the disk are generally similar, especially when compared to those of the $\\beta$ Pic disk, but localized brightness asymmetries are seen, especially in the SE side. The disk at $r > 100$ AU is $\\sim2\\times$ brighter in the NW than the SE.  The disk is blue relative to the star, perhaps indicating a relative surplus of small particles compared to grain size distributions in neutral or red disks seen around earlier-type stars.  The radiation pressure from a late-type star like AU Mic may be insufficient to force these small grains out of the disk.  An extended feature at $r = 96$ AU on the NW midplane is a background galaxy.  Modelling indicates that the inner ($r < 50$ AU) region is seen very close to edge-on ($\\theta < 1^{\\circ}$) while the outer is inclined $\\sim3^{\\circ}$ toward us.  The disk appears to be clear of material within 12 AU of the star, consistent with the infrared SED.  The disk has a fairly sharp midplane with an intrinsic vertical FWHM of $<1$ AU at $r < 50$ AU which increases to $\\sim3$ AU at $r = 100$ AU.  The degree of forward scattering by the grains is comparable to that seen in other debris disks, namely $\\beta$ Pic."
    },
    "0410/astro-ph0410185_arXiv.txt": {
        "abstract": "We have studied the properties of an infrared cluster embedded in the gas and dust of the southern part of the Carina Nebula (NGC\\,3372), where the probable existence of current star formation has already been predicted. We used mid-infrared (A \\& C bands) and near-infrared (J\\,H\\,K$_{\\rm s}$) images from the MSX and the 2MASS surveys respectively, combined with an optical H$_\\alpha$ narrow-band filter image obtained at the CTIO. The infrared star cluster has at least 17 members, and its parameters, radius and stellar density are in very good agreement with high- to intermediate-mass star formation scenarios. The detected IR sources have roughly the same intrinsic infrared excess determined from their position in colour-colour and colour-magnitude diagrams, suggesting that these objects could be related to pre-main sequence stars of high to intermediate mass. Furthermore, we present a low-dispersion spectrum of LS\\,1883 (O9.5\\,V) star located near the centre of the IR cluster. The position of this object in the colour-colour and colour-magnitude IR diagrams lies close to the reddening vector of an ZAMS O9\\,V spectral type star, and it seems to be the first star of this cluster to emerge. All these facts are consistent with the current star-forming scenarios associated with highly embedded star clusters. ",
        "introduction": "The Carina Nebula (NGC\\,3372) is one of the most interesting massive star-forming regions in the Galaxy which can be observed from the Southern hemisphere, containing a large population of hot stars \\citep{1971ApJ...167L..31W,1973ApJ...179..517W}. This extraordinary region is available to observations using medium-sized telescopes and it hosts several well-known young open clusters, such as Trumpler 14, 15 and 16, Collinder 228 and 232, and Bochum 10 and 11, \\citep{1993AJ....105..980M,1995RMxAC...2...51W}.  NGC\\,3372 is an evolved H{\\sc ii} region, with dust and neutral gas evacuated from its core by the action of strong winds from its hot massive stars. However, embedded star-forming cores suggest that star formation has not ceased \\citep{2000ApJ...532L.145S}. These cores allow us to study the very initial phases of star formation which, in general, are deeply embedded in dust and/or are located in heavily reddened lines of sight. The near-infrared (near-IR) and mid-IR domains are therefore necessary to study them \\citep{1991fesc.book....3L}. Recently, \\cite{2002MNRAS.331...85R} discovered several groups of diffuse IR sources and studied the interaction between the young stars and the molecular material that surrounds them. \\cite{2004A+A...418..563R} found 12 candidate embedded clusters using mid-IR Midcourse Space Experiment (MSX) data; some of these sources are also proposed as candidates by \\cite{2000ApJ...532L.145S}.  Among all these candidates, G287.84--0.82 is a bright nebulosity near LS\\,1883 (CPD-59$^{\\circ}$\\,2661; $\\alpha_{\\rm 2000}$=10h\\,45m\\,53.6s, $\\delta_{\\rm 2000}$=-59$^{\\circ}$\\,57\\arcmin\\,02.2\\arcsec), which seems to be a prime candidate of current star formation. \\cite{2001A&A...376..434D} and \\cite{2003A&A...400..533D} searched for candidates to star clusters and groups between $|b|$\\,$<$\\,10$^{\\circ}$ and 230$^{\\circ}$\\,$<$\\,{\\it l}\\,$<$\\,350$^{\\circ}$ using the Two-Micron All-Sky Survey (2MASS)\\footnote{http://www.ipac.caltec.edu/2mass/}. They have suggested that this region could be a partially resolved cluster in a reflection nebula in agreement with the previous classification vdBH-RN43 \\citep{1975AJ.....80..208V}. They have also given an estimate for the angular diameter of D=1.5\\,arcmin at a distance of 2.5\\,kpc (equivalent to a linear diameter of 1.1 pc), and a possible multiplicity of about 10 members. \\cite{2003A&A...399..141M}, who merged several catalogues of reflection nebula in order to create a uniform catalogue containing 913 objects, also classified it as an object of this type.  The aim of this paper is to determine the properties of this deeply embedded IR star cluster, by using mid-IR 8\\,-\\,21 $\\mu$m images obtained with the MSX and the 2MASS IR atlas \\citep{1997ilsn.proc...25S} in the J, H and K$_{\\rm s}$ image bands. Also, we have included new spectroscopic data of LS\\,1883 and SS73\\,24, two bright stars which appear projected towards the nebulosity, in order to determine the actual association of these sources with the embedded cluster. ",
        "conclusions": "We have derived some of the optical and near-IR properties of G287.84--0.82 confirming that it is an IR star cluster, embedded in the head of a dusty pillar in the southern part of the Carina nebula (NGC\\,3372). Also, we have established some important properties of the brightest optical sources in the field of the cluster. These can be summarized as follows. \\begin{enumerate} \\item The surface density of IR excess sources located in the IR cluster is one order of magnitude greater than the background density. \\item The density of IR excess sources outside the 30\\,arcsec radius is roughly five times lower than the background density. \\item The cluster has at least 17 members, and its parameters, radius and stellar density, are in good agreement with high to intermediate mass star formation scenarios. \\item C-C and C-M diagrams show roughly the same IR excess as most of the IR cluster member candidates. \\item These IR excesses suggest that these objects could be intermediate-mass pre-main sequence stars. \\item LS\\,1883 is located near to the IR cluster core. We confirm the O9.5\\,V spectral type for this star based on the presence of He\\,{\\sc ii} absorption lines. This spectral classification agrees with the position in the C-C and C-M diagrams for a moderately reddened O9\\,V star. LS\\,1883 seems to be the first ZAMS star emerging from the cluster, and it could be producing a blister-like H{\\sc ii} region. Optical narrow band images reveal an ionizing gradient in the nebula surrounding the star. \\item The positions of SS73\\,24 in the C-C and C-M diagrams suggest that this star has a strong intrinsic near-IR continuum emission, although we cannot establish any relationship between SS73\\,24 and the IR cluster. \\item The spectral characteristics of SS73\\,24 suggest that this star could be related to a strong mass-loss stage in an interacting binary. \\item The mid-IR MSX images show emission from the IR cluster and from warm dust associated with SS73\\,24. \\end{enumerate}  Some of these facts are consistent with the current star formation scenarios associated with highly embedded IR star clusters, where the most massive component is an O-type star. In order to be able to reveal the true nature of each individual IR cluster member, spectroscopic observations and images with higher spatial resolution in the near-IR and mid-IR wavelength range are needed.  \\vskip 0.5cm  {\\it Acknowledgments} We are very grateful to Michael Taylor and Silvia Giovagnoli for their aid in the clarification of some passages of the text. We specially want to thank Elena Terlevich, who read this paper and suggested to us a better way to present different aspects of this work. Substantial improvements to the paper resulted from a careful revision and suggestion from an anonymous referee. This research made use of data products from the Midcourse Space Experiment. Processing of the data was funded by the Ballistic Missile Defense Organization with additional support from NASA Office of Space Science. This research has also made use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the IR Processing and Analysis Centre, funded by the NASA and the NSF. We want to thank the director and staff of CTIO for the use of their facilities and kind hospitality during the observing run. We acknowledge the use at CASLEO of the CCD and data acquisition system supported under US NSF grant AST-90-15827 to R.M. Rich. Also, we thank Julia Arias for providing us with the low resolution spectra of LS\\,1883 and SS73\\,24."
    }
}